UBC Theses and Dissertations

UBC Theses Logo

UBC Theses and Dissertations

Stellar populations in spiral galaxies : broadband versus spectroscopic viewpoints MacArthur, Lauren Anne 2005

Your browser doesn't seem to have a PDF viewer, please download the PDF to view this item.

Item Metadata

Download

Media
831-ubc_2005-105313.pdf [ 49.72MB ]
Metadata
JSON: 831-1.0085579.json
JSON-LD: 831-1.0085579-ld.json
RDF/XML (Pretty): 831-1.0085579-rdf.xml
RDF/JSON: 831-1.0085579-rdf.json
Turtle: 831-1.0085579-turtle.txt
N-Triples: 831-1.0085579-rdf-ntriples.txt
Original Record: 831-1.0085579-source.json
Full Text
831-1.0085579-fulltext.txt
Citation
831-1.0085579.ris

Full Text

Stellar Populations in Spiral Galaxies Broadband versus Spectroscopic Viewpoints by Lauren Anne MacArthur B.Sc, University of Guelph, 1999 M.Sc, University of British Columbia, 2001 A THESIS SUBMITTED IN PARTIAL FULFILMENT OF THE REQUIREMENTS FOR THE DEGREE OF DOCTOR OF PHILOSOPHY in The Faculty of Graduate Studies (Astronomy) THE UNIVERSITY OF BRITISH COLUMBIA July 25, 2005 © Lauren Anne MacArthur, 2005 11 Abstract This thesis addresses the stellar population content in the bulges and disks of spiral galaxies using broad-band and spectroscopic data. The results can be used to con-strain models of galaxy formation in addition to establishing a comprehensive, model-independent, picture of colour and line-index gradients in spiral galaxies. Building upon my Masters study of structural parameters in spiral galaxies, I use the largest collection of multi-band (optical and IR) surface brightness profiles for face-on and moderately-tilted galaxies to extract radial colour profiles. The colour gradients are then translated into age and metallicity gradients by comparison with stellar population synthesis (SPS) models considering a range of star formation histories, including recent bursts. Based on their integrated light, we find that high surface brightness (SB) regions of galaxies formed their stars earlier than lower SB ones, or at a similar epoch but on shorter timescale. At a given SB level, the star formation histories are modulated by the over-all potential of the galaxy such that brighter/higher rotational velocity galaxies formed earlier. This formation "down-sizing" implied by our results is inconsistent with current implementations of semi-analytic structure formation models. In order to alleviate concerns that our colour gradients could be affected by dust reddening, we designed a similar spectroscopic investigation and explored the dust sen-sitivity of absorption-line indices. The latter test makes use of the latest SPS, models incorporating a multi-component model for the line and continuum attenuation due to dust. For quiescent stellar populations (e.g. spheroids and globular clusters), dust extinction effects are small for most indices with the exception of the 4000 A break. For models with current star formation, many indices may suffer from dust reddening and any departures depend on age, dust distribution, and the effective optical depth. However, a number of useful indices are only weakly affected by dust extinction (es-Abstract iii pecially relative to typical measurement uncertainties), and can thus be safely used in spectroscopic studies of dusty systems. Motivated by our previous results, we embarked on a long-term project to deter-mine age and metallicity gradients from absorption features in spiral galaxy spectra from their centers and extending well into their disks for the first time. A pilot sample of 8 barred and unbarred nearby spiral galaxies was observed with Gemini/GMOS and line indices with S/N > 40 per A were extracted out to ~ 1 -1.5 disk scale lengths. Emission contamination and a suite of instrumental effects were fully taken into account. Reli-able line-indices compared with the latest SPS models reveal that; i) late-type bulges and inner disks are generally young (light-weighted SSP ages < 1 to 6 Gyr) with no age gradients, and ii) late-type spirals have metallicities close to solar at their cen-ter decreasing rapidly outward (with gradients of ~ —0.3 to —0.7 dex per rd). Disk contamination into the bulge is an issue but the inferred young ages exclude the in-terpretation of early rapid collapse or merger origin of late-type bulges. While secular evolution processes are likely the predominant mechanism for the bulge build-up, the strong observed metallicity gradients are not currently supported by such models. Our analysis has demonstrated the feasibility of age and metallicity determinations from longslit spectroscopy of gas rich, star-forming, systems. However, a systematic com-parison with galaxy properties requires a larger statistical sample. New GMOS longslit spectra acquired recently will augment our data base and contribute to the build-up of this instrumental data base for the study of bulge and disk formation models. iv Contents Abstract ii Contents iv List of Tables vii List of Figures viii Acknowledgements xiii Acronyms xiv 0.1 List of Acronyms xiv 1 Introduction 1 1.1 Galaxy Formation and Evolution 1 1.1.1 Bulge Formation Scenarios 5 1.2 Observational Constraints 9 1.2.1 Colour Gradients in Spiral Galaxies 10 1.2.2 Spectroscopic Age and Metallicity Indicators 12 1.2.3 Structural Components of Disk Galaxies 16 1.3 Thesis Outline 17 2 Colour Gradients in Spiral Galaxies 19 2.1 Introduction 19 2.2 The Data 21 2.3 Galaxy parameters 26 2.4 Stellar Population Models 28 Contents v 2.5 Dust effects 39 2.6 Colour-Colour Profiles 43 2.7 Model Fitting 47 2.8 Results 50 2.8.1 Local and Global Trends in Age and Metallicity 50 2.8.2 Global Trends in Age and Metallicity Gradients 53 2.8.3 Age Gradients 56 2.8.4 Metallicity Gradients 58 2.9 Discussion 58 3 Dust Sensitivity of Absorption-Line Indices 66 3.1 Introduction 66 3.2 Description of Indices Studied 68 3.2.1 The Lick/IDS System 68 3.2.2 Higher Order Balmer Indices 71 3.2.3 The 4000 A Break 72 3.2.4 The Near-IR Ca UTriplet Indices 73 3.2.5 The Rose Indices 74 3.3 Models 75 3.4 Results 80 3.4.1 Dust Sensitivity of the Lick Indices 80 3.4.2 Dust sensitivity of the 4000 A Break 87 3.4.3 Dust Sensitivity of the Ca UTriplet Indices 89 3.4.4 Dust Sensitivity of the Rose Indices 90 3.5 Age and Metallicity Fitting 91 3.6 Discussion 108 4 Spectroscopic Age and Metallicity Gradients in Spiral Galaxies . . 112 4.1 Introduction 112 4.2 Longslit Spectroscopic Observations 114 Contents vi 4.2.1 Gemini GMOS 114 4.2.2 Galaxy Sample 115 4.2.3 Lick Standard Star Sample 118 4.2.4 Integration Times and Observing Strategy 122 4.3 Data Reduction 123 4.3.1 Characterization of CCD device(s) 124 4.3.2 High-Frequency CCD Effects and Defects 126 4.3.3 Low-Frequency Effects 128 4.3.4 Sky Subtraction 131 4.4 Data Analysis 139 4.4.1 Velocities 139 4.4.2 CCD Gap Locations 142 4.4.3 De-Redshifted Sky Lines 143 4.4.4 H II Region Masking 144 4.4.5 ID Spectra 152 4.4.6 Resolution and Velocity Dispersion Effects 153 4.5 Results from Indices 168 4.5.1 Line-Strength Profiles 168 4.5.2 Measured Indices Compared to Models 188 4.5.3 Age and Metallicity Fits 201 4.5.4 Colours Compared to Models 215 4.5.5 Age k Metallicity Fits from Colours 215 4.6 Comments on Individual Galaxies 221 4.6.1 Summary 231 5 Conclusions and Outlook 233 Bibliography 241 vii List of Tables 2.1 2MASS H - K Isophotal Colours 24 3.1 Lick Index Definitions 70 3.2 W097 Higher-Order Balmer Line Indices 72 3.5 Age & Z Fits and Errors as a function of fv: H/3 vs. Mgb & Mg 2 . . . 101 3.6 Age & Metallicity Fits and Errors as a function of fv: H/3 vs. [MgFe]' & (Fe) 102 3.7 Age & Z Fits and Errors as a function of fv: H/3 vs. G4300 & T i 0 2 . • 103 3.8 Age & Z Fits and Errors as a function of fv: H/3 & Dn(4000) vs. Fe4668 104 3.9 Age & Z Fits and Errors as a function of fv: ESA vs. (Fe) & Fe4668 . 105 4.1 Galaxy Sample 119 4.2 Galaxy Structural Parameters 122 4.3 Standard Star Sample 122 4.4 Observation sequences and conditions 134 4.5 Galaxy Rotation Velocities 141 4.6 CCD Gaps: Effective Wavelength Ranges and Indices Affected 143 4.7 De-redshifted Sky lines & Indices Affected 144 4.8 H II Region Emission Lines 145 4.9 Dispersion Relations for Resolution Functions 162 4.10 Bulge and Disk Scale Parameters 164 4.11 Bulge and Seeing FWHMs 165 4.12 Galaxy Ages and Metallicities 231 viii List of Figures 2.1 W50 vs. W20 for the Bottinelli (1990) and Theureau (1998) Samples . 25 2.2 W50 from Bottinelli (1990) vs. W21 from Haynes et al. (1991) and l / m a x from Courteau (1997) 26 2.3 Correlations between galaxy parameters 27 2.4 Correlation between galaxy parameters and morphological type . . . . 29 2.5 Comparison of the stellar population model tracks for different initial mass functions 32 2.6 Time evolution for the exponential and Sandage star formation histories 34 2.7 Comparison of model tracks for exponential and Sandage star formation histories 35 2.8 Comparison of the stellar population model tracks with bursts at 1 Gyr 37 2.9 Comparison of the stellar population model tracks with bursts at 4 Gyr 37 2.10 Bulge versus Disk Colours 42 2.11 Near-IR - optical colour-colour plots for the Courteau et al. sample sep-arated by morphological type 43 2.12 Near-IR - optical colour - colour plots for the Courteau et al. sample sep-arated by rotational velocity 44 2.13 Near-IR-optical colour - colour plots separated by rotational velocity for the BdJOO sample 44 2.14 Near-IR - optical colour - colour plots for the Courteau et al. sample sep-arated by iJ-band magnitude 45 2.15 Near-IR - optical colour - colour plots for the Courteau et al. sample sep-arated by H-band central surface brightness 45 2.16 Examples of Monte Carlo realizations for age and metallicity fits . . . . 49 List of Figures i x 2.17 Average age and metallicity as a function of local Hm and if-band surface brightness 50 2.18 "Effective" average age as a function of central surface brightness and total Hm, if-band galaxy magnitude 51 2.19 "Effective" average age as a function of scale length and rotational velocity 52 2.20 "Effective" metallicity as a function of central surface brightness and total Hm, if-band galaxy magnitude 52 2.21 "Effective" metallicity as a function of scale length and rotational velocity 53 2.22 Age as a function of radius 54 2.23 Metallicity as a function of radius 54 2.24 Age gradient as' a function of central surface brightness 55 2.25 Age gradient as a function of ii m , if-band central surface brightness and total magnitude . .' 57 2.26 Age gradient as a function of disk scale length and rotational velocity . 58 2.27 Metallicity gradients as a function of ii m , if-band central surface bright-ness and total magnitude 59 2.28 Metallicity gradients as a function of disk scale length and rotational velocity 59 2.29 Metallicity gradient as a function of age gradient 60 2.30 Effective age and metallicity as a function of morphological type . . . . 62 2.31 Age and metallicity gradients as a function of morphological type . . . 62 2.32 Age and metallicity gradient histograms as a function of barredness . . 63 3.1 Colour excesses E(B — V) and E(V — K) as a function of f y for adopted dust models 78 3.2 Comparison of solar metallicity SEDs with fy = 0, 3, & 6 at ages of 0.5 & 13 Gyr for SSP and exponential SFH models 79 3.3 Lick and W097 index differences as a function of fy for solar metallicity SSPs 81 3.4 Same as Figure 3.3 but for for an exponential SFH 84 List of Figures x 3.5 Comparison of the H/3 index measurement for different dust content . . 86 3.6 Same as Figure 3.5 but for Mgi 87 3.7 D(4000) and Dn(4000) differences as a function of f y for SSPs and ex-ponential SFHs 88 3.8 Calcium triplet index differences as a function of f y for solar metallicity SSPs 90 3.9 Same as Figure 3.8 but for for an exponential SFH 91 3.10 Rose index differences as a function of f y for solar metallicity SSPs . . 92 3.11 Same as Figure 3.10 but for an exponential SFH 93 3.12 Examples of the Monte Carlo fits for D„(4000) vs. Fe4668 diagnostic plot 96 3.13 H/3 vs. Mg6 and Dn(4000) versus Fe4668 diagnostic plots 97 3.14 (Fe) vs. Mgb diagnostic plot 100 3.15 H/3 vs. Mg 2 , H/5 vs. [MgFe]', H/3 vs. Fe4668, and R8A vs. Fe4668 diag-nostic plots 106 3.16 H£ vs. G4300, H/3 vs. T i 0 2 , BSA vs. (Fe), and H<5A vs. Fe4668 diagnostic plots 107 3.17 H/3 vs. (Fe) diagnostic plot for dust-free exponential SFH models for metallicities Z = 0.0004, 0.02, and 0.05 108 3.18 B.5A vs. Dn(4000) diagnostic plot for dust-free SSPs and exponential SFH for solar and Z = 0.05 metallicity 110 4.1 Sample GMOS longslit observation 115 4.2 Efficiency curve for the B600-G5303 grating on GMOS-N 116 4.3 Observational Set-Up 120 4.4 Observational Set-Up - continued 121 4.5 Sky brightness as a function of the time to twilight or dusk 135 4.6 Pseudo-rotation curves for each galaxy. 140 4.7 Example of an H II Region 146 4.8 Ha Masking for NGC 173 & 628 148 4.9 Ha Masking for UGC 2124 & 7490 149 List of Figures xi 4.10 Ha Masking for NGC 7495 & NGC 7610 150 4.11 Ha Masking for NGC 7741 & IC 239 151 4.12 Spectra as a function of radius for NGC 173 154 4.13 Spectra as a function of radius for NGC 628 155 4.14 Spectra as a function of radius for UGC 2124 156 4.15 Spectra as a function of radius for NGC 7490 157 4.16 Spectra as a function of radius for NGC 7495 158 4.17 Spectra as a function of radius for NGC 7610 159 4.18 Spectra as a function of radius for NGC 7741 160 4.19 Spectra as a function of radius for IC 239 161 4.20 Comparison of Model Grids at Different Resolutions 167 4.21 Index values as a function of radius for NGC 173 173 4.22 Index values as a function of radius for NGC 628 175 4.23 Index values as a function of radius for UGC 2124 177 4.24 Index values as a function of radius for NGC 7490 179 4.25 Index values as a function of radius for NGC 7495 181 4.26 Index values as a function of radius for NGC 7610 183 4.27 Index values as a function of radius for NGC 7741 185 4.28 Index values as a function of radius for IC 239 187 4.29 Comparison of SSP vs. EXP SFH Index-Index Grids 189 4.30 Index-index plots for NGC 173 192 4.31 Index-index plots for NGC 628 193 4.32 Index-index plots for UGC 2124 194 4.33 Index-index plots for NGC 7490 195 4.34 Index-index plots for NGC 7495 196 4.35 Index-index plots for NGC 7610 197 4.36 Index-index plots for NGC 7741 198 4.37 Index-index plots for IC 239 199 4.38 (Fe)-Mgfc diagnostic plot for all 8 galaxies . . 200 List of Figures xii 4.39 Age & Metallicity fits for NGC 173 207 4.40 Age & Metallicity fits for NGC 628 208 4.41 Age & Metallicity fits for UGC 2124 209 4.42 Age & Metallicity fits for NGC 7490 210 4.43 Age & Metallicity fits for NGC 7495 211 4.44 Age & Metallicity fits for NGC 7610 212 4.45 Age & Metallicity fits for NGC 7741 213 4.46 Age & Metallicity fits for IC 239 214 4.47 Colour-colour plots for UGC 2124 216 4.48 Colour-colour plots for NGC 7490 216 4.49 Colour-colour plots for NGC 7495 217 4.50 Colour-colour plots for NGC 7610 217 4.51 Colour-colour plots for NGC 7741 . . 218 4.52 Age and metallicity fits from indices and colours for UGC 2124 218 4.53 Age and metallicity fits from indices and colours for NGC 7490 219 4.54 Age and metallicity fits from indices and colours for NGC 7495 219 4.55 Age and metallicity fits from indices and colours for NGC 7610 220 4.56 Age and metallicity fits from indices and colours for NGC 7741 220 5.1 Best Fit Model Spectrum 238 Acknowledgements Xll l First and foremost, I would like to extend my eternal gratitude to my supervisor, Stephane Courteau. He has stuck with me through my five years as a graduate student resulting in two degrees (Masters and PhD). My love for astronomy began and flour-ished thanks to his enthusiasm and encouragement. Another source of great inspiration to me over the past five years is Eric Bell, who never tired of my endless questions during and after our collaboration on Chapter 2 of this thesis. A big shout-out to Jesus Gonzalez, without whose spectroscopic data reduction expertise Chapter 4 would not have been possible. Thanks also for teaching me your unique approach to science (Gris-som & Horatio would approve!) Many thanks go out to my collaborators in external projects; Jim Rose, Ricardo Schiavon, Scott Trager, and Guy Worthey, for allowing me to expand my scientific horizons and for their careful readings of early drafts of the pa-per resulting from Chapter 3 of this thesis. I would also like to extend my appreciation to the members of my thesis committee; Paul Hickson, Jaymie Matthews, and Tom Mattison, who provided useful comments throughout the research and were extremely supportive of my work. My external examiners; Robert Kennicutt (ApJ editor-in-chief, Steward Observatory), Douglas Scott (UBC Physics & Astronomy) and Rajiv Gupta (UBC Mathematics), are also thanked for their useful remarks and enthusiastic support. Cheers also to all of the astronomers I have met along the way. I sincerely hope that our paths will cross again soon! Finally, to my friends and family, words cannot convey my appreciation for the role you have all played in my life. Any success and happiness that comes my way is largely thanks to you all. xiv List of Acronyms For easy reference, presented here is a list of the acronyms used throughout this thesis along with their definitions. 0.1 List of Acronyms 2MASS Two Micron All Sky Survey (http://www.ipac.caltech.edu/2mass/) 2dFGRS Two Degree Field Galaxy Redshift Survey (http://msowww.anu.edu.au/2dFGRS/) A D U Analog-to-Digital Unit A G N Active Galactic Nucleus A G B Asymptotic Giant Branch B / D Bulge-to-Disk Ratio CCD Charge Coupled Device C D M Cold Dark Matter C M B Cosmic Microwave Background COBE COsmic Background Explorer (http://aether.lbl.gov/www/projects/cobe/) CSB Central Surface Brightness (/i0) DEC Declination DEEP Deep Extragalactic Evolutionary Probe (http://deep.ucolick.org/) List of Acronyms xv D M Dark Matter FIR Far-Infrared (15-170/mi) FOV Field Of View F W H M Full Width at Half Maximum GC Globular Cluster GMOS Gemini Multi-Object Spectrograph GOODS The Great Observatories Origins Deep Survey (http: //www.stsci.edu/science/goods/) IC Index Catalogues (supplements to the NGC) IMF Initial Mass Function IR Infrared wavelength (~ 1/im-1 mm) IRAF Image Reduction and Analysis Facility (http://iraf.noao.edu/) H D M Hot Dark Matter HSB High Surface Brightness ( /X 0 ,B ^ 22 mag arcsec-2) HST Hubble Space Telescope (http://hubble.nasa.gov/index.php) KPNO Kitt Peak National Observatory LSB Low Surface Brightness (/i0)s ~ 22.5 mag arcsec-2) M W Milky Way NGC New General Catalog NIR Near-Infrared wavelength (~ 1 -10 fj,m) RA Right Ascension List of Acronyms xvi R M S Root-Mean-Square S / N Signal-to-Noise Ratio S A M Semi-Analytic Model S B Surface Brightness S D S S Sloan Digital Sky Survey (http://www.sdss.org/) S E D Spectral Energy Distribution S F Star Formation S F R Star Formation Rate S F H Star Formation History S I N G S The Spitzer Infrared Nearby Galaxies Survey (http://sings.stsci.edu/) S N ( e ) Supernova(e) S P S Stellar Population Synthesis S S P Simple Stellar Population T P - A G B Thermally-Pulsing Asymptotic Giant Branch U G C Uppsala Galaxy Catalog U V Ultra-Violet wavelength (10nm-l/im) X V I S T A Interactive Image and Spectral Reduction and Analysis (Package developed at the Lick Observatory and maintained at http://astronomy.nmsu.edu/holtz/xvista/) 1 Chapter 1 Introduction 1.1 Galaxy Formation and Evolution The principal motivating theme underlying this thesis is to understand how the mul-titude of galaxies in our Universe formed and evolved into the systems we observe today. It has long been a puzzle how the early, homogeneous, universe evolved into the complex structures we see today. This overall cosmic evolution conundrum has been compounded by a global dearth of empirical information about evolutionary ob-servables from the dark ages until now, for both the dark and luminous components of cosmic structures. Among the many missing pieces to this puzzle, stellar populations in galaxies are tell-tale indicators of the evolution of the baryonic components at all ages. While we focus on fully evolved (virialized), nearby (and therefore spatially resolved) galaxies in this study, a full understanding of the broad subject of galaxy formation and evolution requires that we start from the very early universe. A general outline of the currently favored paradigm is provided here. There have been dramatic new discoveries from measurements of the cosmic mi-crowave background (CMB) that have enabled scientists to detect a small, but non-zero, anisotropy indicative of tiny fluctuations in the density of the early universe. These asymmetries led to gravitational centers, seeding the accumulation of matter that would eventually account for all of the large scale structure (number and distribu-tion of galaxies in space and time) we see in the universe today. There has also been great progress in understanding the matter and energy distribution in the universe. Information from the rotation curves of galaxies, supernovae, and measurements of the curvature of the universe has led us to a remarkably consistent picture. However, in this picture, only about 4% of the universe can be accounted for by ordinary baryonic Chapter 1. Introduction 2 matter; the protons, neutrons, atoms, and molecules that make up our world. A further 23% is made up of exotic particles we have yet to identify, the so-called dark matter (DM), and the remaining 73% is a mysterious dark energy, also not understood. Clearly the formation and evolution of galaxies will be affected by the presence of a substantial dark matter component. With the advent of new data, a greater understanding of the matter/energy density of the universe, and the initial conditions, it is imperative that we better understand the nature of the galaxies we see today. This will allow an accurate comparison with older galaxies to expose evidence for evolutionary changes. In this thesis, we develop techniques to characterize the stellar populations in spiral galaxies, with an eventual aim of being able to understand their formation and evolution. With ever-improving observational technologies, ambitious new surveys, and very large data bases becoming available (SDSS, 2MASS, 2dFGRS, GOODS, DEEP) there is great potential for major advances in this field. In order to gain an appreciation for the outstanding questions regarding the formation of galaxies and their evolution into the structures we see today, we present here a brief historical account of the field. The foundation for theories of structure formation in the universe dates back to the late 1920s with the first studies of perturbations in a dust-filled Friedmann-Lemaitre universe (homogeneous and isotropic with a cosmological constant) (de Sitter 1930; Lemaitre 1931). The general picture involves the amplification, due to gravitational in-stability, of small inhomogeneities in an otherwise homogeneous distribution of matter. But a full treatment of linear perturbation theory (Peebles 1980; Bardeen 1980), led back to the starting point: what are the initial amplitude and spectrum of the density fluctuations necessary to seed galaxy formation? By the late 70's, Alan Guth had begun developing his inflationary universe theory which provided theoretical predictions for the origin and spectrum of primordial density fluctuations. Inflation does not, however, offer a prediction for the amplitude of the fluc-tuations. An empirical measurement of the amplitude and spectrum of the primordial density fluctuations was provided by Smoot et al. (1992) based on microwave data from Chapter 1. Introduction 3 the Cosmic Background Explorer satellite (COBE). COBE measured the temperature anisotropics in the cosmic background radiation to be (AT/T)rms « 1.1 ± 0.2 x 10~5 which provided the necessary normalization to the structure formation models. Addi-tionally, the spectrum measured by COBE provides an excellent match to the scale-invariant spectrum of density perturbations predicted by inflation 1 . Equipped with a plausible inflationary theory and powerful computers, cosmolo-gists in the 1990s were then in a much better position to study the evolution of the density fluctuations of the early universe. Cosmological N-body simulations, in which the matter distribution is described as a distribution of N particles that interact via gravity, provided a direct method of following the (non-)linear evolution of the density perturbations. However, given that the matter density of the universe is dominated by non-baryonic (dark) matter, whose physical nature and composition remains elusive, the simulations must adopt a basic framework providing a prescription for the dark matter particles. The two most extensively explored theories are referred to as the hot and cold dark matter models (HDM and CDM, respectively). In HDM models, the dark matter particles are light, fast moving (relativistic), and interact only via the weak force. A primary motivation for the HDM models was the existence of a known candidate; the neutrino. CDM particles on the other hand are collisionless and have negligible random velocities. No known candidates for the CDM particle have been detected to date but a number of search efforts are under way (Primack 2000; Cline 2004). Both the HDM and CDM theories have encountered a number of difficulties, but only the CDM models have stood up to most of the observational constraints available to date. Major weaknesses of the HDM models result from their lack of power on small scales. Accordingly, structure on the scale of galaxies can only form after the collapse of cluster-sized objects, in a "top-down" scenario, by means of fragmentation. Simula-tions of HDM universes with reasonable values for Z f o r m , the redshift at which galaxies 1 While the inflationary theory still remains somewhat speculative, and a myriad of alternative theories have been suggested, the remarkable agreement between the C O B E measurements and its predicted spectrum along with its solutions to the long standing flatness, horizon, and monopole problems, make the inflationary universe theory difficult to dismiss. Chapter 1. Introduction 4 begin to collapse, cannot be reconciled with the current measurements of the galaxy correlation function and observed clustering. Based on pure HDM models, galaxy for-mation should occur on timescales too close to the present time. Additionally, the COBE measurements of anisotropy in the CMB rule out DM models in which free-streaming wipes out small scale power, which is a characteristic of HDM particles. In contrast, CDM models have the largest amplitudes on smallest scales which results in a hierarchical, "bottom-up", process of structure formation: large structures form out of smaller ones which collapse first. CDM models have also faced their share of inconsistencies when comparing predictions with observations2, but their overall success at reproducing many aspects of the large scale structure of the universe make them the preferred choice for cosmologists at the present time. In particular, CDM model predictions of the clustering of gas and galaxies as a function of redshift have been tested and verified observationally. Attempts have been made to extend the above models to the formation of individ-ual galaxies, taking their results as initial conditions. Formation models begin with a "proto-galaxy" in the form of an isolated, uniformly rotating, sphere composed of pri-mordial gas and dark matter. The mass and angular momentum distribution of these spheres are taken from the above N-body simulations. The initial angular momentum of each proto-galactic sphere is acquired through tidal torquing from neighboring sys-tems (Peebles 1969; Fall & Efstathiou 1980). The assumption of uniform rotation is based on an early observation that the angular momentum distribution of the galactic disk closely resembles that of a sphere in solid-body rotation (Mestel 1963). Another assumption adopted in these models is that angular momentum transport during the collapse is negligible. This assumption may hold for dissipationless, dark matter domi-nated systems, which is true of low surface brightness systems (LSBs) (Dalcanton et al. 1997), but in systems that are dominated by baryons, viscous dissipation must certainly play a major role in the evolution of the galaxy. Nevertheless, these simple models are 2 A s of this writing, the principal challenges faced by C D M models are the confrontation of cosmologically-motivated halo density profile shapes at the center of galaxies based on high resolution rotation curves, angular momentum in halos and galaxies, and dark matter substructure (Primack 2004a,b). Chapter 1. Introduction 5 able to reproduce many of the universal properties of low and high surface brightness disk galaxies; exponential disk profiles with a wide range of scale lengths, asymptoti-cally flat rotation curves, and the Tully-Fisher relation between circular velocity and luminosity of spiral galaxies, but they do not simultaneously reproduce the observed local luminosity function. When gravitational instabilities are considered in models which incorporate viscous dissipation, galaxies often develop a bulge-like component whose properties also depend on the total angular momentum of the system (Saio & Yoshii 1990; Struck-Marcell 1991). However, the modeling of galaxies has been hampered by our limited knowledge of processes involving angular momentum transfer between the halo and disk, and star formation and feedback mechanisms in a multi-phase interstellar medium. Lacking detailed prescriptions for these fundamental astrophysical processes, it is difficult to develop a self-consistent model for the dynamical and chemical evolution of individual galaxies. Another class of models, referred to as semi-analytical (SAM) models, have attempted to incorporate analytical parameterizations of both well-constrained and elusive astrophysical processes into numerical simulations (Somerville & Primack 1999; Cole et al. 2000; Baugh et al. 2005). These require a number of adjustable parameters, which are fixed by reference to a set of local galaxy data. Based on the properties of present-day galaxies we can attempt to roll back the clock and predict how galactic structures may have appeared in the past. SAMs, however, have had limited success at reproducing observations that were not set by the tuning of parameters, and at making new hard predictions. 1.1.1 Bulge Formation Scenarios Also of crucial importance to the study of the formation and evolution of disk galaxies and spheroids in general is a scenario for the formation of the bulge. A successful galaxy formation theory must determine the formation timescales for bulges and disks and reproduce the observed stellar ages and chemical distributions in galaxies as a function of galaxy properties. Current theories for the formation of disk bulges generally fall under two basic pictures: monolithic collapse (e.g. Eggen, Lynden-Bell & Sandage Chapter 1. Introduction 6 1962; Larson 1974; Carlberg 1984) and hierarchical merging (e.g. Kauffmann 1996; Cole et al. 2000). In the monolithic collapse model, primordial gas clouds collapse under their own self-gravity in the very high redshift universe. Early N-body simulations (Larson 1974) demonstrated that while the gas collapses, it loses energy due to dissipative processes and falls towards the center of the galaxy. As the gas sinks, it accumulates enriched gas from the local stellar population. Consequently, successive stellar generations not only form closer to the center, but are also increasingly rich in metal content3. Thus, for spheroidal stellar systems (i.e. elliptical galaxies and spiral bulges) the monolithic dissipative collapse of a gas cloud at high redshift is consistent with strong metallicity gradients (on the order of AlogZ/Alog(r pc)~0.5 dex/decade within the central ~ 2 kpc, but flattening in the outer parts, Larson 1974), with metallicity decreasing with increasing radius, shallow positive age gradients (the centers are slightly younger than the outer parts), and old stellar populations. This result was later verified (Carlberg 1984), and it was further demonstrated that the violent relaxation that occurs in such a collapsing system naturally produces a specific shape for the light profile of the spheroid (the so-called "r 1/ 4" or "de Vaucouleurs" profile), as well as dynamical support from random motions (velocity dispersion). In this "monolithic" scenario, the disk would accrete onto the already formed bulge, and thus be younger than the mean bulge age, and their structural relations would be uncorrelated. On the other hand, in the hierarchical collapse models, spheroid formation results from the interaction and/or merging of smaller fragments (likely disk galaxies) in which star formation has already begun. This violent merging again produces the r 1 / , 4-law profiles and velocity dispersion dynamical support, but any stellar population gradients 3 A "metal" here refers to any element heavier than Helium. The "metallicity" of a system is often expressed in terms of fractional abundance by weight and denoted Z (and X & Y refer to hydrogen and helium respectively, so X + Y + Z = 1), where the Sun (denoted by the symbol ©) has ZQ = 0.02. Alternatively, metallicity is expressed as the logarithmic abundance of metals by number relative to hydrogen and compared to that of the Sun and denoted [Z/B] = \ogw{n{Z)/'n(H)) -l o g 1 0 ( n ( Z ) / n ( H ) ) © . Since iron has the largest fractional contribution to the total metallicity of any other element that can be measured in absorption, the metallicity of a stellar system is often given as [Fe/H]. If the system of interest has solar abundance ratios of all elements, then [Z/H] = [Fe/H]. Chapter 1. Introduction 7 present in the merging objects will be flattened or seriously altered by a burst of central star formation due to gas in the merging objects being funneled to the center (Barnes & Hernquist 1996). The dominant signatures in the hierarchical scenario depend on the time evolution of the merger rate: if mergers were most frequent at high redshift, as predicted by present-epoch cosmological models, hierarchical signatures should be similar to those of dissipative collapse. More recent mergers and tidal disruptions would result in flatter metallicity gradients, and a significant fraction of young or intermediate-age stars. Thus, the primary difference between the above two scenarios relates to the relative ages of spheroidal systems. For example, if bulges formed via monolithic collapse at high redshift, the bulges we see today would be uniformly old, and significantly older than their corresponding disks. Additionally, the subsequently forming disk would not have a strong influence on the bulge component. Accordingly, no structural correlation would be expected between the bulge and disk components. However, in the hierarchical scenario, while the majority of spheroids form at high redshift where the merger rate is highest, mergers continue to the present day and thus spheroids should have a large range of ages (e.g. Kauffmann 1996). These models further predict that spiral bulges will form later than ellipticals (and thus should have an age distribution that does not extend as old as the oldest ellipticals), and a correlation between bulge mass and age such that later-type bulges should form earlier. Finally, another category of models for the formation of spiral bulges has been proposed, the so-called "secular evolution" models (see Kormendy & Kennicutt 2004 for an excellent review). In these models, spiral bulges form as a result of angular momentum redistribution due to disk instabilities. Models including viscous transport find that the efficiency of transporting disk material into the central regions is greatly enhanced by a bar or oval distortion, which can be triggered by a global dynamical instability in the disk induced by interaction with a satellite (Martinet 1995). Disk material gets heated vertically up to 1 - 2 kpc above the plane via resonant scattering of stellar orbits by the bar-forming instability. A "bulge-like" component with a nearly Chapter 1. Introduction 8 exponential profile will emerge due to relaxation induced by the bar. Gas redistribution by the bar can cause its own dissolution. Secular accumulation or satellite accretion of only 1-3% of the total stellar disk mass near the center is sufficient to induce dissolution of the bar into a lens or triaxial component and later into a spheroid (e.g. Kormendy 1982; Pfenniger 1993; Friedli et al. 1994; Norman, Sellwood, k Hasan 1996). It is estimated that about two-thirds of disk galaxies currently have a bar, especially as revealed in the infrared (Zaritzky, Rieke, k Rix 1993; Sellwood k Wilkinson 1993; Martin 1995) and that most spirals have probably harbored a self-destructive bar at one time or another during their evolution (Friedli & Benz 1993; Sellwood k Moore 1999), thus lending some support for bulge formation models of secular evolution. Note, however, that the number of bar formation - gas inflow - star formation - bar dissolution cycles that can occur in a galaxy may be limited, as the central mass formed in the first cycle can prevent subsequent formation of a bar (Wyse 1999)4. Radial migration of disk material (stars and gas) can also be induced by non-axisymmetric spiral waves (Sellwood k Binney 2002), which would necessarily flatten metallicity gradients in disks. In general, the secular evolution scenario predicts that bulges will be similar in metal-licity and age to their (parent?) disks, with flattened gradients compared with mono-lithic collapse predictions, and will be structurally linked to their disks (e.g. Courteau, de Jong, k Broeils 1996). Evidence for the latter in late-type spiral galaxies includes the continuation of spiral structure into center of galaxies and structural and photo-metric correlations between bulge and disk parameters (e.g. MacArthur, Courteau, k Holtzman 2003). Secular evolution is a viable mechanism for producing the small, central accumula-tions of material in late-type disks. Bigger bulges, however, could not be formed this way without disrupting the disk. The energy required to heat up the central material is far greater than the total bar and disk's mechanical energy supply. Accretion of a 4 Only progressively larger bars in the centers of exponential bulges would be allowed to form in a recurring scenario as a result of the disrupting dynamical effect of a growing nucleus (H.-W. R ix 1999, private communication, as reported in Carollo 1999). Chapter 1. Introduction 9 small satellite to explain the bigger bulges of SO - Sa's provides an appealing alternative (Pfenniger 1993; Walker, Mihos & Hernquist 1996; Balcells et al. 2003). It is important to note that it is entirely possible that all of the above mentioned scenarios occur in nature. The goal is thus to determine the relative importance of each scenario and their respective influence on the formation of the entire Hubble sequence. One might expect that, at early times the principle processes were violent and a com-bination of dissipative collapse and merging, while today, internal secular processes are dominant. The current weaknesses in model constraints are largely due to the paucity of data for early and - especially - late-type bulges (Zhang & Wyse 2000; Kormendy & Kennicutt 2004), a predicament we aim to begin to resolve in this thesis. 1.2 Observational Constraints Current observational results (O/H gradients, SFRs, constant IMF, gas distribution) favor an inside-out formation for the Milky Way, but the details about age and metal-licity gradients in our own Galaxy are conflicting (Freeman & Bland-Hawthorn 2002; Kormendy & Kennicutt 2004). Integrated studies of stellar populations in external galaxies can restrict the range of possible interpretations, but formation timescales and chemical evolution of external spiral galaxies have barely been explored (e.g. Matteucci 2002). Detailed information on stellar populations, as tracers of initial conditions, in nearby spirals is scanty and relies mostly on colours and emission diagnostics. However, colours are plagued by a degeneracy between dust and age/metallicity effects (Bell & de Jong 2001; MacArthur et al. 2004 [Chapter 2]) and emission lines are sensitive only to the last episode of star formation. These issues have become even more important now that cosmologically-motivated numerical simulations are capable of resolving the formation of disks and bulges. Furthermore, semi-analytic galaxy evolution models use recipes to describe bulge and disk evolution, but if different types of galaxies have different bulge formation mechanisms, these models may give misleading results (e.g. Somerville & Primack 1999; van den Bosch 2000), unless a variety of bulge formation mechanisms are modeled (e.g. Cole et al. 2000). Chapter 1. Introduction 10 1.2.1 Colour Gradients in Spiral Galaxies Existing studies of spiral bulges and disks have placed tentative constraints on the source of their colour gradients and which galaxy parameters most strongly correlate with SFHs and, hence, on specific formation scenarios, but some of the results are conflicting. In a study of colour gradients in early-type bulges (< Sb) using HST and ground-based data, Peletier et al. (1999) inferred that early-type bulges are old (absolute ages not well determined) with a small age spread (< 2 Gyr) for all early-type bulges measured at the bulge effective radius. The old and narrow age range makes it unlikely that these bulges would have formed via secular evolution. The same conclusion does not hold for their (few) later-type galaxies. Peletier et al. (1999) conclude that the intrinsic colour gradients of early-type bulges are caused mainly by metallicity gradients (in agreement with other similar studies; e.g. Mehlert et al. 2003), consistent with the monolithic collapse scenario. However, based on the lack of r 1 / 4 shaped bulges found in a later study of bulge SB profiles of SO - Sbc galaxies using HST near-IR high resolution imaging, Balcells et al. (2003) have revised their interpretation and conclude that these bulges could not have formed from violent relaxation in mergers, but rather are more consistent with secular formation processes. On the other hand, in an HST study of spatially resolved colours of high redshift (z ~ 0.5) galaxies in the Hubble Deep Field, Abraham et al. (1999) rule out metallicity gradients as major contributors to galaxy colour distributions. They do find, in agreement with Peletier et al. (1999), that large bulges are significantly older than their disks and therefore rule out secular evolution formation processes in favor of a gradual disk formation by accretion of gas. However, the flat metallicity distribution of the bulge is not easily explained in this scenario (particularly since the lack of "r 1/ 4" bulges rules out major merging as a cause of gradient flattening). The most comprehensive study of colour gradients in disk galaxies to date is that of Bell & de Jong (2000, hereafter BdJOO), who compute age and metallicity gradients for a sample of 121 nearby low-inclination SO-Sd galaxies. They conclude that galaxy colour gradients are due in most part to gradients in age and metallicity in their stellar Chapter 1. Introduction 11 populations (in the sense that inner regions are older and more metal rich than outer regions) and contend that dust reddening mainly affects the metallicity gradients, but is likely too small to affect their conclusions. In comparing the SFHs with galaxy parameters, they conclude that the SFH of a galaxy is primarily driven by surface density and that the total stellar mass of a galaxy is a less important parameter that correlates significantly with metallicity, but not age. Kauffmann et al. (2003a,b) also conclude that the recent SFHs, as probed by the K5A absorption line index and the 4000 A break5, of over 100,000 galaxies are more strongly correlated with surface mass density than stellar mass. BdJOO argued that these correlations could be the result of a surface density-dependent star formation law, coupled with galaxy mass-dependent chemically enriched gas outflows. Bell & Bower (2000) further explored this idea by constructing a family of simple models for spiral galaxy evolution for comparison with the observational trends in SFH with galaxy parameters. Indeed they found that the data are consistent with the proposition that the SFH of a region within a galaxy depends primarily on the local surface density of the gas but that additional ingredients, such as galaxy mass dependent infall (outflow) of primordial (metal-enriched) gas and/or formation epoch, are required to fully explain the observational results. Dust in Galaxies and the Importance of Multiple Passband Information Dust absorbs and scatters blue light and re-emits it at longer, far-IR wavelengths. The emission wavelength depends on the size and temperature of the dust grains. The overall effect of dust on the observed light distribution in a galaxy is strongly dependent on the dust geometry. In the simplest case of a homogeneous screen of dust, the observed galaxy is dimmed and reddened in the UV-to near-IR wavebands. However, for more complicated geometries where there is mixing between the stellar and dust components, the situation becomes much more complex as differential optical depths come into play. A comprehensive review of the current status of dust effects in galaxies is given by Calzetti (2001). A detailed picture of opacity and radiation transfer in galaxies is clearly not avail-5We will see in Chapter 3 that the 4000 A break may however be affected by dust extinction. Chapter 1. Introduction 12 able and multi-wavelength data sets, scanty at present, are of crucial importance for such studies. The near-IR wavebands, the K-band in particular, are largely free of dust extinction effects. Multi-wavelength data sets including optical and near-IR observa-tions, such as the one we present here in Chapter 2, provide a possible indicator of the dust content and distribution in galaxies. However, great care must be taken in the interpretation of multi-wavelength information as dust extinction and scattering, stellar population mix, age, and metallicity effects can all account for some of the differences seen in the light distributions of galaxies at different passbands. Studies using multi-wavelength data that solve for the intrinsic colour of stellar populations and reddening and extinction effects from dust simultaneously have measured typical central optical depths ry(O) in the range 0.3-2.5 for types Sab-Sc (Peletier et al. 1995; Kuchinski et al. 1998; Xilouris et al. 1999), while lower values of Ty(0) ~ 0.15 have been observed for low surface brightness late-type galaxies (Matthews & Wood 2001). Optical depths of this order can certainly have effects on the observed broadband colours. Given that the effects due to dust redenning will decrease as the baseline over which the measurement is made narrows, measuring individual absorption and emission features has the potential of breaking through the degeneracy caused by dust. Thus spectroscopic measurements could provide a complementary database to compare with the broadband photometry. No such database currently exists, a situation we aim to remedy beginning with a pilot sample presented in Chapter 4. 1.2.2 Spectroscopic Age and Metallicity Indicators Numerous efforts have been made to develop a more sensitive tool for probing compos-ite stellar populations than is possible with broadband photometry. In particular the Lick/IDS system of absorption-line indices (Gorgas et al. 1993) was designed to probe stellar population features that have differing sensitivities to population parameters. Line strength measurements of stellar systems can be translated into metallicities and ages using population synthesis models (e.g. Worthey 1994; Maraston 1998; Bruzual &; Chariot 2003), thus enabling a direct comparison with galaxy formation scenarios. The line-index system and population synthesis models are discussed in detail in later Chapter 1. Introduction 13 sections (§3.1-§3.3). We focus here on the major results obtained to date based on spectroscopic studies that have made use of the above technique to probe stellar pop-ulations of galaxies. Major efforts have been invested to determine mean ages and chemical compositions of galaxies, with a clear emphasis on elliptical galaxies whose composition is simpler than that of disks (e.g. Gonzalez 1993; Trager, Faber, Worthey & Gonzalez 2000; Caldwell, Rose & Concannon 2003). These studies have shown that the relative abundance of Mg to that of Fe, Mg/Fe, is enhanced in the centers of massive galaxies, and that lower mass ellipticals have younger luminosity-weighted mean ages. Since the so-called a-elements (i.e. O, Mg, Ca, Ti , and Si) are promptly released by massive, short-living (< 3x 107 yrs) progenitors exploding as type II supernovae (SNe), while most Fe comes from type la SNe, whose progenitors span evolutionary timescales from over ~ 3 x l 0 7 yrs to many Gyrs (e.g. Greggio & Renzini 1983), a high cv-to-Fe ratio ([a/Fe]) implies that star formation ceased before the bulk of type la SNe had enough time to enrich interstellar medium (ISM) with Fe. Such a short star formation timescale appears to be at variance with the predictions of current hierarchical models for the formation of elliptical galaxies (Thomas & Kauffmann 1999), which predict star formation to continue for several Gyrs. Indices centered on carbon-sensitive features ( C N i , C N 2 , Fe4668 - often referred to as C 2 4 6 6 8 due to its sensitivity to carbon) have also been observed to lie above the stellar population synthesis (SPS) model predictions (Henry & Worthey 1999). However, the interpretation of carbon-enhancement is difficult due to the many, and poorly understood, production sites for carbon. Elliptical galaxies have significant gradients in metal-line indices centered on Mg and Fe features, but the H/3 profile is flat. In terms of SPS models, this implies a negative metallicity gradient of order —0.3 dex/Alog(r) and a positive age gradient of 3 Gyr/Alog(r) (Gonzalez 1993). Such gradients could be consistent with either the monolithic or hierarchical scenarios. They also have a large spread in average age, with field elliptical having ages of ~ 4-6 Gyr and cluster ellipticals with ages of ~ 12-15 Gyr. Such an age spread is not consistent with the pure monolithic collapse scenario, Chapter 1. Introduction 14 but does agree with predictions of the hierarchical collapse models. It has also been suggested that the most massive ellipticals must be formed by violent merging, whereas the low-to-intermediate mass ellipticals have gradients driven by dissipation. In a study of SO galaxies, Fisher, Franx, & Illingworth (1996) found that SO galax-ies generally follow trends in central Mg and Fe line strengths similar to the behavior observed in elliptical galaxies. They studied a subsample of 9 highly inclined galaxies from which they inferred bulge and disk gradients along both major and minor-axis profiles (outside of the central region, the minor axis should be free from disk contami-nation, thus is representative of a pure bulge profile). Metal line strengths were seen to decrease with radius along the major and minor axes in the bulge-dominated central re-gions. However, at larger radii, the major-axis metal line strength profiles flatten while the minor-axis bulge profiles fall to lower values. By comparing with the Worthey (1994) SPS models, they found that these shallow gradient profiles in the disk correspond to an average metallicity gradient of A[Fe/H]/A(r//i) = —0.08 ± 0.06. These gradients are a factor of 2 - 3 smaller than those derived for the disks of late-type spiral galaxies from H II regions (Garnett 2002), and are consistent with previous investigations that showed a trend for disk metallicity gradients to decrease toward earlier Hubble types. The SO disks also show Mg/Fe ratios lower than those found in the central regions. The mean size of their bulge metallicity gradients is A[Fe/H]/Alog(r) = —0.7 ± 0.4, which is steeper than typical elliptical galaxy gradients. Fisher, Franx, & Illingworth (1996) conclude that their findings are best explained in terms of formation via dissipative collapse (or merging) at early times, but do not favor any type of secular evolution in the formation of SO bulges. Current studies of later-type spiral bulges are few, and those that exist have con-flicting results. Goudfrooij, Gorgas, & Jablonka (1999) presented longslit spectra along the minor axis of the bulges of a sample of 16 edge-on spiral galaxies with Hubble types SO/a-Sc. Their results imply that stellar populations in bulges are old, with ages similar to cluster ellipticals (which, as mentioned above, are the oldest between field and cluster Es), and encompass a range of metallicities with later types having Chapter 1. Introduction 15 lower metallicity. They find that the bulge a/Fe element ratio is typically super-solar, also similar to those found in elliptical galaxies. They conclude that their findings are more compatible with predictions of the dissipative collapse model (either monolithic or heirarchical) than with those of secular evolution models. On the other hand, Proctor &; Samson (2002) present a spectroscopic analysis of the bulges of 32 edge-on galaxies with Hubble types ranging from E-Sbc. Using longslit spectra of disk bulges collected with the Palomar 5-m telescope, they measure the full suite of 24 Lick indices to constrain the ages, metallicities, and abundance-ratios of the bulge stellar populations using models that take into account varying abundance ratios6. They find that bulges are less a-enhanced and have lower average age than early-type galaxies (Es and SOs), in stark contrast with the Goudfrooij et al. (1999) findings. Also in contrast to Goudfrooij et al. (1999), they rule out primordial collapse models of galaxy formation (for all spheroids) in favor of the hierarchical models. Additionally, they find large differences between derived ages and parameter correlations between early and late-type bulges suggesting a difference in their SFHs, and that disk inflow must play an important role in the SFH of late-type bulges. It is unclear where the discrepancy arises for the two studies, but we note that the Goudfrooij et al. (1999) results are said to be preliminary and they reserve final conclusions for the final work (not published to date). Additionally, their results are based on a small number of indices which, as we will demonstrate in Chapter 4, can be dangerous, particularly for spiral galaxies which could suffer from emission line contamination. The only study to date presenting radial age and metallicity indicators for face-on late-type bulges and inner disks is the preliminary work of Trager, Dalcanton, & Weiner (1999), who present results for two early-type (Sab with central velocity dispertion ovd > 100 kms - 1 ) and two late-type (Sc with ovd < 100 kms - 1 ) spirals. They find that the early-types bulges are consistent with being old and metal rich (Age > 10 Gyr, 6 I t is relevant to point out, for our discussion in Chapter 4, that the spectra from Proctor & Samson (2002) only cover the bulge and do not reach into the disk. Thus neither the potential disk contamination nor the continuity of bulge to disk1 features can be addressed in this study. Chapter 1. Introduction 16 [Fe/H]>0) with bulges older than their disks by several Gyr, suggesting a bulge-first formation scenario. On the other hand, the later-type bulges are younger, more metal-poor, and closer in age to their inner disks, supporting a disk-first scenario. They do point out, however, that emisssion-line and disk contamination could bias the results. Another interesting study worth note is that of Bergmann, J0rgensen, & Hill (2003) who studied stellar populations in a sample of low surface brightness (LSB) galaxies using a combination of emission-line (EW(Ha), [N II] A6584/Ho:) and absorption-line (Mgb, H/3, (Fe)) based diagnostics. Realizing the need for the largest possible aperture to identify faint absorption features, they used the GMOS spectrograph on Gemini North (8-m telescope) to obtain longslit spectroscopy of their LSB sample. Bergmann, Jorgensen, & Hill (2003) find that the formation and chemical enrichment histories of these galaxies are diverse, with some galaxies having low metallicity and very young mean stellar ages while others show old, super-solar metallicity stellar populations. Correlations between several of the gas-phase and stellar population age and metallicity indicators are used to argue that the star formation history in LSB galaxies must have been fairly smooth, and not bursty. Their spectra are not radially resolved, thus their results are representative of the integrated (bulge + some disk) population. 1.2.3 Structural Components of Disk Galaxies Finally, we note that structural signatures, such as bulge and disk shapes (power-law index; scale lengths; truncation radii), may also provide direct links to formation timescales (Courteau, de Jong, &: Broeils 1996; Freeman Sz Bland-Hawthorn 2002). The most extensive structural studies of nearby spirals in the optical and infrared broadbands (de Jong 1996; MacArthur, Courteau, & Holtzman 2003) find definitive bulge/disk correlations and sharp truncation radii which may be linked to the angular momentum properties in the early proto-cloud (Freeman & Bland-Hawthorn 2002). A major open-question in disk modeling is the origin of Freeman Type-II profiles (which show strong depressions in the stellar surface density at the bulge/inner disk transition); interpretations range from dust opacity, to bar-induced stellar rings, and genuine inner disk truncation. The latter suggestion has enjoyed sufficient attention lately to warrant Chapter 1. Introduction 17 a careful investigation (Freudenreich 1998; Andersen et al. 2002). If true, inner disk truncation would represent a major departure from standard dynamical models. Surface brightness profiles of barred and unbarred galaxies both exhibit Type-I and Type-II signatures (MacArthur, Courteau, & Holtzman 2003) and the dust-only in-terpretation for Type-II profiles cannot be verified due to the low spatial resolution of far-infrared maps. The power of absorption spectra to disentangle dust, age, and metal-licity effects is needed to unravel the stellar populations occupying the central regions in Type-I and Type-II of barred and un-barred galaxies and elucidate the Freeman-type enigma. The same data can also trace population gradients as a function of barredness. The general picture emerges that larger bulges may have formed via mergers or dis-sipational collapse with no secular component, while smaller bulges are more consistent with having formed secularly. Many discrepancies, however, do exist and further obser-vational constraints are needed to confirm or contest this general view. In particular, as recent reviews on the formation of spiral bulges attest (e.g. Wyse 1999; Kormendy & Kennicutt 2004; Peletier et al. 2005), constraints on the stellar populations in the bulges and inner disks of late-type spirals are the most crucial missing piece to the puzzle of bulge formation. 1.3 Thesis Outline In order to extend to spiral galaxy bulges and disks the modern studies of abundances and kinematics in spheroids we have embarked on a systematic survey of bulge and disk absorption-line and emission features (line indices, equivalent widths) using a most efficient longslit spectrograph (GMOS: see Hook et al. 2004 for a description) mounted on the northern 8-m Gemini telescope. We specifically aim to: 1) Determine the (model-independent) range of colours and line indices in extra-galactic bulges and disks along the Hubble sequence; 2) Combine several of the gas-phase and stellar population age and metallicity in-dicators to investigate the star formation and chemical enrichment histories in disk galaxies and place constraints on basic evolutionary scenarios of disk galaxies; Chapter 1. Introduction 18 3) Map structural and dynamical properties of disk galaxies as a function of radius (inner truncation, Freeman Type-II profiles, M / L ratios) that imaging or emission-line spectroscopy alone cannot resolve; 4) Study the role of bars in the formation and evolution of spiral bulges. It is clear that major developments in the theories of the formation and evolution of galaxies have been linked predominantly to major advances in observational technolo-gies. There are a number of ambitious observational programs underway, in particular with the goal of probing to increasingly higher redshifts. However, given the tremen-dous difficulties in relating the unresolved high redshift objects to their present day analogues, it is imperative that we fully explore and understand the current state of these structures. This is a primary concern of the current study. Ultimately, we wish to accurately describe the distribution of stellar population parameters in nearby late-type spiral galaxies as a key to understanding their formation and evolution. As a first, and crucial step, this thesis provides the necessary data and analysis tools for this endeavor. The outline for the thesis is as follows: In the next chapter we describe our colour gradients analysis of the largest multi-band image collection of spiral galaxies which includes the Courteau-Holtzman database used for structural analysis in my Masters thesis (published in MacArtfmr, Courteau, & Holtzman 2003, hereafter Paper I) and the compilation of Bell & de Jong (2000). The following chapter, chapter 3, presents an analysis of the potential effects of dust on line-index measurements as a prerequisite to our analysis in chapter 4, where we present a pilot study of radially resolved spectroscopic absorption-line index measurements of late-type spiral galaxies. By comparison with the latest models of stellar population synthesis, we derive age and abundance profiles well into the galaxy disks for the first time. Finally we summarize our results and conclude with observations for future work in chapter 5. 19 Chapter 2 i Colour Gradients in Spiral Galaxies 2.1 Introduction Colour gradients in galaxies reveal information about the nature of their stellar popu-lations via age and metallicity trends, and the amount and distribution of dust. The technique of using broad-band colours as a probe of the stellar populations and star formation histories (SFH) of galaxies, pioneered by, e.g. Searle, Sargent, & Bagnuolo (1973) and Tinsley & Gunn (1976), is still reflected in the modern studies of Peletier & Balcells (1994), de Jong (1996), and Bell & de Jong (2000; hereafter BdJOO). Early analyses, however, suffered significantly from degeneracies between age and metallic-ity. Worthey (1994) quantified the age-metallicity degeneracy that exists in optical broad-band colours as Alog(Age)/Alog(zT) ~ —3/2. This implies that the composite spectrum of an old stellar population is nearly indistinguishable from that of a younger but more metal-rich population (and vice versa). This degeneracy can be partially bro-ken with infrared photometry (e.g. H or K band) in addition to optical colours (de Jong 1996). Cardiel et al. (2003) have quantified the relative ability of different colour and absorption-line index combinations to constrain physical parameters of composite stellar populations. Their results show that inclusion of an infrared band improves the pre-dictive power of the stellar population diagnostics by ~ 30 x over using optical colours alone. The interpretation of broad-band colour gradients also relies on a careful map-ping of the dust extinction within a galaxy (Witt, Thronson, & Capuano 1992; de Jong 1996). While dust opacity is much reduced at redder wavelengths, its effects may still be non-negligible in the determination of the stellar content of late-type spiral galaxies and must be considered in the final interpretation of colour gradients. 1The analysis presented in this chapter is an expansion of the work published in The Astrophysical Journal Supplement Series as MacArthur et al. 2004, ApJS, 152, 175. Chapter 2. Colour Gradients in Spiral Galaxies 20 In spite of genuine progress in recent years in the study of photometric colour gra-dients in spiral galaxies (see §1.2.1 for a discussion of previous studies), our ability to model and interpret them in terms of formation models is still limited due to the lack of extensive multiband data and the intricacies involved in converting observed quantities to reliable ages and metallicities. In this chapter, we use the largest catalog of deep optical and NIR galaxy colours to date to revisit the comparison of broad-band colour gradients with stellar population models using a range of SFHs and basic assumptions about the dust distribution. Spectral gradients will be investigated in §4. We follow the approach developed by BdJOO, exploring additional parameter ranges and using an ex-tended database by combining the BdJOO data with our own collection of deep, optical and NIR surface brightness profiles (Courteau, Holtzman, & MacArthur 2005, in prep.; hereafter Paper III), for a total sample of 172 galaxies. For the combined database we determine local average ages and metallicities in radial bins for each galaxy and com-pute gradients in age and metallicity as a function of radius. We pay special attention to the effects of fitting out to different physical extents for individual galaxies and the distinction between inner and outer galaxy gradients. In particular, we find that false trends can be inferred if the radial extent of the gradient fits is not taken into account, which in turn leads to erroneous conclusions about the galaxy parameters that drive their SFHs. The outline of this chapter is as follows. In §2.2, we describe the database from which colour gradients are computed. Our colour gradients are presented in Paper III, and further details of the BdJOO sample can be found in their §2. In §2.3 we explore the range of galaxy parameters in our sample and their intrinsic correlations which must be considered when comparing trends in age and metallicity gradients with galaxy parameters. In §2.4, we discuss the stellar population models to be compared to the data and the different star formation prescriptions adopted. Dust models and its potential effects on our results are discussed in §2.5. Optical-NIR colour-colour profiles and their matching population models are presented in §2.6 and trends with galaxy parameters are explored. The technique by which the data are fitted to the SSP models Chapter 2. Colour Gradients in Spiral Galaxies 21 is presented in §2.7. Local and global age and metallicities are presented in §2.8.1 and their gradients are discussed in §2.8.3 and §2.8.4 and contrasted with previous results from BdJOO. We conclude with a discussion of the mechanisms that control stellar evolution in spiral galaxies and compare our results with existing models of galaxy-evolution in §2.9. 2.2 The Data The database from which colour gradients are measured is a combination of the data in BdJOO and the compilation by Courteau, Holtzman, & MacArthur (2005, in prep., Paper III) of 1063 digital images in the B, V, R, and H passbands of 324 nearby late-type spiral galaxies collected at the Lowell Observatory and Kitt Peak National Observatory. For the current analysis we consider only the face-on and moderately tilted galaxies (i < 60°) in the sample. Al l galaxies were selected from the Uppsala General Catalogue of Galaxies (UGC, Nilson 1973) to have: • Hubble type Sb-Sd • Zwicky magnitude < 15.5 • blue Galactic extinction A B = 4x E(B — V) < 0.5 mag (Burstein & Heiles 1984), • blue major axis < 2'. 2. For the computation of homogeneous structural parameters and colour gradients we use the isophotal map determined at i?-band and applied onto all other images (BVH) of a galaxy. The SB profiles are reliable (with SB errors < 0.12 mag arcsec-2) down to ~26 mag arcsec -2 for optical bands and ~ 22 mag arcsec -2 at iJ-band. For the purpose of this colour-based analysis, we further require that the galaxies have measured surface brightnesses in at least two optical and one near-IR (H) band observations out to at least 1.5 iJ-band disk scale lengths and with SB errors of less than 0.12 mag arcsec -2. Chapter 2. Colour Gradients in Spiral Galaxies 22 This leaves us with 51 galaxies (25 type I, 17 type II, 9 transition)2 with extended reliable colour gradients. The SB profiles were corrected for Galactic extinction using the reddening values of Schlegel et al. (1998). We do not attempt to correct our SB profiles for internal extinction, but discuss the possible effects of dust on our results in §2.5. The SB profiles were degraded to the worst seeing FWHM before computing colour gradients. For a full description of the data see Paper III. In order to increase the signal-to-noise (S/N) of our colour profiles, the SB profiles were averaged into radial bins scaled by the NIR-band disk scale length, / I N I R , and we required at least 3 data points per bin. The bin sizes were defined as follows: gridi : 0.0 < r/hmR < 0.5, grid2 : 0.0 < r/ hNm < 0.25, 0.5 < r/hwR < 1.5, 0.25 < r/hmR < 0.5, 1.5 < r/ hNm < 2.5, 0.5 < r/hmR < 0.75, 2.5 < r/hNm < 3.5, 0.75 < r/h-um < 1.0, 3.5 < r/hmR < 4.5, 1.0 < r/hmR < 1.5, 4.5 < < 5.5 1.5 < r/hNm < 2.0, 2.0 < r/hmR < 3.0 The second, finer, grid (grid2) was used to see if the coarser binning of gridi hides any important features in the colour profiles, especially near the center where gradients are steepest. The difference in using the two different binning schemes is small. In particular, this test confirmed that the measured gradients do not depend on the central pixels, which are likely to be more affected by a central concentration of dust or nuclear starburst in spiral galaxies. Hence we subsequently use gridi only (as this matches the binning scheme of BdJOO). In order to look for trends in colour gradients with galaxy parameters, we col-lected, for as many galaxies as possible, their morphological type, absolute magnitude, iJ-band disk central surface brightness (CSB), if-band disk scale length, inclination, 2 For type I SB profiles (Freeman 1970), the inner profile always lies above the SB of the inward extrapolation of the outer disk, whereas type II systems have a portion of their brightness profiles lying below the inward disk extrapolation. We defined a transition case for luminosity profiles that change from type II at optical wavelengths to type I in the infrared. See Paper I for a definition of the different profile types. Chapter 2. Colour Gradients in Spiral Galaxies 23 total if-band magnitude, and rotational velocity VI0i (see below). The structural pa-rameters (disk scale length, h, and CSB, p,0) were computed using our bulge-to-disk decomposition scheme (Paper I). Total galaxy magnitudes encompass the light out to SB levels given above, and include an extrapolation of the SB profile to infinity with an exponential fit to the last quarter of the SB profile3. The extrapolated magnitude increment (0.07 H-mag on average) is added to the isophotal magnitude which is cor-rected for Galactic extinction (Schlegel et al. 1998). Stellar masses are computed using the prescription of Bell & de Jong (2001). We merged our sample with the 121 spiral galaxies from BdJOO which spans a wider range in Hubble type (SO-Irr) and includes a few low surface brightness (LSB) galaxies. For the computation of colour gradients, BdJOO required at least two optical (BVRI) and one if-band observations per galaxy. The data reductions and radial binning (using only gridi) are very similar to our own. For 3 galaxies in common with both data sets, there is excellent agreement in the overall calibration and radial profiles at least for the BVRH bands (Paper III). In order to compare the two data sets directly when looking for trends in age and metallicity with galaxy parameters, we converted our iJ-band magnitudes and surface brightnesses into if-band using the 2MASS H — K colour transformations derived from Jarrett et al. (2003). Since we are interested mainly in qualitative trends in age and metallicity with galaxy parameters, coarse transformations are adequate and we took the mean H — K values for each Hubble type, derived from their Figure 16 and listed here in Table 2.1. The modified H magnitude will be denoted as Hm, to represent the K-band equivalent of the if-band measurement. We do, however, consider the disk scale lengths in the H and K to be directly comparable. Mollenhoff & Heidt (2001) find hj/hx = 1-01 ± 0.19 for a sample of 40 spiral galaxies and the relationship will be even tighter for hn/hx, hence a direct comparison of the H and K band disk scale lengths is justified. 3 This accounts for outer disk truncations that would not be fit properly using the entire profile from our B / D decompositions in Paper I. Chapter 2. Colour Gradients in Spiral Galaxies 24 Hubble Type Sa Sab Sb Sbc Sc Scd Sd mean H — K 0.30 0.26 0.30 0.28 0.28 0.24 0.21 Table 2.1: Mean 2MASS H - K isophotal colours for non-barred spiral galaxies. Adapted from Figure 16 of Jarrett et al. (2003). It has been suggested that age and metallicity gradients may depend upon the potential of the galaxy and thus vary as a function of VTOt (Dalcanton & Bernstein 2002; Garnett 2002). In order to search for such a trend we augmented out catalog with VI0t measurements obtained from H I or Ho; line widths (which we take to be 1^50/2, half the line width at 50% peak flux) collected from the literature. These data enable us to test for a distinct signature in the colour profiles from the more prominent dust lanes seen in galaxies with VTOt > 120 kms" 1 (Dalcanton & Bernstein 2002) and/or if there is a flattening of the metallicity of galaxies above V r o t ~ 150 k m s - 1 (Garnett 2002). The line widths were collected from various sources in the literature; Courteau (1997) [7gals]; Bottinelli et al. (1990) [74 gals]; Theureau et al. (1998) [11 gals]; Haynes et al. (1999) [1 gal]; de Blok & Bosma (2002) [1 gal]; de Blok, McGaugh, & Rubin (2001) [4 gals]. Different measures and measurement techniques are employed for the different samples, so we must derive transformations between difference sources and measures to ensure uniformity between them. For the Bottinelli et al. (1990) and Theureau et al. (1998) samples, the line width at 20% and 50% of the peak flux, W50 and W20, were provided, though not for all galaxies. In cases where only 14^ 20 was available, a conversion to W50 was made from a linear least-squares fit between these two quantities derived from those for galaxies for which both measurements were available (see Fig 2.1). The conversions between the fully corrected line widths are as follows: For the Bottinelli et al. (1990) sample: W50 = 0.98 * W20 - 19.21 [N = 3325]. For the Theureau et al. (1998) sample: W50 = 1.00 * W20 - 17.01 [N = 2055]. Chapter 2. Colour Gradients in Spiral Galaxies 25 200 400 600 200 400 600 W20(Bot90) W20(The98) Figure 2.1: W50 versus W20 for Botinelli et al. (1990) [left] and Theureau et al. (1998) [right] samples. The dotted lines are the linear least-squares fits to N = 3325 points in the Bottinelli et al. (1990) sample and N = 2055 in the Theureau et al. (1998) sample. The bottom panels show the % differences between the two values. Haynes et al. (1999) give a different measure of the line width denoted W21 and defined as the full width between the velocity channels at the 50% level of each horn. We can convert the Haynes et al. (1999) W21s to the Bottinelli et al. (1990) WbOs as (see Fig. 2.2): W50(Bot) = 1.02 * W21(Haynes) + 1.39 [N = 275], and Courteau (1997)'s maximum velocity rotation measures, Vma,x, and Botinelli et al. (1990)'s W50 are matched with: WbO(Bot) = 0.99 * Vmax - 9.31 [N = 122]. For four of the LSBs in the BdJOO sample we obtained VmAX values from de Blok, McGaugh, & Rubin (2001) and de Blok & Bosma (2002). A comparison between the Courteau (1997) V^ax values and the Bottinelli et al. (1990) W50s for galaxies in Chapter 2. Colour Gradients in Spiral Galaxies ZOO 400 W 2 1 „ - - ( H a y n e s 9 9 ) ZOO 400 V m „ ( C o u r t 9 7 ) Figure 2.2: W50 from Bottinelli et al. (1990) sample versus corrected W21 from the Haynes et al. (1991) sample [left] and V m a x from the Courteau (1997) sample [right]. The dotted lines are the linear least-squares fits to N — 275 and N = 122 points in common with the Bottinelli and Haynes and Courteau samples, respectively. Bottom panels show the % differences between the two sample values. common (see Fig. 2.2) reveals little difference, and we assume here that V m a x values for the LSBs also map directly into W50. Multiple measurements of a galaxy from different samples vary typically by less than 10%, more than accurate enough for our purposes. In order to ensure accurate V r o t measurements, we restricted ourselves to galaxies with inclinations greater than 35° (see Courteau 1997). We are left with 98 galaxies in our joint sample that have reliable rotational velocity measurements. 2.3 Galaxy parameters Trends in age and metallicity (local values and gradients) with galaxy parameters, may bear an imprint from intrinsic correlations within the sample itself. In Figure 2.3 we plot the galaxy parameters used in this analysis (h, fio, M H i n , K , and VTOT) against each other. The most notable correlations are the luminosity-rotation speed relation, MH^K versus log(V^ot) (the so-called "Tully-Fisher" relation), and the size-luminosity relation, Chapter 2. Colour Gradients in Spiral Galaxies 27 m O LB •IB -20 -22 -24 -26 M H„.K 0 0.5 1 222120 19 18 17 16 l 0 g ( h H , K ) M o - H o, .K Figure 2.3: Correlations between galaxy parameters: rotational velocity, Vrot (kms - 1 ) , central surface brightness ^o,Hm,K, total magnitude, MHm^K, and disk scale length hH^-The dotted black line is at l o g i n C ^ t = 120 kms - 1 ) . The red points with dotted error bars represent the galaxies from our sample while the blue points with solid error bars correspond to the BdJOO sample. Chapter 2. Colour Gradients in Spiral Galaxies 28 MHm,K versus \og(h). A l l correlations wi th fj,0 are weaker and have large scatter. More detailed scaling relations of spiral galaxies, which match our results, are also presented in Courteau et al. (2005). In Figure 2.4 we plot the 4 basic galaxy parameters as a function of morphological type T 4 . Trends are seen with galaxy parameters for T > 4 such that later-type galax-ies have smaller rotational velocities and total magnitudes, and lower central surface brightness. O n the other hand, there is no clear trend of disk scale length with Hubble type (and hence the trend with total magnitude is largely driven by the trend with C S B , see Courteau et al . 2005). Da ta points are few for T < 4, but deviations from the T > 4 trends seem to occur. A larger spread is (barely) noticeable in VTOt, but is further supported by the larger spread in MHintK for which we have more galaxies (unfortunately, reliable rotational velocities could not be found for the earliest-type galaxies.) The increasing trend in / / n wi th decreasing T seems to level off at T < 4, but the spread is st i l l significant. Finally, earliest-type galaxies have the shortest scale lengths, as expected. We shall return to these correlations in the analysis of trends in age and metallicity gradients (§2.8). 2.4 Stellar Population Models In order to translate colour information into constraints on the underlying ages and abundances of the stellar populations, the colour gradients must be compared with stellar population synthesis models. In their most basic form, commonly referred to as simple stellar populations (SSPs), these models provide evolutionary information for a coeval population of stars born with a given composition and ini t ia l mass function ( I M F ) . The closest physical analog to such SSPs are globular cluster systems from which the SSP models are calibrated (e.g. Schiavon et al . 2002a). Several such SSP models have been produced by a number of independent groups and are in a constant state of modification as improvements to many of the input parameters (e.g. stellar 4 The numerical classification, T, corresponds to the following Hubble types: E0 (T=-5), SO (T=-l ) , Sa (T=3), Sb (T=5), Sc (T=7), Sd (T=9), and Sm ( T = l l ) Irr (T=13), with the mid-types filling in the remaining numbers. Chapter 2. Colour Gradients in Spiral Galaxies 29 Figure 2.4: Correlation between galaxy parameters and morphological type. The dotted line is at logio(V^o t = 120 k m s - 1 ) . Chapter 2. Colour Gradients in Spiral Galaxies 30 libraries, model atmospheres, convection, mass loss, mixing) come to light (see §3.3 for further discussion and application of these models). There are discrepancies among the different models that, depending on the application, may result in significantly different interpretations of observations of unresolved stellar populations. In order to determine if these model discrepancies could affect our analysis, we have compared two independent sets of SSPs; the 2003 implementation of the Bruzual & Chariot (2003) models (hereafter referred to as GALAXEV) , and the Project d'Etude des Galaxies par Synthese Evolutive (PEGASE 2.0) models of Fioc & Rocca-Volmerange (1997). The stellar populations of interest here are those of spiral galaxies spanning the full Hubble sequence (SO-Irr). The approximation of a single stellar population clearly does not apply to evolved, complex stellar populations of spiral galaxies. However, if one assumes that galaxies are composed of a superposition of SSPs, born at different epochs, rates, and metallicities, one can use the SSPs to develop model grids that mimic the range of plausible galactic stellar populations and these can be compared with galaxy colour profiles. The stellar population model grids are created by taking the single burst SSPs, with constant stellar IMF and fixed metallicity, and convolving them with a given SFH. A few simplifying assumptions are inherent in this formulation. We are assuming that, the IMF does not change as a function of time or galactic environment. The validity of this assumption may be questioned from a theoretical perspective as the IMF is expected to vary systematically with star formation environments (Chabrier 2003, and references therein). However, significant observational evidence of star formation in small molecular clouds, rich and dense massive star-clusters, as well as ancient metal-poor stellar populations, reveals a remarkable uniformity of the IMF (Kroupa 2002). The use of fixed-metallicity SSP implies no chemical evolution. This is clearly an unrealistic assumption with regards to galaxy evolution (as it also ignores feedback effects from stellar winds and supernovae), but it still allows us a reasonable comparison of relative metallicities and ages within and among galaxies (Abraham et al. 1999; BdJOO; Gavazzi et al. 2002). The specific choice of IMF, parameterized as £(logm) oc m - x , can also potentially Chapter 2. Colour Gradients in Spiral Galaxies 31 affect our results. To gauge the importance of this effect on our colour-based analysis we compared model grids obtained with some of the most widely used IMF characteri-zations; (a) the single-slope Salpeter (1955) with x = 1.35, 0.1 < m/MQ < 100, (2.2) (b) the Kroupa (2002) IMF (PEGASE models only) with: x0 = 0.3, 0.1 < m/MQ < 0.5 x1 = 1.3, 0.5 < m / M Q < 120, (2.3) and (c) the Chabrier (2003) single-star Galactic disk IMF (GALAXEV models only) which is parameterized as: e ( l o g H 4 e X P [ - " ° * " ^ m ' ' ' ] ' f ° r 0 J - m / M 8 S 1 : (2.4) [ m - 1 3 , for 1 < m / M 0 < 100 where mc and a2 = ((logm — (logra))2) denote the mean mass and the variance in logm, respectively. The Chabrier (2003) IMF is preferred over the Kroupa (2002) and Salpeter (1955) IMFs because of its better agreement with number counts of brown dwarfs in the galactic disk and its theoretical motivation. The integrated spectrum, F\(t), of a stellar population with SFR \&(r) is computed as the convolution of a single burst stellar population of given metallicity, f\{t), with the given SFR: Fx(t)= f\(t-t')fx{t')dt'. (2.5) Jo The left panel in Figure 2.5 compares the model tracks from the G A L A X E V models with the Salpeter (1955) (blue grid, eq. 2.2) and Chabrier (2003) (red grid, eq. 2.4) IMFs, and the right panel compares those from the PEGASE models for Salpeter (1955) (red grid) and Kroupa (2002) (green grid: eq. 2.3) IMFs. Iso-metallicity tracks are connected by solid lines, while iso-age tracks are connected by dashed lines. Chapter 2. Colour Gradients in Spiral Galaxies 32 The difference between the various IMFs on the model grids are small, thus the specific choice of IMF (within current observational constraints) will not affect our results. We opt to use the Salpeter (1955) IMF for the remainder of this analysis (primarily to facilitate comparison with other studies). Also over-plotted in the right B - R B - R Figure 2.5: Comparison of the stellar population model tracks for different initial mass functions; the Salpeter (1955) IMF (blue grid, eq. 2.2), and the Chabrier (2003) IMF (red grid, eq. 2.4) for the G A L A X E V models [left panel] and of the Salpeter (red grid) versus the Kroupa (2002) (eq. 2.3) IMFs for the PEGASE models [right panel]. The G A L A X E V models are also plotted on the right panel (green grid) for comparison between two different SSP models. Iso-metallicity tracks are connected by solid lines, while iso-age tracks are connected by dashed lines. panel of Figure 2.5 are the G A L A X E V model with Salpeter IMF tracks (blue grid) for direct comparison with the PEGASE model tracks for the same IMF (although note that the upper mass cut-off is 100 M 0 in the G A L A X E V models and 125 M Q in the PEGASE models). While differences in the grids between the different models are noticeable (most likely due to different treatments of the thermally pulsing asymptotic giant branch phase, see e.g. Maraston 2003), they are only significant in the bluest regions where the SB errors of our observations are also large (i.e. the outer regions of our spiral galaxies where sky subtraction errors become large, are, for the most part, bluer). Slight differences in the age and metallicity gradients would be inferred using Chapter 2. Colour Gradients in Spiral Galaxies 33 the different models, but we are mainly interested in relative quantities, and thus are not concerned with the small absolute differences between models. For the rest of this analysis we refer only to the GAL A X E V models. We have explored two different SFR regimes, both parameterized by a star formation timescale r. The first is the simple exponential SFR, \&exp(£), = — e~t/r, (2.6) r where Cexp - 1 _ e _ A / r - (2.7) The second is the so-called "Sandage" SFR, ^san(^) first parameterized by Sandage (1986), *san(*) = c ^ e - t 2 ' ^ (2.8) r 2 and CSan - 1 _ e _ 1 / 2 ( j 4 / r ) 2 ( 2 - 9 ) where c e x p and Cs a n are normalization constants (to 1M Q ) . Figure 2.6 shows the time evolution of the exponential [top panel] and Sandage SFRs [bottom panel]. The ex-ponential SFR falls off at a given rate for all positive values of r. However, in order to cover the full colour-colour space spanned by the observations, negative values of r were also included which allow for increasing SFRs. The Sandage SFR is characterized by a delayed rise in the SFR, followed by an exponential decline, the rate of which is determined by the value of r. For values of r > A, where A is the age of the galaxy (fixed to 13 Gyr in this work), the SFR is still rising at the present time. The average age of a stellar population of a given r is computed as: = S r - (2-10) Jo * ( « ) d i For the exponential SFR, this equates to, ( A U = A . T i z * ^ m (2.n) Chapter 2. Colour Gradients in Spiral Galaxies 34 Age (Gyr) Figure 2.6: Time evolution for the exponential (eq. 2.6) [upper panel] and Sandage (eq. 2.8) [lower panel] star formation histories (solid curves). The dotted curve is a Sandage-style burst of star formation in which 10% of the total mass of stars are formed. See Figure 2.9 for the effect of such a burst on the population model grids. Chapter 2. Colour Gradients in Spiral Galaxies 35 and for the Sandage SFR, (^ )san = A -1 - e 2 ^ (2.12) Figure 2.7 shows the resulting model grids for the exponential (blue), and Sandage (red) SFRs overlaid on each other. Lines of constant average age (A) are the dashed, roughly vertical lines (labeled in Gyr), while lines of constant metallicity are solid and run closer to the horizontal (labeled in Z, the mass fraction in elements heavier than helium, where Z Q = 0.02). The overall shape of the two model grids are quite similar, i 1.5 —I 1 . 1 , 1 1 1 1 r-Sandage SFH < ^ > (Gy r) Exponential SFH " L£j? 0.000\ YMOl GALAXEV B - R Figure 2.7: Comparison of the stellar population model tracks for exponential (blue grids, eq. 2.6) and Sandage (red grids, eq. 2.8) star formation histories. the main difference being that the Sandage SFH has a younger average age for a given B — R and the metallicity is slightly higher for a given R — H. The Sandage grids also cover a wider range in ages (which would increase the magnitude of our measured age gradients). The reported radial age gradients inferred from colour gradients (e.g. BdJOO) seem to be at odds with the lack of age gradient observed in open clusters of the Galaxy (Freeman & Bland-Hawthorn 2002). Would it be possible that the apparent radial age Chapter 2. Colour Gradients in Spiral Galaxies 36 gradients observed in spiral galaxies are simply due to small amounts of young frostings of star formation on top of an underlying old population? In order to test for such effects in our age determinations, we produced new model grids for the same underlying SFH but with an additional burst of star formation at different times and fractional masses. The burst was added as a Sandage (1986) type profile (eq. 2.8) with r = 0.15. The average age for the models with a burst superimposed on the underlying SFR is now computed as, where bfrac is the mass fraction of the burst and tburst is the time since the burst, i.e. the burst of star formation occurs tburst Gyr before the current age of the galaxy, A. Figures 2.8 and 2.9 demonstrate the effect of such a burst of star formation superimposed on an underlying exponential SFR with ages (hurst) of 1 and 4 Gyr respectively. The left panel in both figures shows the effect if 10% of the galaxy's total mass was involved in the burst, and the right panel shows a 50% (by mass) burst. The models certainly do not cover the same extent in the colour-colour space for recent bursts. With only 10% of the total galaxy mass in the burst (Fig. 2.8; left panel) the grid does not extend as far to the red as those without bursts. The effect is strongest (~ 0.3 mag) in the optical colours. Also, the burst grid does not extend as far into the blue optical colours (< 0.1 mag) for the most metal-rich regions. This is because we are essentially adding more older stars to the SFHs that are rising at 13 Gyr (r < 0 for ^ e x p , or r > 13 Gyr for ^>San] This effect is stronger at higher metallicity because the evolution of B — R colour with time steepens for higher metallicity stellar populations at ages ~ 1-3 Gyr in the G A L A X E V models.) The lines of constant age also steepen. "Frosting" also shifts the R — H colours to redder values (~0.1mag) at a given metallicity. This, in fact, prevents the burst grid from reaching red-ward enough as required by the observational data for the inner-most parts of the galaxies (see Figs. 2.11-2.15). Either the bursts are accompanied by significant amounts of dust, or 1 Gyr old bursts can be ruled out 5 T h i s timescale is justified on the basis that starbursts are consistent wi th a constant S F lasting 10-100 M y r (Meurer 2000). (A) = A (1 — bfrac) Jo t^(t) ^ + bfrac Itburst^ ~ tburst)^ bur st(t ~ tburst) dt (1 - bfrac) Jo"4 *(<) dt + bfrac f £ u r s t ^burst(t ~ W s i ) dt (2.13) Chapter 2. Colour Gradients in Spiral Galaxies 37 Figure 2.8: Comparison of the stellar population model tracks for a pure exponential SFH (blue grid, eq. 2.6) and one with a 1 Gyr old Sandage (red grid, eq. 2.8) burst contributing 10% of the total mass [left plot] and 50% of the mass [right plot]. Figure 2.9: Comparison of the stellar population model tracks for a pure exponential SFH (blue grid, eq. 2.6) and one with a 4 Gyr old Sandage (red grid, eq. 2.8) burst contributing 10% of the total mass (left plot, and see Figure 2.6) and 50% of the mass [right plot]. Chapter 2. Colour Gradients in Spiral Galaxies 38 for the central parts of late-type galaxies. The 50% by mass burst in Figure 2.8, right panel, shows that the grid covers only a narrow range in colour - colour space that does not extend far enough into the red or the blue to agree with the observations. For older bursts of star formation, the model grids have colours fully consistent with the no-burst models. This is shown in Figure 2.9 for a pure exponential SFH (blue) overlaid with an exponential plus t^st = 4 Gyr Sandage-like burst (red). The grids are almost identical with respect to their shape and location (with slight narrowing of the 50% mass burst grid). The only difference is the average ages (eqs. 2.10 & 2.13), being younger for a given colour of the burst grid. Thus, adding a burst of star formation that is older than ~ 1 Gyr has a similar effect as changing the overall form of the SFH; the grid shape and colour coverage are the same, but the average age can change significantly from one scenario to another. Not only do the age determinations change, but the age gradients can also be different for a given SFH (compare Sandage and exponential SFH model grids in Fig. 2.7). Since the actual SFHs of real galaxies are unknown, the absolute values of the gradients measured with this colour-based technique cannot be trusted. We can, however, trust that a gradient exists in a galaxy, just not its magnitude. We can also compare relative gradients among different galaxies if we assume self-similar SFHs and no significant bursts of star formation within the last 1-2 Gyr. A reasonable estimate of the validity of this assumption can be obtained from the study of Kauffmann et al. (2003a) who used the R5A absorption line index and the 4000 A break age indicators along with the Bruzual & Chariot SSPs to constrain the SFHs, dust attenuation, and stellar masses of over 100,000 galaxies from the Sloan Digital Sky Survey (SDSS). For each galaxy they estimated F^,.^, the fraction of the total stellar mass formed in bursts in the past 1-2 Gyr. Their Figure 5 shows the fraction of galaxies with F 6 u r s t > 0 at the 50% and 97.5% confidence levels in stellar mass bins ranging from 8-121ogi 0(M o). From this and estimates of the galaxy stellar masses (Bell & de Jong 2001), we infer the fraction of galaxies in our sample that have undergone bursts of SF in the past 1-2 Gyr is ~ 1% (~ 12%) at the 97.5% (50%) confidence level. The fraction of galaxies that have Fburst 10% will be even smaller Chapter 2. Colour Gradients in Spiral Galaxies 39 and we conclude that our assumption of no recent burst is valid. 2.5 Dust effects While the presence of dust in late-type spiral galaxies is well-established, its distribu-tion and effective optical depth remain poorly constrained. Extinction by dust could potentially mimic a colour gradient and its effects must be considered in determin-ing the stellar content of galaxies from colours. We attempt to quantify the effects of dust extinction on colours with simple dust geometry extinction models, adopting the Milky Way (MW) extinction curve and albedo values from Gordon, Calzetti, & Witt (1997). The extinction vector for the simplistic foreground screen model with Av = 0.3 is shown (as the arrow) in the upper left corner in Figures 2.11-2.15. The vector is in the same general direction as the observed gradients of the galaxies, but given the unrealistic nature of this model, we do not consider it further (but include it for vi-sual comparison). Also shown in the top left corner in Figures 2.11-2.15 is the more realistic "face-on triplex" dust model of Disney, Davies, & Phillipps (1989) and Evans (1994). This model assumes that the stars and dust have exponential distributions in both the radial and vertical directions with radial scale lengths / i* and hd, and vertical scale lengths z* and zd, respectively. Although scattering is not taken into account, Byun, Freeman, & Kylafis (1994) and de Jong (1996) have shown that its effects are negligible in face-on galaxies (as photons are just as likely to be scattered into the line of sight as they are out of it.) We therefore use the dust absorption curve to compute the dust optical depth; rx(0) = — TV(0)(1-O a) (2.14) Ty where T\ and a\ are the dust optical depth and albedo respectively. Studies of multi-wavelength data that solve simultaneously for the intrinsic colour of stellar populations and reddening and extinction effects from dust find central optical depths, ry(O), in the range 0.3-2.5 for types Sab-Sc (Peletier et al. 1995; Kuchinski et al. 1998; Xilouris et al. 1999). The HST and ground-based colour gradient study of early-type bulges (< Sb) by Peletier et al. (1999) reveals dust extinctions of Ay = 0.6-Chapter 2. Colour Gradients in Spiral Galaxies 40 1.0 mag (or A H = 0.1-0.2 mag) in galaxy centers but negligible extinction beyond one effective radius. Multi-band SB profile modeling of massive edge-on disks suggests that the dust is confined to a thin extended plane, with Zd/z^O.7 and hd/h* ~ 1.4 (Xilouris et al. 1999; Matthews & Wood 2001). Masters, Giovanelli, & Haynes (2003) studied internal extinction in the near-IR for a sample of 15,244 2MASS galaxies by examining the inclination dependence of various photometric parameters, and concluded that galaxies with Mx - 51og(h) > -20, -20.7, and -20.9 in the J, H, and ifs-bands respectively are spared any extinction, but that disk opacity increases monotonically with disk luminosity. Application of the triplex model for these galaxies favors small dust-to-star scale height ratios (zd/z* ~ 0.5, in good agreement with Xilouris et al. 1999) and face-on central opacities of r°(0) = 0.7 and 0.3 at H and Ks respectively. We adopt the same vertical and radial scale length ratios for the dust and stars; using zd/z* = 0.7 and hd/h* = 1.4 instead stretches the gradients by only 0.018 in B — R, 0.019 in R — H, and 0.021 in R — K, and the central V-band optical depth to Ty(0) = 1 (or pole-to-pole F-band optical depth of 2). The higher end of the measured range of ry(0) = 2 is shown in the upper left corner of Figures 2.12 & 2.13 (right panels). The triplex models lie parallel to, and thus could contribute significantly to, the observed gradients (note there is no overall calibration in the dust models so they can slide to any region on the plot). For the gradients to be entirely due to dust requires extremely high central optical depths (ry(0) ~ 5) to reproduce the data. Also, the triplex models alone cannot reproduce the significant colour gradients observed from the half-light radius (denoted by open circles) outward. Thus, while dust is likely a contributor to colour gradients, age and metallicity effects must also be invoked. Dalcanton & Bernstein (2002) have analyzed optical and infrared colour maps of 47 extremely late-type (bulge-less) edge-on spirals spanning a wide mass range (40 < V c < 250 kms - 1 ) . They find that higher mass (rotation velocity) galaxies have more prominent dust lanes and have redder colours than the stellar population grids (whereas the less massive galaxies have colours consistent with the population grids). Analysis Chapter 2. Colour Gradients in Spiral Galaxies 41 of their R — Ks colour maps, and a comparison of colours between their edge-on sample and the face-on sample of de Jong (1996) suggests that dust plays little role in all but the most massive galaxies in their sample (V c > 120 kms - 1 ) . Previous analyses have ruled out dust effects on the colour gradients in galaxies based on the fact that they see no correlation between the gradients and inclination. Dust effects on gradient profiles as a function of bulge-to-total ratio (B/T) and inclination were modeled by Byun, Freeman, & Kylafis (1994). Their Figure 6 shows B — I versus radius for model galaxies with different B / T ratios and rv(0) ranging from 0.0-10.0 for inclinations 0-85°. In the 0-50° range the differences in the profiles are very small and would not be detected in the observations. Our measured radial gradients also show no inclination dependence, but we do not consider this sufficient grounds to rule out significant dust extinction. Based on optical-IR imaging of SO-Sbc galaxies, Peletier & Balcells (1996) found that "dust-free" colours of galaxy disks are not significantly different from their bulges. They derived bulge colours from minor-axis wedges in images of early-type edge-on spirals. The wedges are presumed dust-free above the disk plane. They conclude that (dust-free) bulge and disk colours are very similar with A(B — R) = 0.045 ± 0.097 and A(R - K) = 0.078 ± 0.165. We find a significantly different result; as can be seen in Figure 2.10 (compare with Fig. 2 in Peletier & Balcells 1996), our bulges are much redder than their disks with A(B - R) = 0.29 ± 0.17 as A(R - K) = 0.30 ± 0.17. This could be due to genuine dust extinction in our galaxies or there is a fundamental difference between our respective samples and/or analysis methods. Their disk SB profiles were measured along 10°-wide wedge apertures centered 15° away from the disk major axis, to avoid the prominent dust lanes near the major axis of their inclined galaxies. Naturally, this technique is sensitive to a vertical disk colour gradient. Dalcanton & Bernstein (2002) have suggested that all thin disks are embedded in a red stellar envelope. Whether this envelope is redder or bluer than the thin disk depends on the rotational velocity of the galaxy: redder envelopes for VTOt < 120 k m s - 1 and bluer otherwise, but the colour of the envelope is similar for all disk galaxies. The Chapter 2. Colour Gradients in Spiral Galaxies 42 bulge bulge Figure 2.10: Disk colours (average of 1.5-2.5 disk scale length radial bin) as a function of bulge colours (average of 0.0-0.5 disk scale length radial bin) for optical B - R [left panel] and optical-NIR R— K [right panel] colours. The solid horizontal lines represent a one-to-one mapping (for reference only). Hm is the "modified" i7-band magnitude for the Courteau et al. sample, converted to if-band with 2MASS H - K colours (see text for details) for direct comparison with the if-band data of the BdJOO sample. redder thin disks of the galaxies with VI0t > 120 k m s - 1 are attributed to strong dust lanes observed in their B — R colour maps, which disappear in the VTOt < 120 k m s - 1 galaxies. Clearly, these red stellar envelopes, and hence the presence of vertical colour gradients in disk galaxies renders the interpretation of Peletier & Balcells (1996) "disk" colours difficult. The 10° offset from the major axis may result in a measurement of the "red envelope" stars which are intrinsically older than the thin disk stars (though not necessarily redder since thin disk stars could be reddened by a central dust lane). While optical-IR colour imaging alone is not sufficient to break the degeneracy be-tween dust and stellar population effects on the colour gradients, the tentative consensus to date is that dust is generally not a significant contributor to galaxy colours in low-mass/low-luminosity spiral galaxies, but is likely important in more massive/brighter galaxies. It must be borne in mind that future studies of high-resolution IR and FIR imaging and absorption-line spectroscopy may radically alter this view. Chapter 2. Colour Gradients in Spiral Galaxies 43 2.6 Colour-Colour Profiles Optical-NIR colour-colour profiles are shown for the Courteau et al. sample in Fig-ures 2.11-2.12 and 2.14-2.15 (see Figs. 1 -5 in BdJOO for similar plots with their sample and our Fig. 2.13 which is only a subset of the BdJOO sample with available VI0t values). Typical observational errors are shown as crosses in the lower right corners of Figures 2.11-2.15. From left to right the crosses represent calibration errors, and average sky subtraction errors for the innermost and outermost bins. The galaxy B - R B - R Figure 2.11: Near-IR-optical colour-colour plots for the Courteau et al. sample separated by morphological type. The galaxy centers are indicated by solid symbols: circles for Type-I galaxies, triangles for Type-II, and asterisks for Transition galaxies. Open circles denote the half-light radius of the disk. The line types for the galaxy colour profiles represent bar strength: solid for bar-less (A), dashed for mild bars (AB), and dot-dashed for strong bars (B). Average errors due to the global calibration, as well as sky uncertainties for the innermost and outermost points are shown as crosses. Over-plotted are the Bruzual & Chariot (2003; GALAXEV) stellar population models for an exponential SFH at different metallicities. Lines of constant average age (A) (see eq. 2.10) are the dashed, roughly vertical lines, while lines of constant metallicity are solid and nearly horizontal. In the upper left corner of each panel, triplex and foreground screen dust models are plotted. For the triplex model, the Milky Way (MW) dust extinction curve and albedo from Gordon, Calzetti, & Witt (1997) was adopted. The galaxy center and disk half-light radius are denoted by solid and open circles respectively. Chapter 2. Colour Gradients in Spiral Galaxies 44 2.5 1.5 1 1 1 ' 1 I 1 1 1 1 1 ' ' ' „ *.«>> -1 <A> (Gyr) MW 12.9 h * / h - " 1 1 0 ^ - ^ 0.05 fc/* - 1 9 calib Inner outer 66 S V M < 120 GALAXEV _ 2.5 h 1.5 ' i ' Aj.0) - 2 M W / / h y n . - 1 ja z./z. » 1 —I : ; ; <A> (Gyr) , £ . 5 , 1 . . ; JrJi^Z o.o5' • 7.6 • A, - 0.3 6.4^^1 8-5 ^ - < ^ V ^ i/ * i ^ 0 ^ 2 5.o^<^; &*Ef*ffi\ \t -'0.008 fit S<-^ —' * ' J 7 T Ji 0.004 " i / i Jr r r / /' / >0 f f t • i A " ' - £ / 0.0004 frOOOl + . Type I - Type II ~~y * T r a n s calib inner outer '_ • 120 l , , , , l S < 280 i i i i GALAXEV^ i 0.5 1 1.5 B - R 0.5 1 1.5 B - R Figure 2.12: Near-IR-optical colour-colour plots for the Courteau et al. sample separated by rotational velocity, VTOt ( k m s - 1 ) . Symbols and line-types are as in Figure 2.11. Figure 2.13: Near-IR-optical colour-colour plots separated by rotational velocity, VT ( k m s - 1 ) , for the BdJOO sample. Chapter 2. Colour Gradients in Spiral Galaxies 45 3 p 1 . 1 1 1 1 1 1 1 1 1 1 1 1 1 3 p 1 1 1 1 1 1 1 1 1 1 1 1 1 r U i i i i I i , , , I i i i i I U i i i i I i i i i I i i i i I 0.5 1 1.5 Z 0.5 1 1.5 2 B - B B - R Figure 2.14: Near-IR-optical colour-colour plots for the Courteau et al. sample separated by if-band magnitude, MH. Symbols and line-types are as in Figure 2.11. Figure 2.15: Near-IR-optical colour-colour plots for the Courteau et al. sample separated by if-band central surface brightness, PO,H- Symbols and line-types are as in Figure 2.11. Chapter 2. Colour Gradients in Spiral Galaxies 46 centers are indicated by solid symbols: circles for type-I galaxies, triangles for type-II, and asterisks for transition galaxies (as defined in Paper I). Open circles denote the half-light radius of the disk. The line types for the galaxy colour profiles represent bar strength: solid for bar-less (A), dashed for mild bars (AB), and dot-dashed for strong bars (B). It is conceivable that the type-II "dip" could be due to dust extinction and/or stellar population effects, possibly linked to the presence of a bar, and or inner-disk truncation (see Paper I). However, we do not see any distinction between the different SB profile types in the colour gradients. This argues against dust as a major factor for the type-II phenomenon as dustier galaxies would be redder and have larger gradients. No conclusion can be drawn about the effects of a bar due to the small size of our sample and the fact that non-barred galaxies may have once harboured a bar that dissolved after mixing the stellar population. Over-plotted on Figures 2.11-2.15 are our convolved Bruzual & Chariot (GALAXEV) stellar population models with an exponential SFH at different metallicities. Lines of constant average age (A) (see eq. 2.10) are the dashed, roughly vertical lines, while lines of constant metallicity are solid and nearly horizontal. In the upper left corner, triplex and foreground screen dust models are plotted with the galaxy center and disk half-light radius denoted by solid and open circles respectively. All of the galaxies show significant colour gradients that are consistent with gradients in both age and metallicity with the central parts having older mean ages and being more metal rich. Figures 2.12 & 2.13 for the Courteau et al. and BdJOO samples reveal a largely consistent picture of the radially-resolved colours of spiral galaxies, though the Courteau et al. galaxies extend to higher metallicity (0.1-0.2 mag in R — H). The origin of this discrepancy is not clear, but we note that the optical versus J — K colours of galaxies are not well fit by single metallicity models (e.g., Fig. 15 of Bell et al. 2003). It may well be that, due to TP-AGB prescriptions, systematic errors of 0.2 mag between H and K predicted colours are inevitable. Thermal emission from the telescope at if-band may also affect sky flats of the type used by BdJOO. Bearing in mind these potential sources of system-atic uncertainty, we feel that the degree of similarity between the Courteau et al. and Chapter 2. Colour Gradients in Spiral Galaxies 47 BdJOO galaxy samples is satisfying. The small number of earlier-type galaxies (Sab - Sbc) in our sample makes it difficult to infer any trends in age and metallicity with galaxy type, but Figure 2.31 shows hints that later-types are slightly less metal-rich than the earlier types. In Figures 2.12 & 2.13 we divide the Courteau et al. and BdJOO samples, respectively, at Vrot > 120 kms - 1 , the threshold above which all edge-on galaxies show prominent central dust lanes (Dalcanton & Bernstein 2002). Due to inclination restrictions, we could only retrieve reliable H I line-widths for 28 of our 51 galaxies and 70 out of the 121 BdJOO sample. A trend with VI0t is detected in Figures 2.12 & 2.13 with faster rotators being more metal rich and having older mean ages; also consistent with them having a higher dust content. However, according to the triplex models, the true signature of dust is not simply a redder colour, but also a stretched colour profile (compare triplex profiles in Figures 2.12 & 2.13 for V-band optical depths of TV(0) = 1 [left panels], and Ty(0) = 2, [right panels]). For our sample (Fig. 2.12), the length of the gradients appears to be about the same in both velocity bins. For the BdJOO sample (Fig. 2.13) the gradients could be more extended in the larger VTOt bin, but the stretching is not predominantly along the R — K axis, as in the triplex dust models. Clearly we cannot attribute with absolute certainty the redder colours of the faster rotators to dust effects. Ages and metallicities are determined by fitting the SPS models to the radial galaxy colours using a maximum-likelihood approach similar to that of BdJOO (see their §3 for further details). We compute a finely spaced grid in r using equations 2.5 & 2.6, and interpolate (linearly) between the SPS metallicities. Treating the colours and errors of each annulus separately, we compute an average age and metallicity per annulus by minimizing the x2 statistic: 2.7 Model Fitting (2.15) Chapter 2. Colour Gradients in Spiral Galaxies 48 where N is the number of passbands (at least 3; 2 optical plus 1 IR), 5ptot,i is the total error in passband i, and fxc is the best normalization between the model and data computed as a weighted average: Errors for the individual age and metallicity measurements as well as their gradients are estimated using a Monte Carlo approach. One hundred realizations of the model fits are obtained for each radial bin using errors drawn from a normal distribution of the observational errors (which include calibration, sky subtraction, and flat fielding errors added in quadrature). For each realization, gradients are computed as weighted linear fits to the parameter (age and metallicity) determinations as a function of the radial scale length. The weights in the fits are taken as the A% 2 = 1 interval for each annulus-model fit. The error for the measured gradients and individual ages and metallicities are taken as half the interval containing 68% of the 100 Monte Carlo realizations (i.e. the l a confidence interval). Figure 2.16 shows two examples of our fits in (A) -\ogw(Z/ZQ) space. A few galaxies in both samples fall outside the model grids and thus cannot be fit reliably. These were removed from the sample for the model fits, but are worthy of some discussion. As can be seen in Figure 2.11, five of our galaxies lie significantly redward of the model grids in their central R — H colour. The dust vector indicates that such red colours could result from a central dust concentration. Another explanation could be an extremely (unrealistically) metal-rich central stellar population. Three of our galaxies lie considerably redward of the model grids in their B — R colours. Given the current WMAP measurement of the age of the universe (13.7 Gyr; Spergel et al. 2003), it would be unrealistic for these galaxies to have extremely old central regions. Even if the model age calibrations were unreliable, the oldest age of the models of 20 Gyr still do not extend as far red in B — R as these outliers. Dust likely contributes to the red colours, but they could also result from erroneous seeing measurements or photometric calibration errors. In the BdJOO sample 3 very late-type (Sdm/Irr) galaxies (see left N / J o b s , i - A ' m o d e l , i ( ( ^ ) . - Z ) ] ' tot,: N l (2.16) Chapter 2. Colour Gradients in Spiral Galaxies 49 -2 -1.6 -1 -0.5 0 -2 -1.5 -1 -0.5 0 l o g 1 0 ( Z / Z Q ) l o g l o ( Z / Z 0 ) Figure 2.16: Examples of Monte Carlo realizations for age and metallicity fits. Gray curves represent the best fits using errors drawn from a normal distribution of the observational error, with the thick black line being fit with the nominal measurement error. The blue and purple dotted lines are the Ax 2 = 1 interval for each annulus-model fit, with the thick dashed lines being the corresponding interval for the nominal measurement error. panel of Fig. 2.13) and the extended tails of another 3 galaxies lie blue-ward of the model grids in their R — K colours. These low surface brightness galaxies may suffer from calibration and sky uncertainties couple with the weakness of the model predictions at such low metallicities. Finally, the BdJOO sample has 3 early-type galaxies whose entire profiles lie red-ward of the grids in their B — R colours. Again, extremely old ages are improbable, but significant dust absorption could contribute to the red colours of these three galaxies. Thus, unlike some of the studies cited in §2.5, these observations remind us of our fragile understanding of dust effects and the challenge we face in trying to separate them from stellar population effects. Chapter 2. Colour Gradients in Spiral Galaxies 50 2.8 Results 2.8.1 Local and Global Trends in Age and Metallicity Figure 2.17 shows the local average age [left] and metallicity [right] as a function of local surface brightness. There is a clear trend in the local age and metallicity as a function of local SB in the sense that regions of higher SB are older and more metal rich, but with large scatter, as also found in BdJOO. This correlation suggests that the local potential plays an important role in the SFH. 23 22 21 20 19 18 17 16 15 14 23 22 21 20 19 18 17 16 15 14 / L4 H K (mag arc-sec - 2) p^ K (mag arcsec - 2) Figure 2.17: Average age, (A) [left], and metallicity, \og10(Z/ZQ) [right] as a function of local Hm and if-band surface brightness, where Hm is the "modified" if-band SB for the Courteau et al. sample, converted to if-band with 2MASS H — K colours (see text for details) for direct comparison with the if-band data of the BdJOO sample. The Hm data is distinguished with dotted error bar lines. The different symbols represent total Hm, ff-band magnitude bins for each galaxy. The dotted lines in the metallicity plot [right] denote the limits of the population model grids. Points at or near these limits for metallicity, and those at 12.9 Gyr for the average age, should be interpreted with caution. Figure 2.18 shows an effective average age, (A)eg, for a galaxy, taken as the measured mean age in the 0.5 -1 .5 h averaged radial bin, as a function of central surface brightness [left plot] and total galaxy magnitude [right plot]. There is a clear trend in (A)eg as a function of both HQ^H^K and MHINTK with more luminous and higher CSB galaxies Chapter 2. Colour Gradients in Spiral Galaxies 51 22 21 20 19 18 17 16 -18 -19 -20 -21 -22 -23 -24 -25 -26 -21 M 0 . H . .K ( m a g arcsec"2) M H__ K (mag) Figure 2.18: "Effective" average age (average of 0.5-1.5/i bin), (A)ee as a function of central surface brightness [left] and total Hm, K-band galaxy magnitude [right] for all 158 galaxies. See caption of Figure 2.17 for the definition of Hm. having older {A)eg (as expected from Fig. 2.17). However, in both cases, the roughly linear increasing trend at fio,Hm,K ~ 18.5 mag arcsec-2 and Mnm,K ~ ~ 22.5 mag seem to flatten for brighter values. For MHm^K < — 22.5 there is large scatter and a much weaker (if any) trend of increasing age with total magnitude (but still containing the oldest galaxies). In Figure 2.19 we plot (A)eg as a function of hH,K [left] and VTOt [right]. As expected from the correlation of Mam,K with Vrot (Fig. 2.3), a trend of increasing (A)eg with Kot is detected with a similar turnover in the slope of the correlation at higher rotational velocities (VTOt^t 120 kms - 1 , the location of the vertical dotted line). The left panel of Figure 2.19 may show a weak trend of increasing age with h, however, the scatter is very large. Note also that the spread in age is larger for smaller h galaxies, and reduces with increasing h. Figures 2.20 & 2.21 similar trends for the effective metallicity, log10(Zeg/ZQ) versus l^o,Hm,K-, MHjnyK, and V^t, but no trends are seen with h. Effective metallicity increases with total magnitude up to MHmtK ~ —22, but likely saturates for brighter galaxies at roughly solar metallicity (with large scatter). Similarly, in the right panel of Figure 2.21, Chapter 2. Colour Gradients in Spiral Galaxies 52 log (h (kpc)) A < V V r o t ( k m s-i) Figure 2.19: "Effective" average age as a function of scale length [left] for all 158 galax-ies and rotational velocity [right], VTOt for all 98 galaxies for which we have reliable VTOt measurements. 21 20 18 18 1? MO,H ,K ( M A S a r c s e c " 2 ) 19 - 20 -21 -22 -23 -24 -25 -26 -21 M H ,K (mag) Figure 2.20: "Effective" metallicity \ogw(ZeS/ZQ) as a function of central surface bright-ness [left] and total Hm, if-band galaxy magnitude [right] for all 158 galaxies. See caption of Figure 2.17 for the definition of Hm. Chapter 2. Colour Gradients in Spiral Galaxies 53 log(h (kpc)) V . ( k m s-i) Figure 2.21: "Effective" metallicity, \ogi0(Zeg/ZQ), as a function of scale length [left] for all 158 fitted galaxies, and rotational velocity [right], VI0t for all 98 galaxies for which we have reliable VTOt measurements. \og10(Zeg/ZQ) increases with VI0t up to ~ 120 k m s - 1 and then levels off with smaller scatter. There is no dependence of log 1 0(Z eff/Z o) with h. 2.8.2 Global Trends in Age and Metallicity Gradients One of the conclusions from BdJOO is that the amplitude of the age gradients increases from HSB to LSB galaxies. However, this is likely an artifact of their linear gradient fitting technique out to a different number of radial bins (scaled by the galaxy disk scale length), anywhere from 2-4 bins (i.e. out to 1.5-3.5 scale lengths). Assuming that the LSB galaxy photometry generally does not extend to as many scale lengths as the HSB photometry, non-linear gradients in the galaxies (e.g. steeper in the inner regions, flattening at large radii), could mimic a trend with SB. To demonstrate this, in Figures 2.22 & 2.23, we plot average age, {A}, and metallicity, \og10(Z/ZQ), respectively, as a function of radius. In each case, the left panels give radius in terms of disk scale lengths, and the right panels have the physical radius in kpc. Modulo fairly large uncertainties, we see that gradients are not linear over the extent of the galaxy. The inner and outer slopes are often quite different, sometimes even changing sign. Also, the Chapter 2. Colour Gradients in Spiral Galaxies 54 Rad (h H K ) Rad (kpc) Figure 2.22: Age as a function of radius for all galaxies. The left plot gives the radius in terms of the measure scale length hH<K while the right plot gives the radius in physical units (kpc). Dotted lines represent our sample whereas solid lines are the BdJOO sample. Rad (h H K ) Rad (kpc) Figure 2.23: Same as Figure 2.22 but for metallicity. Chapter 2. Colour Gradients in Spiral Galaxies 55 galaxies are not all measured out to the same number of scale lengths. When plotted against number of scale lengths, most galaxies have steeper (negative) inner gradients that flatten off, or even turn around, in the outer parts. With the exception of extremely early or late-types, the outer slopes are quite similar for most galaxies. However, when plotted as a function of physical scale (kpc), the slopes can be quite different, with a rough trend for the earliest types being steepest, and leveling off for later-types. Figures 2.22 & 2.23 demonstrate that fitting gradients out to a different number of scale lengths yields misleading results. In general, the larger the fit baseline, the shallower the measured gradient. We verify this assertion in Figure 2.24 where we plot age gradients fit out to different numbers of radial bins for all the galaxies as a function of CSB [0-1.5hH,K [top left], 0-2.5hHtK [top right], 0-3.5hH,K [bottom left], and 0-4.5HH,K [bottom right]). Indeed, fitting fewer bins results in systematically larger 1 1 . 1 1 1 . 1 , m i • 1 • I ' fit: C 1 i i , 1 , i , ) - ZShw _ i , , . i , - i — i — I — i — i — r ~ fit: 0 - 4.5hH , . . I , 20 18 16 22 20 18 MO.H K ( m a § arcsec- 2 ) 16 Figure 2.24: Age gradient as a function of central surface brightness, fio^^K for all galax-ies. The gradients are fit out to a different number of radial bins: 0 - 1 .5 A ^ y [top left], 0-2.5hH,K [top right], 0-3.5hH>K [bottom left], and 0 -4 .5 h H , x [bottom right]. Chapter 2. Colour Gradients in Spiral Galaxies 56 (usually negative, but occasionally positive) age gradients. The more extended the baseline, the flatter the gradient. Because LSB galaxy profiles do not extend as far as HSBs (in scale lengths, as indicated by the lower CSB galaxies dropping out of the longer baseline fits), they do appear to have larger gradients (which led to the incorrect interpretation by BdJOO). Given the dangers of fitting gradients out to different radial extents, in the remainder of this analysis we consider "inner" and "outer" gradients fitted over the Ist-2nd radial bins and 2nd-4th bins, respectively6. This restriction greatly reduces the available galaxy parameter range, especially for the LSB galaxies. Also note that inner gradients will be much more affected by dust (if present) than outer gradients. 2.8.3 Age Gradients In Figure 2.25 we plot the inner [left panels] and outer [right panels] gradients mea-sured in disk scale lengths [upper panels] and kpc [lower panels] as a function of CSB [left plot] and total magnitude [right plot]. Unlike BdJOO, we do not see a trend of d(A)/dr (Gyv/h) (inner or outer) with CSB. Note that the low CSB galaxies are miss-ing in the outer gradient plots so one cannot assess any trend for the outer gradients with great confidence. The current data suggest that outer gradients are generally smaller than those within. When age gradients are plotted as a function of kpc, the inner gra-dients are slightly steeper for higher CSB galaxies (/io,ifm,/c ~ 17.5 mag arcsec -2). The 3-4 outliers at high positive gradients are the SO galaxies whose colour gradients are inverted (bluer inward). Such anomalous gradients have been seen before in SO galaxies (e.g. Emsellem et al. 2002) and are often interpreted as the result of a recent gaseous accretion followed by (central) star formation. These are also the few galaxies that deviate to much younger ages from the general trend of decreasing effective age with galaxy type (see left plot of Fig. 2.30) which would agree with the above interpretation of recent accretion and SF. There is a slight trend (with large scatter) for more luminous galaxies to have 6 Note that these two definitions cannot be interpreted as bulge and disk gradients as the inner fit goes out to 1.5 disk scale lengths, and bulge scale lengths are typically ~0.13 /i<* (Paper I). Chapter 2. Colour Gradients in Spiral Galaxies 57 22 21 20 IS 18 1? 16 15 22 21 20 19 18 17 18 15 -18 -20 - 2 2 - 2 4 - 2 8 -18 -20 - 2 2 -24 - 2 8 MO.H,.K ( m a S arcsec"2) M H _ K (mag) Figure 2.25: Age gradient as a function of Hm,K-band central surface brightness, po,Hm,K [left] and total magnitude, MHmjK [right] for all galaxies. The fit ranges represent "in-ner'gradients [left panels] fit out to 1.5H,K-band disk scale lengths, hn,K, and "outer" gradients [right panels] fit from 1.5 to 3.bhn,K- The upper and lower plots show gradients as a function of scale length and kpc respectively. steeper d(A)/dr (Gyr/h) negative gradients (inner and outer), and the fainter galaxies extending to positive gradients. However, this trend disappears when the gradients are plotted as a function of kpc. In this scale the outer gradients are greatest (negative or positive) for the fainter galaxies, and become negligible for the brighter ones. A similar, somewhat stronger trend of steepening gradient (in scale lengths) with increasing disk scale length is seen in Figure 2.26 [left]. The trend in d(A)/dv (Gyr/kpc) is reversed (ignoring the 3-4 steep positive gradient points) such that bigger galaxies have smaller gradients. This may reflect the larger galaxies as being less likely to have had any recent activity/interactions, as the small gradients combined with the older ages for the bigger galaxies suggest an overall old age for the largest galaxies. Similar trends are seen with VTOt (as expected from the correlations between VTOt with h and M in Fig. 2.3), but they are weaker (largely because of the smaller sample). Chapter 2. Colour Gradients in Spiral Galaxies 58 - 0 . 5 0 0.5 1 - 0 . 5 0 0.5 1 50 100 150 200 250 50 100 150 200 250 log(hHK (kpc)) V r o t (km s-i)) Figure 2.26: Same as Figure 2.25 except for disk scale length, \ogw(hH,K) [left], and rotational velocity, VI0t [right]. 2.8.4 Metallicity Gradients Figures 2.27 & 2.28 recast the information presented in Figures 2.25 & 2.26 discussed above, but now in terms of metallicity gradients. Metallicity gradients are less well measured and are more sensitive to dust effects than age gradients (according to the triplex dust models discussed in §2.5). Nevertheless, the trends in the metallicity gra-dients with galaxy parameters are very similar to those with age gradients (see §2.8.3). The trends in the magnitudes of the gradients as a function of galaxy parameters are comparable except the correlation changes sign. This suggests a direct anti-correlation between age and metallicity gradients within a galaxy which is seen in Figure 2.29 (al-though with some scatter), such that stronger age gradients are associated with weaker metallicity gradients (and vice versa). 2.9 Discussion Our main findings thus far can be summarized as follows: (i) Our relative age determinations are robust under the assumption that the underly-Chapter 2. Colour Gradients in Spiral Galaxies 59 Figure 2.28: Same as Figure 2.26 except for metallicity gradients. Chapter 2. Colour Gradients in Spiral Galaxies 60 Figure 2.29: Metallicity gradient as a function of age gradient for all 158 galaxies. ing SFR of all disk galaxies is similar, and that no major star bursts have occurred in the past 1-2 Gyr (which is justified on the basis of Kauffmann et al. (2003b)); (ii) Dust cannot be ruled out as a contributor to colour gradients in spiral galaxies, although it is unlikely that the gradients are largely due to dust extinction; (iii) Age and metallicity correlate strongly with local surface brightness: higher SB regions tend to be older and more metal-rich (see Fig. 2.17). This indicates that the local potential plays a significant role in the SFH of spiral galaxies; (iv) Age and metallicity, measured at an effective radius of Ih, increase with ear-lier Hubble type, MHm,K, Kot, and A*o,Hm,i<: but the trends flatten for T < 4, MHintK< - 22.5 mag, VTOt > 120 kms - 1 , and £t0,fwr £ 18-5 mag arcsec-2 (see Figs. 2.18-2.21 & 2.30). Age also correlates weakly with scale length (with scat-ter decreasing with galaxy size, see right plot of Fig. 2.19); (v) Age and, to a lesser extent, metallicity gradients show radial structure, with gen-erally steeper gradients in the inner parts of the galaxy. Care must thus be taken when defining the gradient fit region (see Figs. 2.22-2.24); (vi) Age gradients, as measured as a function of scale length, show correlations with Chapter 2. Colour Gradients in Spiral Galaxies 61 luminosity, size, and rotational velocity (Fig. 2.25), especially in the inner disk. Trends in metallicity gradients with galaxy parameters are weaker (Figs. 2.27 & 2.28); (vii) Age gradients do not correlate with CSB, contrary to the findings of BdJOO, which we attribute to inconsistencies in their fit radii (see left plot of Fig. 2.25). This is not at odds with statement (iii) since po and h are not strongly correlated. These observations are consistent with a picture where (i) higher surface brightness regions of galaxies formed their stars earlier than lower surface brightness regions, or at a similar epoch but on shorter timescales, and (ii) the SFHs at a given SB level, which lead to age gradients, are modulated by the overall potential of the galaxy such that brighter/higher rotational velocity galaxies formed earlier. These trends reach saturation for the brightest and highest CSB galaxies. An earlier formation time for more massive galaxies is clearly in conflict with current models of hierarchical galaxy formation which predicts that more massive galaxies form late. This discrepancy could be remedied if there is a mechanism at work that prevents the gas in lower mass galaxies from cooling and forming stars at early times. Feedback processes are often invoked as a possible solution, but no prescription has yet been found that agrees with all observational constraints (e.g. Kauffmann et al. 2003b; Bell et al. 2003). In a related manner, the semi-analytic models of hierarchical galaxy formation of Kauffman (1996) predict a correlation between bulge-to-disk ratio and luminosity-weighted age, such that the bulges of late-type spirals should be older (up to ~ 4 Gyr) than those of early-type spirals, although this correlation could be erased by any signif-icant inflow from the disk after the last major merger. Figure 2.30 [left] shows a strong correlation of effective age with galaxy type, but in the opposite sense of Kauffmann's prediction; later-type galaxies have systematically younger effective ages. Clearly, mix-ing by a bar could erase the predicted age trend. If bars trigger radial flows the gradients of strongly barred galaxies would be flat (the flattening would occur on short enough timescales). To look for such mixing effects, we plot age and metallicity gradients as Chapter 2. Colour Gradients in Spiral Galaxies 62 A < V 6 8 10 Type Type Figure 2.30: Effective age [left] and metallicity [right] as a function of morphological type for all 158 galaxies. >> o b -o \ A < V -a a a. J£ \ b b -o \ A < V -a I 1 I ' I 1 I 1 I T I 1 I J I 1 I 1 I 1 I 1 I ' I 1 I 1 I i . i . i . t . i i ^ inner: 0 - 1.5hHK . I i I i I i I i L I 1 I 1 I 1 I 1 I 1 I 1 I outer: 1.5 - 3.511,^  i • i • i . I . I . I . i - 4 - 2 0 2 4 0 10 -4 -2 0 2 4 6 8 10 Type -a o no 0 O 0.5 a i 1 i 1 i 1 i 1 i 1 i • i ii ' • ' • i • ' • i ' • i • i • i • i • ' • ' b -a o OB O -1.5 I ' I ' I 1 I 1 I 1 i !i.,|f| inner: 0 - liShj^ i . I . I . I . I . I . i t « | i j f i outer: 1.5 - 3.511^ J -4 -2 0 2 4 10 -4 -2 0 2 4 Type Figure 2.31: Age gradient [left] and metallicity gradient [right] as a function of morpho-logical type for all 158 galaxies. The upper panels plot radius in terms of the measured scale length hn,K while the bottom panels plot radius in physical units (kpc). Chapter 2. Colour Gradients in Spiral Galaxies 63 a function of barredness in Figure 2.32. No trends are seen when the gradients are 0.6 0.4 0.2 5 0-6 o 3 0.4 £ 0.6 0.4 0.2 0.6 0.4 o 0.6 2 0.4 £ 0.2 EL. 0.6 0.4 0.2 _l 1 1 1 1 1 1 1 1 , • i . i . ij-rrTr71TL A -j L . 1 . " _| i | i | i | i | i " | , rt-i, r+JTl 1 ! . iVl - K f 1 | 1-AB "-i . i , i ,• -I 1 1 1 1 1 1 1 1 1 • i L ,» i rrr i i v* ' l ' l -B : Til i I r -I i I i I i I i I i : P 1 i 1 i 1 , 1 J l I 1 1 1 1 '-i 't-f-r 1 i ' _| i | i | i | f | i 1 i 1 i 1 i 1 h 1 i |*',|, •' 1 '-i 1 i ! i ' _| I ] I | I | ' i | I " i . i . i i i n r, j ' l ] 1 j l_T V - h , 1 i _ - 5 - 4 - 3 - 2 - 1 0 1 2 3 - 5 - 4 - 3 - 2 - 1 0 1 2 3 d<A>/dr (Gyr/h) . • 1 1 1 1 - i ..^rr 1 I 1 I 1 I 1 I -A 1 L i i 1 i 1 I + -. 1 1 • i • i+1 • 1 1 1 1 1 1 T. AB 1 L i 11 i 11 i• 1 innerj1" ~ \ 1 i 1 i 1 i. 1 ° ] : ! ; ! , 1 ; 1 -L . J 1 l 1 I 1 I 1 I . , 1 i + i i , 1 • . 1 i ' i 1 • i i i i T + i I i T 1 1 1 1 . . I . I . I . I -. 1 i 1 i 1 1 outer - , 1 , 1 , j 1 1 1 1 1 1 1 1 . L . 1 , 1 . 1 -- 3 - 2 - 1 0 1 2 3 4 - 3 - 2 - 1 0 1 2 3 4 d<A>/dr (Gyr/kpc) 0.4 0.2 2 °'4 § 0.2 e u fc. 0.4 0.2 _ l 1 1 1 1 1 1 1 1 , i i i _ A J rr-p i " —1—I I 11 ] 11 Ml—1— i i 11-_ AB J I T H - I ' J _ J . , , 1 1 1 l _ B : % o 3 0.5 O to u 6- 0 0.5 1 1 1 1 1 1 1 1 [ 1 1 1 1 j 1 1 ! ! i ,t, J—i rh 111—i—i—1—i_ta=u—1_ dlog(Z/Z 0 ) /dr (/h) 1 1 1 1 1 I 1 I I 1 1 • 1 1 1 1 1 1 1 1II -rJ 1  | 1   | I' lTI i i f i _ 1 AB j T r I I' r ' , j • inner r- 1 • , 1 , , , , 1 _ , rJ , B 1 -0.5 0.5 _ I 1 I 1 1 1 1 I , , I - , 1 , , , , 1 , -r-L , i U _ 1 | 1 1 1 1 | 1 X 1 1 j _ 1 | 1 1 1 1 | 1 1 I'I4 ' outer | • , i , , , , i , , ,1, i " P i l l -0.5 0 0.5 -1 -0.5 dlog(Z/Z e ) /dr (/kpc) 0.5 Figure 2.32: Age gradient [left] and metallicity gradient [right] histograms as a function of barredness for all 158 galaxies. The upper panels plot the radius in terms of the measured scale length hHtK while the bottom panels plot the radius in physical units (kpc). The fit ranges represent "inner"gradients [left panels] fit out to 1 . 5 a n d "outer" gradients [right panels] fit from 1.5 to 3.5 hH>K. The vertical dotted lines are located at zero gradient for reference. plotted in scale lengths, but against kpc, the strongly barred galaxies may have smaller age gradients (inner and outer), though this observation is based on small statistics. This agrees with Martin & Roy (1994) who found similar flattening of O/H metallicity gradients with bar strength (their trend also disappears when the gradients are plotted in scale lengths). Thus, under all assumptions, our observations cannot be reproduced by the Kauffmann (1996) models. The same model deficiency responsible for the back-wards age-size correlation, as compared against observations, could also be the cause of the discrepancy in bulge morphology versus age seen in the Kauffmann (1996) models. Prantzos & Boissier (2000; hereafter PBOO) present chemo-spectrophotometric mod-els of spiral galaxy evolution. They adopt a Schmidt-type law for the SFR, which is proportional to the gas surface density and varies with galactocentric radius (dynamical time), assume an inflow rate (of unenriched gas) that decreases exponentially with time and increases with surface density and galaxy mass, and assume that the gas settles into Chapter 2. Colour Gradients in Spiral Galaxies 64 an exponential disk (bulges are not modeled), but do not consider radial flows (e.g. by viscosity which could create gradients, or due to a bar which could flatten radial gradi-ents on small timescales). The SF efficiency and infall rate are free parameters tuned to match the Milky Way observational constraints. The models are then extended to other disk galaxies by adopting the Cold Dark Matter(CDM)-based scaling relations of Mo, Mao, & White (1998). As such, their galaxy disk radial profiles are fully described by just two parameters: rotational velocity, VTOt (assuming a constant disk-to-halo mass ratio), and the halo spin parameter, A. A third parameter describing the formation redshift would be required for a description fully consistent with the CDM hierarchical models of galaxy formation but, as of yet, there is no clear definition for the time of for-mation for an individual galaxy. Hence, PBOO assume that all galaxies started forming at the same epoch (13.5 Gyr ago) but evolve at different rates. The PBOO models give predictions for O/H abundances and gradients as could be measured in bright H II regions in nearby galaxies. Note that O/H determinations in H II regions probe the present-day ISM abundances and are insensitive to abundance evolution with time. Still, their model predictions and observations of O/H gradients can be compared, at least indirectly, to our stellar luminosity weighted gradients. PBOO find that the absolute central abundance for a given A > 0.03 correlates with VTOt (and total magnitude) such that faster rotators have larger central abundances, but the trend saturates for V ,^t ^ 220 kms - 1 . Furthermore, at a given V ^ , the central abundance decreases with increasing A but again saturates for galaxies with VTOt ^ 220 k m s - 1 above which central abundances are high, regardless of A. By extension, we should see a trend of increasing central abundance with VI0t, with significant spread (larger at lower Kot) due to different values of A, but this trend would flatten and show less dispersion above VTOt ^ 220 kms - 1 . This is roughly what Figures 2.20 & 2.21 show for the effective metallicity at the half-light radius versus MHjntK and VTOt (central values are too uncertain due to likely higher concentrations of dust and seeing mismatches, and PBOO's models do not consider bulges). Unfortunately, our galaxies do not exceed Vroi ~ 250 kms - 1 , but the general trend above, including the decrease in scatter with Chapter 2. Colour Gradients in Spiral Galaxies 65 Kot, is confirmed. The right side of Figure 2.27 shows our measured metallicity gradients as a function of total galaxy magnitude; gradients in disk scale lengths are shown in the upper panels while gradients in kpc are shown underneath. The left and right panels show the "inner" and "outer" gradients respectively. When plotted against scale length, the inner or outer gradients show no clear trends. However, when plotted in kpc, a trend emerges such that brighter galaxies have smaller metallicity gradients. Since scale length increases with luminosity (see Fig. 2.3), the correlation goes away when the gradients are measured per disk scale length. If this effect is real, the disappearance of a trend of metallicity gradients in units of scale length may suggest a self-similar pattern in disk galaxies. As suggested by Combes (1998), a "universal" slope per disk scale length might be explained by the viscous disk models of Lin & Pringle (1987), although model predictions are not conclusive as of yet. The Garnett et al. (1997) compilation of O/H gradients shows a similar signature, where abundance gradients in dex/kpc are steeper and exhibit greater scatter for lower luminosity disks, but this trend goes away when gradients are expressed in dex/h. Note, however, that the abundance gradients compiled in van Zee et al. (1998) and reported in PBOO do show a clear increase with magnitude when plotted as dex/h, thus thwarting the interpretation of a universal abundance gradient per scale length. We have measured compelling trends for the effective age and metallicity in spiral galaxies and their inner and outer disk gradients as a function of surface brightness, luminosity, rotational velocity, and size. These trends are not well reproduced by cur-rent semi-analytical models of galaxy evolution but will undoubtedly serve as effective constraints for future models. An important limitation inherent in our colour-based analysis is its sensitivity to dust extinction which could mimic gradients in age and, particularly, metallicity. Given these limitations, a spectroscopic analysis of radially re-solved line indices of spiral galaxies could offer great promise to alleviate this problem, and in the following chapters we discuss our pilot study of absorption-line indices in late-type spirals. 66 Chapter 3 Dust Sensit ivi ty of Absorp t ion-Line Indices 1 3.1 Introduction Stellar population studies play a fundamental role in our understanding of stellar evo-lution, initial mass functions (IMFs) associated with star cluster formation, and galaxy formation and evolution. With the aim of determining the star formation histories (SFHs) of stellar systems of all types (from star clusters to entire galaxies), many tech-niques have been developed, with varying success, to determine the luminosity-weighted ages and metallicities of nearby and distant stellar systems. Important caveats hinder or limit the application of these techniques (e.g. Chariot, Worthey, & Bressan 1996; MacArthur et al. 2004 [§2]; Anders et al. 2004; Tantalo & Chiosi 2004). For example, the determination of ages and metallicities of stellar populations has been plagued by the well-known age/metallicity degeneracy (Worthey 1994). Access to both optical and infrared imaging partially lifts this degeneracy, but broad-band colours suffer further from a degeneracy due to dust reddening (e.g. Bruzual, Magris, & Calvet 1988; Witt, Thronson, & Capuano 1992; de Jong 1996; Bell & de Jong 2000; MacArthur et al. 2004 [§2]), which is particularly problematic in the context of gas rich, star forming systems. Another method of breaking the age/metallicity degeneracy of stellar populations in unresolved systems involves measurement of surface brightness fluctuations (Worthey 1993; Liu, Chariot & Graham 2000; Blakeslee, Vazdekis &: Ajhar 2001), which are most sensitive to contributions from the most luminous stars in the system - typically evolved cool giants. Being sensitive to the second moment of the stellar luminosity function, x The analysis presented in this chapter is an expansion of the work published in The Astrophysical Journal as MacArthur 2005, A p J , 623, 795. Chapter 3. Dust Sensitivity of Absorption-Line Indices 67 surface brightness fluctuations provide complementary information to the integrated colours (the first moment of the stellar luminosity function). This method, however, requires the assumption of a smooth underlying light distribution, and thus is generally only applicable to studies of nearby globular clusters, ellipticals, and spiral bulges. This broad-band technique is also not applicable to dusty systems (e.g. spiral galaxies), as the resulting dumpiness will also cause fluctuations in surface brightness and redden-ing will affect the fluctuation colours. Turning to spectroscopy offers the possibility of overcoming the ambiguities of broad-band colour-based analyses (see §2) by studying the variation of individual line strengths which, if defined over a sufficiently narrow wavelength range, should be insensitive to the effects of dust reddening. In this chapter we test this conjecture by exploring the effects of dust absorption on commonly-used spectroscopic age and metallicity indicators. In the past three decades, tremendous progress has been made in the modeling of simple stellar populations (SSPs) (e.g. Tinsley 1972; Bruzual &; Chariot 1993; Worthey 1994; Vazdekis 1999; Fioc & Rocca-Volmerange 1997; Maraston 1998; Bruzual & Char-lot 2003). The current versions provide numerous observables such as high-resolution spectra, magnitudes, colours, mass-to-light ratios, and line-index measurements for sin-gle bursts of star formation (SF) with a given IMF and metallicity at ages ranging from 0-20 Gyr. In particular, the Bruzual & Chariot (2003) models include a prescription for dust attenuation (described in §3.3) that we adopt for this analysis. Numerous studies of absorption-line indices in the integrated spectra of composite systems have sought to disentangle and constrain the ages and metallicities of their stellar populations (e.g. Maraston & Thomas 2000; Trager et al. 2000; Schiavon et al. 2002b; Caldwell, Rose, & Concannon 2003). The majority of these studies have fo-cused on globular clusters (GCs) and elliptical galaxies, where dust is not conspicuous. Recently, however, a number of studies have turned their attention to the stellar pop-ulations of spiral galaxies, where dust is an ineluctable hindrance (Fisher, Franx, & Illingworth 1996; Goudfrooij, Gorgas, &; Jablonka 1999; Trager, Dalcanton, & Weiner 1999; Proctor & Samson 2002; Bergmann, J0rgensen, & Hill 2003; Kauffmann et al. Chapter 3. Dust Sensitivity of Absorption-Line Indices 68 2003a; Falcon-Barroso et al. 2003). Reddening by dust has often been assumed to have a negligible effect on line indices, but no rigorous study has yet been performed to verify this conjecture. This work presents the first such analysis, and provides prerequisite conditions for the validity of our spectroscopic analysis of absorption-line indices in spiral galaxies presented in the following chapter (§4). The outline of this chapter is as follows; a description of the line-indices studied and their applicability as age and metallicity discriminators is given in §3.2. The stellar population synthesis models and dust prescription are described in §3.3. In §3.4 we discuss the dust sensitivities of the Lick & higher-order Balmer line indices, the 4000 A break, the Ca II triplet indices, and the Rose indices. In §3.5 we present age and metallicity fits to the model galaxies in several index-index planes and discuss the errors on the derived physical parameters due to dust extinction. Finally, the results are summarized in §3.6. 3.2 Description of Indices Studied 3.2.1 The Lick/IDS System The Lick/IDS system of 21 spectral line indices was designed to calibrate the strength of fundamental spectral features in stars and composite systems (e.g. Gorgas et al. 1993). The indices measure the strength of a particular spectral feature (either atomic and defined as an equivalent width in A, or molecular and measured in magnitudes) relative to a pseudo-continuum on each side of the feature. The most reliable indices have been calibrated as a function of stellar colour (effective temperature), surface gravity, and metallicity (Gorgas et al. 1993; Worthey et al. 1994) allowing for the construction of semi-empirical population models (e.g. Worthey 1994). The Lick indices are sensitive to the metallicity and age of stellar populations to varying degrees. When compared with population models, diagnostic plots of age versus metallicity sensitive indices, such as H/3 versus Mgb or (Fe), help break the age-metallicity degeneracy. However, measurements of many of the Lick indices are quite sensitive to spectral resolution and, thus, to the velocity dispersion of the system (Gonzalez 1993; Trager et al. 1998; Chapter 3. Dust Sensitivity of Absorption-Line Indices 69 Proctor & Sansom 2002), and their use requires relatively high signal-to-noise data (S/N > 50/A; see Cardiel et al. 1998). In addition, certain indices, H/3 in particular, can suffer "in-filling" from nebular emission contamination, present even in early-type galaxies (Gonzalez 1993; deZeeuw et al. 2002; Caldwell, Rose, & Concannon 2003). In order to measure a given spectral index, three wavelength intervals must be defined: a central bandpass covering the feature of interest, \c1-\c2, and a pair of bracketing bandpasses, \b1-\b2 and Xri-\r2 for the blue and red side of the central bandpass respectively, to define a continuum level. The spectral index is then a measure of the strength of the spectral feature based on the flux in the central bandpass com-pared to that of the continuum level which is linearly interpolated from the bracketing bandpasses. In late-type spectra there is virtually no region entirely free of absorption, but to measure an empirical line strength a true continuum is not necessary, as long as the "pseudo-continuum" bandpass region does not have very strong or highly variable lines. The 21 original Lick indices are defined in Table 3.1. Explicitly, in the pseudo-continuum bandpasses, the average flux in the blue and red passbands are calculated as: where A61, Xb2 and Ar x, Ar 2 are the limits of the blue and red pseudo-continuum band-passes respectively. The local continuum for a given index, FC(A), is then the run of flux defined by drawing a straight line from the midpoint of the blue continuum level to the midpoint of the red continuum level, where A6 0 = (A62 — A&i)/2 and Ar 0 = (Ar2 — Xri)/2 are the centers of the blue and red pseudo-continuum bandpasses respectively. (3.1) (3.2) Atomic indices are measured as (pseudo-)equivalent widths in A as: (3.3) Chapter 3. Dust Sensitivity of Absorption-Line Indices 70 Name Central Bandpass Pseudocontinua Type Features Aci --Ac 2 A&x--Xb2 (Units) Measured Arv -Xr2 CNi 4142.125 -4177.125 4080.125 -4117.625 molecular CN,(0) 4244.125 -4284.125 (mag) C N 2 4142.125 -4177.125 4083.875 -4096.375 molecular CN,(0) 4244.125 -4284.125 (mag) Ca4227 4222.250 -4234.750 4211.000 -4219.750 atomic Ca,(C) 4241.000 -4251.000 (A) G4300 4281.375 -4316.375 4266.375 -4282.625 atomic CH,(0) 4318.875 -4335.125 (A)_ Fe4383 4369.125 -4420.375 4359.125 -4370.375 atomic Fe,C,(Mg) 4442.875 -4455.375 (A) Ca4455 4452.125 -4474.625 4445.875 -4454.625 atomic Ca,(Fe),(C),Cr 4477.125 -4492.125 (A) Fe4531 4514.250 -4559.250 4504.250 -4514.250 atomic Ti,(Si) 4560.500 -4579.250 (A)_ Fe4668 4634.000 -4720.250 4611.500 -4630.250 atomic C2,(0),(Si) 4742.750 -4756.500 (A) H/3 4847.875 -4876.625 4827.875 -4847.875 atomic H/3,(Mg) 4876.625 -4891.625 (A) Fe5015 4977.750 -5054.000 4946.500 -4977.750 atomic (Mg),Ti,Fe 5054.000 -5065.250 (A) M g l 5069.125 -5134.125 4895.125 -4957.625 molecular C.MgH, 5301.125 -5366.125 (mag) (0),(Fe) Mg 2 5154.125 -5196.625 4895.125 -4957.625 molecular MgH,Mg6,C 5301.125 -5366.125 (mag) (Fe),(0) Mg6 5160.125 -5192.625 5142.625 -5161.375 atomic Mg&,(C),(Cr) 5191.375 -5206.375 (A) Fe5270 5245.650 -5285.650 5233.150 -5248.150 atomic Fe,C,(Mg) 5285.650 -5318.150 (A)_ Fe5335 5312.125 -5352.125 5304.625 -5315.875 atomic Fe,(C),(Mg), 5353.375 -5363.375 (A) Cr Fe5406 5387.500 -5415.000 5376.250 -5387.500 atomic Fe 5415.000 -5425.000 (A) Fe5709 5696.625 -5720.375 5672.875 -5696.625 atomic (C),Fe 5722.875 -5736.625 (A) Fe5782 5776.625 -5796.625 5765.375 -5775.375 atomic Cr,Fe 5797.875 -5811.625 (A) Table 3.1: Lick Index Definitions. Chapter 3. Dust Sensitivity of Absorption-Line Indices 71 Name Central Bandpass ' Pseudocontinua \C\-\C2 A&1-A&2 Ari - A r 2 Type (Units) Features Measured Na D 5876.875 -5909.375 5860.625 -5875.625 atomic Na,C,(Mg) 5922.125-5948.125 (A) TiOi 5936.625-5994.125 5816.625-5849.125 molecular TiO,C 6038.625-6103.625 (mag) T i 0 2 6189.625-6272.125 6066.625-6141.625 molecular TiO,C,V,Sc 6372.625-6415.125 (mag) where Fr(X) is the flux per unit wavelength in the index passband. Molecular indices are measured in magnitudes as the mean ratio of Fj(A) to FC(X) in the central bandpass: The above scenario for the index measurements is shown graphically in §3.4 in Fig-ures 3.5 (for the H/3 index) & 3.6 (for the Mgi index) where the vertical green lines delineate the blue and red pseudo-continuum bands (dashed lines) and the central band-pass (solid lines), the dashed lines represent Ft, and Fr, the dotted lines joining the black circles (Xb0 and Aro) is FC(X). 3.2.2 Higher Order Balmer Indices In order to overcome the problem of nebular emission fill in of the H/3 feature, Worthey & Ottaviani (1997; hereafter W097) introduced two pairs of indices that measure the higher-order Balmer lines H7 and B.5. Two definitions for each feature were defined (see Table 3.2), the narrower version (AA ~20 A) is denoted with the subscript "F" (as it encompasses all of the Balmer line absorption from stars of spectral type F at the Lick/IDS resolution of 8-10 A), e.g. B.8F, and the wider definition (AA ~ 40 A) has the subscript "A" (as it includes all of the absorption from A stars), e.g. H-JA- The narrow indices are more age-sensitive, but require higher S/N and resolution. While their age-sensitivity is not as strong as for H/3, the higher-order Balmer lines are much less affected Table 3.1: Lick Index Definitions - continued I mag = -2.5 log10 (3.4) Chapter 3. Dust Sensitivity of Absorption-Line Indices 72 by emission from ionized gases (e.g. Osterbrock 1989). Thus, when combined with a metallicity sensitive index, the W097 indices provide a more reliable age estimate for star-forming galaxies. Note, however, that the W097 higher-order Balmer lines are poorly calibrated (Vazdekis 1999), likely due to the degrading resolution at the blue end of the IDS data. Name Central Bandpass Pseudo-continua Type Aci - Ac2 A&i -A6 2 (Units) AT-! - A r 2 4083.500-4122.250 4041.600 -4079.750 atomic 4128.500 -4161.000 (A) H 7 A 4319.750-4363.500 4283.500 -4319.750 atomic 4367.250 -4419.750 (A) B8F 4091.000-4112.250 4057.250 -4088.500 atomic 4114.750 -4137.250 (A) 4331.250-4352.250 4283.500 -4319.750 atomic 4354.750 -4384.750 (A) Table 3.2: W097 Higher-Order Balmer Line Indices. 3.2.3 The 4000 A Break Another widely used spectral index that obviates the need for high S/N and spectral resolution is the 4000 A break, a flux ratio that brackets the strongest discontinuity in the optical spectrum of a galaxy. The break arises due to the accumulation of a large number of spectral lines in a narrow wavelength region bluewards of 4000 A in stellar types cooler than GO (Bruzual 1983; Gorgas et al. 1999). The main contribution to the opacity just blueward of 4000 A comes from atomic metals (Fe I, Mg I, Ca II) and molecular CN, which decreases for hotter and more metal poor stars. Thus the 4000 A break is weak for young and/or metal-poor stellar populations and strong for old, metal-rich galaxies (Kauffmann et al. 2003a). In its original form (Bruzual 1983), the 4000 A break, denoted D(4000), was defined as the ratio of the average fluxes per frequency unit measured over the spectral ranges 4050-4250 A and 3750-3950 A. Cardiel et al. (1998) demonstrated that this discontinuity can be measured with a relative error of Chapter 3. Dust Sensitivity of Absorption-Line Indices 73 ~10% with a S/N per A ~ l , thus making it better suited for lower quality data. However, the D(4000) does have a few drawbacks due to its long baseline. These were partially alleviated with the introduction of a narrower definition (Balogh et al. 1999) denoted D„(4000) and measured over the ranges 4000-4100 A and 3850-3950 A. The narrow definition was designed to exploit two principal advantages: an improved agreement between multiple measurements of a given galaxy, and a weaker sensitivity to reddening effects. However, using their narrow D„(4000) index, Balogh et al. (1999) still had to invoke dust reddening as a cause for the large D„(4000) values in a number of their z ~ 0.3 galaxies, i.e. D„(4000) is not impervious to dust effects. Nevertheless, often found in the literature is the statement that the Dn(4000) is insensitive to dust attenuation effects. For example, Kauffmann et al. (2003b) use the amplitude of the Dn(4000) in combination with the strength of the R8A index of W097 as diagnostics for the SFH of the host galaxies, from which they infer the dust attenuation by comparing observed to model colours, a method that relies heavily on the assumption of dust insensitivity of the Dn(4000) and ESA indices. We will confirm in §3.4.2 that Dn(4000) is definitely affected by dust extinction to a degree that could seriously compromise SFH determinations. 3.2.4 The Near-IR Ca II Triplet Indices There has been a great effort recently to extend stellar population studies to the near-IR region, focusing on the Ca II triplet as one of the most prominent features in the near-IR spectrum of cool stars (from spectral types of about F5 to M2) (Cenarro et al. 2001a,b; Vazdekis et al. 2003). Three "Lick-style" indices designed to measure the strengths of the Ca II triplet lines (AA8498.02, 8542.09, 8662.14 A), and denoted Cal , Ca2, and Ca3, have been defined and redefined by several authors (e.g. Jones, Alloin, & Jones 1984; Armandroff & Zinn 1988; Diaz, Terlevich, & Terlevich 1989, hereafter DTT; Delisle & Hardy 1992). However, measurement of a reliable continuum for these indices is difficult due to the strong and crowded absorption features in the vicinity of the Ca II lines (largely from Fe I, Mg I, and TiO), as well as significant blending with the hydrogen Paschen series whose absorption is present in stars of types G3 and Chapter 3. Dust Sensitivity of Absorption-Line Indices 74 hotter. To overcome these problems, Cenarro et al. (2001a) defined a new set of Ca II triplet indices specifically designed for measurement in integrated galactic spectra. The new definitions are categorized as "generic" indices, which have advantages over the classical "Lick-style" indices when looking at adjacent absorption lines in regions where the continuum is crowded with spectral features. Their "CaT" index includes the strenghts of all three Ca II lines and uses a combination of 5 continuum bandpasses (see Cenarro et al. 2001a for details on the measurement of generic indices). Additionally, they define the generic "PaT" index, which measures the strength of three of the H Paschen series lines that are free from Ca contamination. Finally, the CaT* index is designed to remove the H Paschen line contamination from CaT making it a reliable indicator of the pure Ca II triplet strength. The index is given by CaT* = Ca —0.93 PaT. Vazdekis et al. (2003) discuss the behavior of these features for SSPs as predicted by recent evolutionary synthesis models. In the current analysis we also explore the dust sensitivity of the three classical Ca II triplet indices, as defined in DTT, and the generic indices of Cenarro et al. (2001a). 3.2.5 The Rose Indices Rose (1984, 1985) developed a complimentary set of optical indices, in part to overcome the difficulties in identifying the continuum in crowded spectral regions. The Rose indices are a measurement of the ratio of a given feature's central line intensity to that of a close reference line, without recourse to the (pseudo-) continuum level Information about equivalent widths is however lost in such relative measurements, and the Rose indices are not ideal for separating age and metallicity effects. Nonetheless, some of the Rose indices are quite sensitive to metallicity (see Vazdekis 1999) and, in combination with a sensitive age indicator (e.g. H/3), could help disentangle the two effects. The Rose indices also provide a unique and sensitive test for the presence of early-type stars (i.e. a post-starburst, and potentially dusty, population) in an integrated spectrum as well as a probe of the relative contribution of dwarf and evolved (red giant branch) stars to the integrated spectrum of a galaxy (Caldwell, Rose, & Concannon 2003). As Chapter 3. Dust Sensitivity of Absorption-Line Indices 75 with the Lick indices, it is usually assumed that the Rose indices are insensitive to dust reddening (e.g. Leonardi & Worthey 2000), an assumption we examine in §3.4.4. Finally, we also investigate dust effects on the pseudo-equivalent width, "Lick-style", indices of Rose (1994) and Jones & Worthey (1995) which have extremely narrow defi-nitions (and are denoted with the subscript "HR"). The H 7 H R index has been identified as one of the most sensitive age distriminators (Jones & Worthey 1995), but these narrow features suffer a worse resolution sensitivity than the W097 indices and thus can only be used with high quality data of low velocity-dispersion systems. Vazdekis & Arimoto (1999) and Vazdekis et al. (2001) confronted the velocity dispersion sensi-tivity of these narrow indices by defining a set of four H 7 indices that take resolution effects into account, allowing for reliable measurements in galaxies with velocity dis-persions up to o ~ 300 kms - 1 . Note, however, that the Vazdekis k, Arimoto (1999) indices require S/N > 200-400/A and have strong error covariances due to overlapping pseudo-continuum and central bandpasses, making them much harder to measure than the W097 indices. The response of these indices to dust extinction is similar to the W097 H 7 . F and the Jones k Worthey (1995) H 7 H R indices, so those results are not shown. 3.3 Models The stellar population synthesis models used for this analysis are those of Bruzual k Chariot (2003; hereafter GALAXEV) . The high resolution models using the Padova evolutionary tracks (Bertelli et al. 1994) and the Chabrier (2003) IMF were adopted. These models include SSP spectra in the wavelength range 3200 - 9500 A with a resolu-tion of 3 A, metallicities ranging from Z = 0.0001-0.05 (or 0.005-2.5 times Z 0 ) , and ages ranging from 0-20 Gyr (in 220 unequally-spaced time steps). For each time step, a number of integrated quantities are provided such as magnitudes and colours in many different filter systems, line-index strengths for all definitions in the Lick/IDS system, W097, the two definitions of the 4000 A break, and the DTT Ca II indices. The line strengths used here are computed directly from the-high-resolution model spectra, i.e. Chapter 3. Dust Sensitivity of Absorption-Line Indices 76 the spectra have not been transformed to the Lick/IDS system (see e.g. W097). The Rose, Ca II, and "HR" spectral indices, which are not provided with the BC03 model outputs, were computed with a public f ortran program by A. Vazdekis2. The G A L A X E V distribution allows for the computation of attenuation effects due to dust according to the two-component model of Chariot & Fall (2000; hereafter CF00). The two adjustable parameters of this model are fy, the total effective V-band optical depth affecting stars younger than 107 yr, and fi, the fraction of the total dust absorption contributed by diffuse interstellar medium (cirrus) dust. That is, where t is the age of any single stellar generation. This model has the singular feature, first introduced by Silva et al. (1998), of accounting for the finite lifetime (~ 107 yr) of stellar birth clouds3. The wavelength dependence of the effective absorption curve (proportional to A - 0 7 ) was constrained to reproduce the observed relation between the ratio of far-infrared to ultraviolet luminosities, and the ultraviolet spectral slope of their starburst galaxies (see also the GRASIL models of Silva et al. 1998). Note that t\ on the left-hand side of equation (3.5) is a function of time and fj,. As a result, even if // = 0 (no cirrus dust component), in the presence of young stellar populations (t < 107 yr), dust will still contribute some extinction (see Fig. 3.1). Note that in the adopted dust model of CF00, the model spectra are representative of the total integrated light of a galaxy, i.e., neither radial nor inclination dependent information is available. Other prescriptions for dust attenuation in galaxies exist (e.g. Witt, Thronson, & Capuano 1992; Byun, Freeman, & Kylafis 1994; de Jong 1996; Gordon et al. 2001), a few of which have been coupled with spectral synthesis and photo-ionization codes. For example, Moy et al. (2001) interfaced the PEGASE population synthesis models of Fioc & Rocca-Volmerange (1997) with the CLOUDY photo-ionization code of Ferland (2002). These models include a treatment for dust attenuation processes, but do not consider nebular emission. They also adopt a simple screen approximation for the dust 2See http://www.iac.es/galeria/vazdekis/ introduced to resolve the apparent conflict between the attenuation of line and continuum photons in their sample of starburst galaxies. { fy(A/5500 A)" 0- 7 for t < 107 yr, fXTV(A/5500 A) - 0 - 7 for t > 107 yr, (3.5) Chapter 3. Dust Sensitivity of Absorption-Line Indices 77 distribution, an unrealistic geometry for galaxies. Similarly, Panuzzo et al. (2003) cou-pled the same CLOUDY code with their spectrophotometric synthesis model GRASIL (Silva et al. 1998), to yield a complete treatment of dust reprocessing with more realistic geometries and account for the age dependence of molecular birth clouds (as in CF00). While these models may be more realistic and provide directional information that is lacking in the CF00 models, the resolution of the GRASIL output is set by the adopted input SSP library. Silva et al. (1998) used a resolution (20 A in the optical) that is insufficient for accurate line-index measurements, whereas Panuzzo et al. (2003), while considering nebular emission lines, coupled GRASIL with SSP libraries of sufficiently high resolution for the task at hand4. The added complexity of the GRASIL models is not warranted for the present study as we are mainly interested in order of magnitude effects from dust, in which case the CF00 models coupled to the BC03 SPS models provide a most suitable output. In a study of 705 non-Seyfert galaxies drawn from the Stromlo-APM redshift survey (Loveday et al. 1996), Chariot et al. (2002) determined parameters for the CF00 models that reproduce the observed integrated spectral properties of the nearby star-forming galaxies. These cover the ranges: 0.2 < Z/ZQ < 4.0, 0.01 < fv < 4.0, and 0.2 < ix < 1.0, for constant and exponential SFHs with ages 107 yr < t < 1010 yr. Two time-scales were adopted for the exponentially declining star formation rate: rexp = 0.1 and 6.0 Gyr. Our model realizations cover roughly the same parameter space. For the SSP models we consider values of \i = 0 and 1 only since, for the SSP ages shown it > 107 yr), different values of fi are equivalent to different values of TV. For the exponential SFHs, for which there are always populations with t < 107 yr present, we present models with 0 < \x < 1 and extend to fv = 8 (so for the /x = 0.3 case, the plots effectively extend to fv = 2.4). To express the effective dust reddening in a more familiar form, in Figure 3.1 we plot the colour excesses E(B — V) and E(V — K) as a function of f y resulting from the dust models used in this analysis. The upper panels 4 The corresponding SSP libraries with 4.5 A resolution in the optical are available upon request from the authors of Silva et al. (1998) for use with the G R A S I L models (L. Silva 2005, private commu-nication), but only the lower resolution SSPs are currently available for download directly from their U R L : h t t p ://web.p d . a s t r o . i t / g r a n a t o / g r a s i l / g r a s i l . h t m l . Chapter 3. Dust Sensitivity of Absorption-Line Indices 78 Figure 3.1: Colour excesses E(B — V) [left panels] and E(V — K) [right panels] as a function of fy for solar metallicity SSPs with fi = 0.0 (circles) & 1.0 (triangles) [top panels] and an exponential SFH with rexp = 13 Gyr for \i = 0.0 (circles), 0.3 (triangles), and 0.9 (squares) [bottom panels]. Different ages are represented by: 0.5 (dotted lines; red shades), 5 (dashed lines; green shades), & 13 Gyr (solid lines; blue shades). are for solar metallicity SSPs with fi = 0.0 (circles) & 1.0 (triangles) and the lower panels show exponential SFH with rexp = 13 Gyr for p = 0.0 (circles), 0.3 (triangles), 6 0.9 (squares) at ages of 0.5 (dotted lines; red shades), 5 (dashed lines; green shades), & 13 Gyr (solid lines; blue shades). Before examining the response of the individual indices to the dust reddening, it is useful to look at the spectral energy distribution (SED) of the models. Figure 3.2 shows the model SEDs in the 3700-6700 A range for SSPs [left panels] and exponential SFHs with rexp = 13 Gyr [right panels]. Al l spectra have been normalized to 5500 A. The SSP SEDs look entirely as expected, with the 0.5 Gyr population [top right panel] closely resembling an early-type (A - F) stellar spectrum with a blue continuum, strong Chapter 3. Dust Sensitivity of Absorption-Line Indices 79 4000 4500 5000 5500 6000 6500 4000 4500 5000 5500 6000 6500 MA) Figure 3.2: Comparison of solar metallicity SEDs with fv = 0 (black), 3 (gray), & 6 (pink) at ages of 0.5 [top panels] & 13 Gyr [bottom panels] for SSP models with p = 1.0 [left panels] and exponential SFH models with rexp = 13 Gyr and /i = 0.3 [right panels]. The spectra are all normalized to their flux at 5500 A. Balmer features, but few metallic features. The 13 Gyr spectrum [bottom right panel] more closely resembles a later-type ( G - K ) stellar spectrum with a redder continuum, weaker Balmer-line strengths, and more metallic absorption features. In both cases, the dust extincted profiles exhibit the overall low frequency reddening of the SED due to the A - 0 7 dependence of the dust model. On the other hand, the unreddened 0.5 Gyr SED for an exponential SFH with rexp = 13 Gyr (Fig. 3.2 [top right panel, black curve]) is even bluer than its SSP counterpart because of the presence of young stars (< 0.5 Gyr) from the ongoing SF. The contrast between the SSP and exponential SFHs is even Chapter 3. Dust Sensitivity of Absorption-Line Indices 80 more striking at 13 Gyr. The exponential SFH SED (Fig. 3.2 [bottom right panel, black curve]) is much bluer than its SSP counterpart [bottom left panel] due to the ongoing SF, but it also shows metallic features due to the presence of older stars (< 13 Gyr). In the exponential SFH case, the dust reddening is more complicated due to the presence of young stars that still live in their birth clouds (< 107 yr). These stars are more extincted than the older stars, the cirrus extinction component, when \x < 1.0, but being intrinsically bright, still contribute significantly to the total flux. In Figure 3.2 [right panels] we show the case for [i — 0.3. Clearly, any measurement made over a long baseline (e.g. colours) will be affected by the dust reddening. However, gauging its effects on absorption-line indices from a cursory examination of the SEDs is not a straightforward task. In the next section we examine individual features and compare the dust-free indices with those computed with dust. 3.4 Results 3.4.1 Dust Sensitivity of the Lick Indices The response of the Lick indices as a function of fy, A*, and age for solar metallicity (ZQ = 0.02) and abundance ratio ([a/Fe] — 0) SSP models is shown in Figure 3.3. In each panel, we plot Aindex versus fy, where Aindex is the difference between the indices measured with and without dust (for the same age & SFH), i.e. Aindex = index(fv) — index(fv = 0). Results are shown for model ages of 0.5 (light gray dotted lines; red shades in electronic edition), 5 (medium gray dashed lines; green shades in electronic edition), & 13 Gyr (black solid lines; blue shades in electronic edition). The two values for p of 0.0 & 1.0 are denoted by circles and triangles, respectively. The horizontal dotted lines in Figures 3.3, 3.4, & 3.7-3.11 denote typical measurement errors for the different indices; ~ 0 . l A for the atomic Lick indices, 0.01 mag for the molecular Lick indices, 0.4 A for the W097 indices, 0.03 for D(4000) and D„(4000), and ~ 0.02-0.05 for the Rose indices (Gorgas et al. 1999; Jones & Worthey 1995; Kauffmann et al. 2003a; Falcon-Barroso et al. 2003; Caldwell, Rose, &; Concannon 2003; J.J. Gonzalez 2004, private comm.), and serve as a guide for the magnitude of the effect Chapter 3. Dust Sensitivity of Absorption-Line Indices 81 0.015 h g 0.01 <0.005 1 i 4& • i &<••-A. , 1 • -A- - -A-1 | • .* • i 1 1 1 1 1 1 J£s --X. " - A -. 1 1 ' 1 ' 1 . I • 1 • I • •m * ' T " T 1 1 ' V ' V Figure 3.3: Lick and W097 index differences, Aindex = index(fv) — index(fv = 0), as a function of fy for solar metallicity SSPs. Different ages are represented by: 0.5 (dotted lines; red shades), 5 (dashed lines; green shades), & 13 Gyr (solid lines; blue shades). Different values of p, are represented as 0.0 (circles), and 1.0 (triangles). The black horizontal dotted lines represent typical measurement errors for the different indices. Chapter 3. Dust Sensitivity of Absorption-Line Indices 82 from dust required to be detectable above the noise. All model realizations presented here have solar metallicity, ZQ = 0.02. We have also considered the other metallicities in the G A L A X E V models. Generally, the dust effects are slightly more conspicuous at Z = 0.05, but diminish with decreasing metallicity of the stellar population. In most cases, the index measurements do not deviate considerably from their dust-free values in the SSP models, as expected. Since we consider SSPs at ages of 0.5 Gyr and greater, by definition, models with /j, = 0 have no dust extinction (thus circle symbols lie at 0 for all values of fv in Fig. 3.3). For fi = 1 models, deviations from the dust-free index measurements do occur. There are a few cases of note that show significant sensitivity to the dust extinction, such as the molecular indices Mgi & Mg 2 , and CNi & C N 2 . This is not surprising since these indices have the longest wavelength baseline (AA~400 & 200 A respectively). As evidenced by the molecular Mg indices, index measurements in the presence of dust can be non-linear and age dependent, thus a simple "dust-correction" poses a formidable challenge. The effect of extinction on the Mg indices goes in the opposite direction for young ages as it does for older ages (compare dotted and dashed/solid lines in Fig. 3.3). Also, the CN indices are more severely affected at younger ages whereas the Mg indices are more affected at young ages for small fv < 3, but at older ages for larger f v > 4. Of the Lick atomic indices, those showing the greatest sensitivity to dust are Fe5015 & Fe4668, both having the longest wavelength baselines of all atomic indices, but their deviations from the dust-free case are always smaller than the typical measurement errors. In addition to the standard Lick indices, we also show the combination indices, (Fe) = (Fe5270 + Fe5335)/2 and [MgFe], first introduced by Gonzalez (1993). For the latter we use the slightly modified definition of Thomas, Maraston, Sz Bender (2003), [MgFe]' = [Mg6 (0.72*Fe5270 + 0.28*Fe5335)]1/2, which has the advantage of being com-pletely independent of the element abundance ratio ([a/Fe]). These combination indices vary very little with fv for the SSP models (Fig. 3.3). The higher-order Balmer indices, shown in Figure 3.3, also show some sensitivity to dust at older ages in the SSP models, particularly the H.'JA index. Measurement errors Chapter 3. Dust Sensitivity of Absorption-Line Indices 83 for these indices are quite large compared to the other Lick indices (due to their narrow baselines), thus any reddening effect would likely not be detectable. Unlike GCs and elliptical galaxies, the assumption of SSP-like star formation is clearly not applicable for spiral galaxies. While the true SFHs of the latter are poorly known in detail, they are often approximated with an exponentially declining SFR, parameterized by a star formation timescale rexp (e.g. Bell & de Jong 2000; MacArthur et al. 2004; §2.4). In Figure 3.4 we show the index responses to dust for an exponential SFH with Texp = 13 Gyr, which is reasonably close to a constant SF rate (see Fig. 2.6 [top panel] between red Texp = 5 Gyr and magenta rexp — 100 Gyr curves) so there is a significant amount of "current" SF at each of the three epochs considered. Models with current SF yield considerably different results than the SSPs, as seen in Figure 3.4. Almost all of the Lick indices show sensitivity to dust, some more so at young ages (e.g. CNi , C N 2 , H/3, and all four W097 indices), others more so at older ages (e.g. Ca4227, Fe4531, Fe4668, Mg6, and NaD), and some show differences of equal magnitude, but in opposite directions at old and young ages (e.g. G4300, Fe4383, and the molecular Mg indices). The situation is undoubtedly complex and can be attributed to the fact that, with current SF occurring at all epochs, there is a significant and non-linear amount of extinction of the hottest, youngest stars which contribute considerably to the total optical flux. As such, dust effects are often stronger for models with smaller p and these effects saturate at optical depths above about fv = 3 (at which the birth clouds become completely optically thick). The situation for mixed stellar populations with significant amounts of current SF is clearly complicated and not easily prescribed. However, for many of the indices (e.g. Ca4227, Ca4455, Fe5406, Fe5709, Fe5782, NaD, TiOi, and Ti0 2 ) , the magnitude of the dust effects is never close to that of current typical measurement errors, and thus would not be detected. Unfortunately, however, the most often used indices (e.g. H/3, Mgb, the composite indices (Fe) and [MgFe]', and the W097 indices) are also the most seriously affected. It is useful to take a closer look at the cause for the erratic behavior in some of Chapter 3. Dust Sensitivity of Absorption-Line Indices 84 Figure 3.4: Same as Figure 3.3 but for for an exponential SFH with rexp = 13 Gyr and values of p are represented as 0.0 (circles), 0.3 (triangles), and 0.9 (squares). Chapter 3. Dust Sensitivity of Absorption-Line Indices 85 the indices in response to the dust extinction. In Figure 3.5 we focus on the H/3 index. Plotted in each panel are solar metallicity SEDs for fv = 0 (black), 3 (gray), & 6 (pink) at ages of 0.5 [top panels], & 13 Gyr [bottom panels]. The index red and blue pseudo-continuum band limits are marked by the green dashed vertical lines and the central band by the solid vertical lines. The dotted lines mark the pseudo-continuum level for each spectrum. The spectra are all normalized to their 4861 A flux (H/3 line center). Figure 3.5 plots SSPs [right panels] with \x = 1.0 exponential SFHs [left panels] with LI = 0.3. There is very little effect in the SSP index measurement, the change in slope of the SED is matched by the change in slope of the pseudo-continuum. How-ever, for the exponential SFH at 0.5 Gyr (Fig. 3.5 [top right panel]), the continuum is increased relative to the depth of the absorption-line in the dust extincted models. This can be understood by recognizing that the young OB stars that suffer significant extinction have lower H/3 values. Thus, by hiding these low H/3 index stars, the index is effectively increased. In fact, this is the case for all of the Balmer lines, hence all of the Balmer-line indices show similar behavior. This also highlights a major limitation when using Balmer-line indices in determining young ages (< 0.5 Gyr). The models become degenerate in age here because there are two possible ages that can be inferred at these Balmer index values, i.e. the Balmer indices are double-valued in the region of young stellar population ages. This point and its resulting limitations are further emphasized in §3.5. For the molecular indices, which have much broader baselines and crowding of spec-tral lines in their passband regions, the situation can be more complicated. For instance, in Figure 3.6 we focus on the Mgi index. The colour codes, line-types, and SFHs are the same as in Figure 3.5. In this case, given the much broader baseline, the slope of the pseudo-continuum in the extincted models changes quite dramatically. In the 0.5 Gyr SSP (Fig. 3.6 [top left panel]) and 13 Gyr exponential SFH (Fig. 3.6 [bottom right panel]) cases it changes from being a negative slope in the dust-free model (black) to positive in the most extincted curve (pink). These pseudo-continua pass through a crowded set of absorption lines in the central passband, and a simple inspection of Chapter 3. Dust Sensitivity of Absorption-Line Indices 86 I • • I • I • • I • • .1 I • .1 I r\ i i I i l l i I i i i l l i 11 -I 4 8 4 0 4 8 6 0 4880 4 8 4 0 4860 4 8 8 0 x(A) Figure 3.5: Comparison of the H/3 index for solar metallicity SSP models with p = 1.0 [left panels], and exponential SFH rexp = 13 Gyr models with / i = 0.3. [left panels]. All panels show models with f v — 0 (black), 3 (gray), & 6 (pink) at ages of 0.5 [top panels] & 13 Gyr [bottom panels]. The index red and blue pseudo-continuum band limits are marked by the green dashed vertical lines and the central band by the green solid vertical lines. The dotted lines mark the pseudo-continuum level for each spectrum. The spectra are all normalized to their flux at 4861A (the center of the H/3 line). Chapter 3. Dust Sensitivity of Absorption-Line Indices 87 4900 5000 5100 5200 5300 4900 5000 5100 5200 5300 X(A) Figure 3.6: Same as Figure 3.5 but for Mgx. The spectra are all normalized to their flux at 5107 A (roughly the center of the index central passband). these figures does not reveal which model would yield larger index measurements. It is clear from these figures that a simple dust correction prescription for all of the indices cannot be achieved. 3.4.2 Dust sensitivity of the 4000 A Break Figure 3.7 shows the response of D(4000) [top panels] and Dn(4000) [bottom panels] to attenuation from dust for SSPs [left panels] and exponential SFHs with rexp = 13 Gyr [right panels]. The narrower index is certainly less affected by dust (note the different Chapter 3. Dust Sensitivity of Absorption-Line Indices 88 Figure 3.7: D(4000) [top panels] and D„(4000) [bottom panels] differences, Aindex = index{jv)—index(jy — 0), as a function of fy for SSPs [left panels] and for an exponential SFH with rexp = 13 Gyr [right panels]. Different ages are represented by: 0.5 (dotted lines; red shades), 5 (dashed lines; green shades), & 13 Gyr (solid lines; blue shades). The horizontal dotted lines represent the typical measurement error. y-axis scales), but it is by no means impervious to dust reddening effects, and can reach deviations with magnitudes much larger than the typical error for this index. Dust effects are smaller at younger ages, but they can be significant at all ages for realistic values of fy. A detailed discussion of the influence of these dust effects on the determination of ages and metallicities is deferred to §3.5. Gorgas et al. (1999) suggest that D(4000) could be corrected for the effects of internal reddening according to, D(4000) c o r r e c t e d = D(4000) o b s e r v e d x M r " " " * " ^ , (3.6) Chapter 3. Dust Sensitivity of Absorption-Line Indices 89 derived using the mean extinction curve from Savage & Mathis (1979). However, this assumes that the dust is distributed as a screen, which is likely unrealistic for galaxies (e.g. Witt, Thronson, & Capuano 1992), and that the value of the colour excess E(B—V) is known. Even knowing E(B — V), as we do for these models, the above formula does not reproduce the dust-free D(4000) value for the CFOO dust model. The numerical constant in equation (3.6) could be adjusted to obtain a good fit, but this number is a function of age, depends on the dust properties, and requires a priori knowledge of E(B — V). Hence, we do not explore this putative correction further, but note that, as we concluded for the Lick indices, there is likely no such simple dust correction for the 4000 A break of mixed stellar populations and complex dust geometries. 3.4.3 Dust Sensitivity of the Ca II Triplet Indices The response of the Ca II triplet indices of DTT and Cenarro et al. (2001a) as a function of fy, p, and age for solar metallicity SSP models are shown in Figure 3.8 and for an exponential SFH with rexp = 13 Gyr in Figure 3.9. The three classical Lick-style Ca indices of DTT (left) show some sensitivity to dust extinction for the SSPs, Ca3 showing the largest sensitivity, but the deviations remain below the typical measurement errors thus would not be detected. The generic CaT indices (right) are virtually unaffected by dust extinction in the SSP models at any age. As usual, the situation is not as straight-forward for the exponential SFH (Fig. 3.9). Here, the only index that is not significantly affected is the Cal index of DTT. All of the generic indices of Cenarro et al. (2001a) are significantly altered due to dust extinction, particularly at young ages. As in the case for the Lick indices in the exponential SFH models (§3.4.1, Fig. 3.4), dust effects are often stronger for models with smaller // and the effects saturate at optical depths above about fv = 4. In general, though, for models that do not have significant amounts of current SF, the dust reddening effects in the Ca II triplet indices are still small compared to measurement errors, and should not cause problems in their interpretation. Chapter 3. Dust Sensitivity of Absorption-Line Indices 90 Figure 3.8: Calcium triplet index differences, Aindex = index(fv) — index(fv = 0), as a function of fv for solar metallicity SSPs. Different ages are represented by: 0.5 (dotted lines; red shades), 5 (dashed lines; green shades), & 13 Gyr (black solid lines; blue shades). Different values of p are represented as 0.0 (circles), and 1.0 (triangles). The black horizontal dotted lines represent typical measurement errors for the different indices. 3.4.4 Dust Sensitivity of the Rose Indices The response of the Rose (1984, 1994) and Jones & Worthey (1995) indices are shown in Figure 3.10 for SSPs and in Figure 3.11 for an exponential SFH with rexp = 13 Gyr. Symbols and line types are as in Figures 3.3 & 3.4. For SSP models there are a few indices that are considerably modified by dust reddening, namely the H.8/Fe IA4045, E6/Fe IA4063, Sr II/Fe IA4045, H7/Gband, and Ca II indices, particularly at older ages. A l l other indices, including the pseudo-equivalent width indices, are essentially unaffected by the dust reddening in the SSPs at any age. For the exponential SFH, shown in Figure 3.11, the response of the indices gets quite complicated. The offset from the dust-free case can be either positive (typically for the older models) or negative (typically for youngest models). Again, for the p = 0.0 models and for some of the Chapter 3. Dust Sensitivity of Absorption-Line Indices 91 0 2 4 6 8 0 2 4 6 8 Figure 3.9: Same as Figure 3.8 but for for an exponential SFH with rexp = 13 Gyr. Different values of p are represented by 0.0 (circles), 0.3 (triangles), and 0.9 (squares). p = 0.3 models, the index offsets saturate for optical depths above fy = 3. Ultimately though, even with current SF, for the majority of the Rose indices, the index offsets due to dust reddening are modest. On the other hand, with a significant amount of current SF, the pseudo-equivalent width indices can be quite affected, more so at young ages for H 7 H R , but more so for old ages for Ca IHR and Fe I H R. The "HR" indices are virtually unaffected for SSP models at all ages. 3.5 Age and Metallicity Fitting In order to translate absorption-line indices to a physical age and metallicity scale, the index measurements must be compared with stellar population synthesis models. In their most basic form, commonly referred to as simple stellar populations (SSPs), these models provide evolutionary information for a coeval population of stars born with a given composition and initial mass function (IMF). Several such SSP models have been produced by a number of independent groups and are in a constant state of flux as im-Chapter 3. Dust Sensitivity of Absorption-Line Indices 92 Figure 3.10: Rose index differences, Aindex = index{fv) — index(fv = 0), as a function of TV for solar metallicity SSPs. Different ages are represented by: 0.5 (dotted lines; red shades), 5 (medium gray dashed lines; green shades in electronic edition), & 13 Gyr (solid lines; blue shades). Different values of \x are represented by 0.0 (circles), and 1.0 (triangles). The horizontal dotted lines represent typical measurement errors for the different indices. Chapter 3. Dust Sensitivity of Absorption-Line Indices 93 Figure 3.11: Same as Figure 3.10 but for an exponential SFH with rexp = 13 Gyr. Different values of fi are represented by 0.0 (circles), 0.3 (triangles), and 0.9 (squares). Chapter 3. Dust Sensitivity of Absorption-Line Indices 94 provements to many of the input parameters (e.g. stellar libraries, model atmospheres, convection, mass loss, mixing) come to light. There are discrepancies among the dif-ferent models that, depending on the application, may result in significantly different interpretations of the observed stellar population signatures. We are not interested here in a detailed comparison of the different SPS models, but rather seek a quantitative guide for the magnitude of the errors due to dust effects on the physical parameters derived from absorption-line indices. To this end, we use the 2003 implementation of the Bruzual & Chariot (2003) SPS models (GALAXEV) for the age and metallicity determinations, the same as are used to compute the model "galaxies" (see §3.3). While most galaxies have likely undergone SFHs that are much more complex than the simple SSP described above, the common practice for line-index measurements is to compute "SSP-ages and metallicities", that is, compare the galaxy indices to the SSP grids, irrespective of the true SFH of the galaxy (which is unknown). We adopt this technique here and thus refer to SSP parameters even for the models with exponential SFHs. Ages and metallicities are determined by fitting the model galaxy index measure-ments to the SSP model grids using a maximum-likelihood approach. We retain the same grid in time as is provided in the model SSPs (220 unequally spaced time steps from t = 0 - 20 Gyr) and interpolate (linearly) between the model metallicities on a fine grid of 120 unequally spaced steps in Z. In order to compare all indices on a similar scale, the atomic indices computed as equivalent widths in A, IEW, are converted to a magnitude, I M A G , as: I M A G = -2.5 log 1 0 ( l - ^ (3.7) where AA is the width in A of the index feature bandpass. The corresponding error in magnitude is: _ 2.5 1og10(e) , . a m a 9 ~ A A I O - 0 - 4 ^ E W ' { ' The two definitions of the 4000 A break, which are dimensionless flux ratios, are con-Chapter 3. Dust Sensitivity of Absorption-Line Indices 95 verted to a magnitude scale as: D(4000) mag 2.5 log 1 0 [D(4000)rat J (3.9) with the corresponding error in magnitude as: <r[D(4000) 1 = 2.5 log10(e) gp(4000) r a f t o ] D(4000) r a t.o (3.10) We compute an age and metallicity by minimizing the following figure of merit: where N is the number of indices (only 2 indices are used here, but this method can be generalized to fit multiple indices simultaneously, see §4.5.3), Oi is the "observed" value (in magnitudes) of index i and <50; is its error in magnitudes, and M(A, Z\ is the SSP model value of index i for a given age and metallicity combination. Errors for the individual age and metallicity measurements are estimated using a Monte Carlo approach. One thousand realizations of the model fits are computed using errors drawn from a normal distribution of the observational errors (taken here as the "typical" observation errors for each index quoted in §3.2). The errors for the measured ages and metallicities are taken as half the interval containing 68% of the 1000 Monte Carlo realizations (i.e. the la confidence interval). Two examples of the fits are given in Figure 3.12 which shows the logarithmic difference between the physical parameters derived from each simulated point relative to the best-fit value for the Dn(4000) -Fe4668 plane (see also the corresponding index-index plot in Figure 3.13 [right panel}). The left panel presents the fit for the dust-free (fi = 0.0) SSP case which indeed finds the correct model age (A = 5 Gyr; Alog 1 0[Age (Gyr)] = 0) and metallicity (Z = 0.02; A[\og10(Z/ZQ)\ = 0), with corresponding la confidence intervals (represented as error bars) of ~0.2 dex for both parameters. The right panel presents the fit for the dusty case (LI = 1.0) at 0.5 Gyr (but note the different y-axis scales). The TV = 0. fit finds the correct physical parameters, but as the dust opacity increases, the best fit wanders around in age (from A = 0.5 Gyr at f v = 0 to A = 10.0 Gyr at fy = 6) and metallicity Chapter 3. Dust Sensitivity of Absorption-Line Indices 96 -0.4 -0.3 0 0.3 0.4 -1.4 -1.2 -1 -0.8 -0.6 -0.4 -0.3 0 0.3 0.4 A l o g 1 0 ( Z / Z 0 ) A l o g 1 0 ( Z / Z e ) Figure 3.12: Examples of the Monte Carlo fits for D n(4000) versus Fe4668 diagnostic plots, represented as the logarithmic difference between the physical parameter derived from each simulated point relative to the best-fit value (also see the corresponding index-index plot in Figure 3.13, right panel). The error bars represent the lo confidence limits. Left: SSP model at 5 Gyr with LI = 0.0. Right: SSP model at 0.5 Gyr with /j, = 1.0 (note the different y-axis scales). (from Z = 0.02 at f v = 0 to Z = 0.003 at f v = 6). Since we do not consider extrapolations, there are "edges" in these plots that represent the limits of the model grids. As a result, the confidence limits can be underestimated for points fit near the model grid limits, which also generally correspond to a region of greater degeneracy in age and metallicity. As a result, any points that lie near the model limits should be treated with caution. Given that we are comparing model galaxies with exponential SFHs against SSP grids, it is inevitable that some of the parameter space will fall outside the model grids and thus cannot be fit reliably. In fact, even for some of the SSP models, the dust extinction can be severe enough to push the locus of the model point in the index-index diagram off the SSP model grids. Both of these situations are observed in the Dn(4000) versus Fe4668 index-index diagram in Figure 3.13 [right panel]. Tables 3.5-3.5 present the results from our age and metallicity fitting for a selected Chapter 3. Dust Sensitivity of Absorption-Line Indices 97 Figure 3.13: H/3 versus Mgb [left] and Dn(4000) versus Fe4668 [right] diagnostic plots. The grids are dust-free BC03 SSP models. Lines of constant age (dashed) are shown for ages 0.2, 0.5, 1.0, 2.0, 5.0, 10.0, and 20.0 Gyr in both panels, but in the H/3 versus Mgb plot [left], the age extends to 0.002 Gyr to illustrate the messy young age, double-valued Balmer line, region (see text). Lines of constant metallicity (solid) are shown for Z = 0.0001 (purple), 0.0004 (dark blue), 0.004 (light blue), 0.008 (green), 0.02 (yellow), and 0.05 (red). The circles are the indices measured from the solar metallicity SSP models with \i = 1.0 for f v = 0 (largest point size) to f y = 6 (smallest point size), and the triangles are the exponential SFH model indices with rexp = 13 Gyr and \x = 0.3 for f y = 0 (largest point size) to f y = 8 (smallest point size). The red, green, and blue shades are for model ages of 0.5, 5, and 13 Gyr respectively. number of index-index combinations: H/3 versus Mgb, Mg 2 , and [MgFe]' (top) and H/3, K6A, and Dn(4000) versus Fe4668 (bottom) for the SSP and exponential SFH model galaxies respectively (tables for other index combinations are available from the author upon request). See Figures 3.13, 3.15, & 3.16 for the corresponding index-index plots. For each index-pair, we list the model age, MA, the effective V-band optical depth, f y , whether the fit is reliable in the "fit" column (see below), the best-fit age with its la-error in brackets, the difference between the measured age in the extincted models and the dust-free age, A A = A g e ( f y ) - A g e ( f y = 0), and the same for the metallicities, which are given as log l o (Z/Z 0 ) and AZ = l o g l o ( Z / Z 0 ) ( f y ) - \ogw(Z/Ze)(fv = 0). Chapter 3. Dust Sensitivity of Absorption-Line Indices 98 Reliable fits are indicated with a check mark, / . For unreliable fits the check mark is replaced with a code indicating the reason for which they could not be fit. The codes are: "OG" if the point lies off the model grids, "DB" if it lies in the young double-valued Balmer line region, and "DM" if it lies in a degenerate metallicity region (which only occurred for the young exponential SFH model in the Dn(4000)-Fe4668 plane). Note that, in the absence of a reliable fit for the dust-free case, even if reliable fits exist for the corresponding dusty models, the A's are not well defined (this only occurred for the 0.5 Gyr exponential SFH model in the Dn(4000)-Fe4668 plane). Examples of the messy double-valued Balmer line region can be seen in the H/3 versus Mg6 plot in Figure 3.13 [left panel] where the ages plotted extend down to 0.002 Gyr, as well as in Figures 3.17 & 3.18 in §3.6. It is obvious from Tables 3.5-3.5 (and see Fig. 3.12) that there is significant uncer-tainty in the fits given our assumed "typical" measurement errors, regardless of the dust content. For example, the 0.5 Gyr SSP model la confidence limits (Tables 3.5-3.5) range from ~ 0.1-0.4 Gyr in age and ~ 0.2 -0.3 dex in metallicity, depending on the index-index grid used. Note that the la confidence limits are not only dependent on the particular index-index plane used, but they also depend on the location of the point within the grid (in other words, to what degree the age and metallicity are separated at the given locus). In general, the degeneracies become more severe at the extreme metallicities (i.e. Z < 0.001 and Z > 0.02), and at young ages, where both the narrower region covered by the metallicity span and the double-valued Balmer lines contribute to the problem. As expected from §3.4, the fits for the 0.5 Gyr SSP models (Tables 3.5-3.5) are virtually unaffected by dust (for the index-index planes shown), with the exception of the Dn(4000)-Fe4668 plane, for which the age and metallicity errors are large (up to 9.5 Gyr and 0.8 dex respectively) and significant compared to the la confidence limits, even for small amounts of dust. Significant errors occur at the highest dust opacities for the 5 Gyr models (in all but the H/3-Mg6 and [MgFe]' planes), but in this case they are of the same order as the la confidence limits. The errors due to dust on age Chapter 3. Dust Sensitivity of Absorption-Line Indices 99 and metallicity are most severe in the 13 Gyr models, with some points ending up off the model grids at the highest fy's (e.g. in the H/3 and H ^ versus Fe4668 grids), but these are again matched by the large la confidence limits for fits in these regions of the grid. In summary, for the SSP model galaxies, if the model indices lie in a region of the index-index diagnostic plot where the age and metallicity are well separated (close to orthogonal), the fits are generally independant of dust extinction for most index combinations. Where the fits are affected by the dust (generally at older ages and for indices with broader baselines), the errors in the derived ages and metallicities are generally of the same order as the la confidence limits and would thus be difficult to detect. The most notable exceptions here are the H/3-Mg 2 plane, where the dust pushes the points for the 0.5 Gyr models into the degenerate double-valued Balmer line region, and the Dn(4000)-Fe4668 plane, where the errors due to dust become large at young ages, and at old ages the points extend off the edges of the model grids. On the other hand, for the exponential SFH models, again as expected from §3.4, the situation is more complicated. Examination of Tables 3.5-3.5 reveals that none of the 0.5 Gyr models can be fit reliably in the Balmer-index versus metallicity-index grids as they all fall in the double-valued Balmer line region of the plots. Reliable fits for the dustier 0.5 Gyr models in the D„(4000)-Fe4668 plane are obtained, but here the A's are undefined due to the poor fit in the dust-free model. Reliable fits are found for most of the 5 Gyr exponential SFH models with "SSP ages" in the range ~ 0.9 -1.0 Gyr (i.e. this is the luminosity-weighted SSP age of a 5 Gyr old stellar population that has been forming stars at a constant rate, roughly, throughout its lifetime). The ages, when reliably fit, generally agree between the different index-index plane fits. When dust has an effect on the ages, it tends to make them slightly youger, but the AA's are always of the same order as the lcr confidence limits for the fits. The metallicity fits for these models, which are all inherently solar metallicity, are in the range logw(Z/ZQ) ~ —0.7 to +0.2 (or Z = 0.004-0.03). They are generally lower(higher) than solar in the 5 (13) Gyr models and increase with dust extinction. Unlike the age fits, the metallicity fits do not agree very well between the different index-index planes. For example, the Chapter 3. Dust Sensitivity of Absorption-Line Indices 100 fv = 0 fit in 13 Gyr models gives log l o (Z/Z 0 ) = +0.32 and -0.12 dex in the H/3-Mgfc and H/3-Fe4668 planes respectively. The ages for the 13 Gyr exponential SFH model fits are also in the ~ 0.9-1.0 Gyr range. While these do have weaker Balmer lines than the 5 Gyr models, their metallicity index values are slightly stronger. This has the overall effect of moving them along a line of roughly constant age. An example of this is shown by the 5 Gyr (green) and 13 Gyr (blue) triangles in the H/3 versus Mgb diagnostic plot in Figure 3.13 [left panel]. These results highlight once again the greater challenge of fitting reliable ages and metallicities for "young" populations. Finally, in Figure 3.14 we demonstrate that the dust effects on the (Fe) versus Mgb plot tends to move the points along the degeneracy, and therefor could not mimic (or hide) abundance ratio enhancements of the kind often observed in elliptical galaxies (e.g. Gonzalez 1993). 4.5 4 3.5 3 A CD 2.5 fa V 2 1.5 1 0.5 0.05 0 0.5 1 1.5 2 2.5 3 3.5 4 4.5 5 5.5 Mgb Figure 3.14: (Fe) versus Mgb diagnostic plot. The lines and symbols are as in Figure 3.13. Table 3.5: Age & Metallicity Fits and Errors as a function of fv: H/3 vs. Mgb and H/3 vs. M g 2 Model Parameters H/3 vs. Mg6 H/3 vs. Mg 2 SFH Age TV fitt Age A A logio(Z/Z 0) AZ fitt Age A A logi O (Z/Z 0 ) AZ SSP 05 0 / 0.51 0.22) 0.00 0.00 ( 0.35) 0.00 D B 0.51 0.43) 0.00 0.00 ( 0.38) 0.00 H = 1.0 2 / 0.51 0.22) 0.00 0.00 ( 0.36) 0.00 D B 0.45 0.36) -0.06 -0.44 ( 0.74) -0.44 4 / 0.51 0.22) 0.00 0.00 ( 0.36) 0.00 D B 0.45 0.36) -0.06 -0.46 ( 0.77) -0.46 6 / 0.51 0.22) 0.00 0.00 ( 0.35) 0.00 D B 0.09 0.19) -0.42 0.29 ( 0.38) 0.29 5 0 / 5.00 1.63) 0.00 0.00 ( 0.13) 0.00 / 5.00 1.50) 0.00 0.00 ( 0.09) 0.00 2 / 5.00 1.50) 0.00 0.00 ( 0.12) 0.00 / 5.00 1.50) 0.00 0.00 ( 0.09) 0.00 4 / 5.00 1.50) 0.00 0.00 ( 0.12) 0.00 / 4.50 1.63) -0.50 0.06 ( 0.13) 0.06 6 / 5.00 1.50) 0.00 0.00 ( 0.12) 0.00 / 4.25 1.57) -0.75 0.11 ( 0.18) 0.11 13 0 / 13.00 6.25) 0.00 0.00 ( 0.21) 0.00 / 13.00 4.38) 0.00 0.00 ( 0.13) 0.00 2 / 11.50 6.00) -1.50 0.08 ( 0.20) 0.08 / 13.25 4.38) 0.25 0.00 ( 0.13) 0.00 4 / 10.50 6.25) -2.50 0.13 ( 0.21) 0.13 / 11.50 4.75) -1.50 0.08 ( 0.13) 0.08 6 / 10.50 5.87) -2.50 0.13 ( 0.20) 0.13 / 9.50 4.75) -3.50 0.19 ( 0.15) 0.19 EXP 05 0 D B 0.08 ; 0.02) 0.00 -0.49 ( 0.50) 0.00 D B 0.06 0.03) 0.00 -0.10 ( 0.52) 0.00 H = 0.3 2 D B 0.11 ' 0.04) 0.03 -0.60 ( 0.21) -0.11 D B 0.64 0.60) 0.58 -0.60 ( 0.53) -0.51 rexp = 13 Gyr 4 D B 0.14 k 0.26) 0.06 -0.62 ( 0.21) -0.12 D B 0.08 0.27) 0.02 0.02 ( 0.53) 0.12 6 D B 0.14 k 0.26) 0.06 -0.60 ( 0.22) -0.11 D B 0.03 0.54) -0.03 0.40 ( 1.08) 0.49 8 D B 0.14 ; 0.22) 0.06 -0.60 ( 0.19) -0.11 D B 0.09 k 0.26) 0.03 -0.05 ( 0.51) 0.05 5 0 / 1.02 0.00 -0.18 ( 0.45) 0.00 / 0.90 , 0.10) 0.00 0.10 ( 0.19) 0.00 2 / 1.02 [ 0.21) 0.00 -0.68 ( 0.90) -0.50 / 0.81 0.10) -0.10 0.22 ( ° - 1 2 ) 0.12 4 / 0.81 ; 0.21) -0.21 0.20 ( 0.41) 0.38 / 0.81 k 0.10) -0.10 0.20 ( 0.12) 0.11 6 / 0.81 I °-21) -0.21 0.20 ( 0.36) 0.38 / 0.81 [ 0.09) -0.10 0.20 ( 0.12) 0.11 8 / 0.81 ; 0.21) -0.21 0.20 ( 0.34) 0.38 / 0.81 ; o.io) -0.10 0.20 ( 0.12) 0.11 13 0 / 0.90 I o-ii) 0.00 0.32 ( 0.07) 0.00 / 0.90 0.00 0.33 ( 0.08) 0.00 2 / 0.90 I 0.10) 0.00 0.33 ( 0.04) 0.01 / 0.90 I 0-10) 0.00 0.33 ( 0.04) 0.00 4 / 0.90 [ 0.10) 0.00 0.33 ( 0.04) 0.01 / 0.90 I 0-10) 0.00 0.33 ( 0.04) 0.00 6 / 0.90 ( 0.10) 0.00 0.33 ( 0.04) 0.01 / 0.90 I °-10) 0.00 0.33 ( 0.04) 0.00 8 / 0.90 ( 0.10) 0.00 0.33 ( 0.04) 0.01 / 0.90 ; o.io) 0.00 0.33 ( 0.04) 0.00 Note. - All ages are in Gyr. The numbers in brackets are la confidence limits based on Monte Carlo sampling of the errors, t A check mark, / , indicates a reliable fit. Unreliable fits are coded as: " O G ' if the point lies off the model grids, " D B ' if it lies in the young double-valued Balmer line region, and " D M ' if it lies in a degenerate metallicity region. Table 3.6: Age & Metallicity Fits and Errors as a function of fv: H/3 vs. [MgFe]' and H/3 vs. (Fe) M o d e l Parameters H/? vs. [MgFe]' H/3 vs. (Fe) S F H Age fv fitt A g e A A log10(Z/Ze) AZ fitt Age A A logl0(Z/Ze) AZ S S P 0.5 0 / 0.51 0.10) 0.00 0.00 ( 0.33) 0.00 / 0.51 0.10) 0.00 0.00 ( 0.25) 0.00 /!= 1.0 2 / 0.51 0.10) 0.00 -0.01 ( 0.33) -0.01 / 0.51 0.10) 0.00 0.00 ( 0.22) 0.00 4 / 0.51 0.10) 0.00 -0.01 ( 0.32) -0.01 / 0.51 0.10) 0.00 0.00 ( 0.24) 0.00 6 / 0.51 0.10) 0.00 0.00 ( 0.34) 0.00 / 0.51 ; o.io) 0.00 0.00 ( 0.24) 0.00 5 0 / 5.00 1.63) 0.00 0.00 ( 0.13) 0.00 / 5.00 ; i.50) 0.00 0.00 ( 0.14) 0.00 2 / 5.00 1.50) 0.00 0.00 ( 0-12) 0.00 / 5.00 ; i.50) 0.00 0.00 ( ° - 1 4 ) 0.00 4 / 5.00 1.75) 0.00 0.00 ( 0-14) 0.00 / 4.75 ; i.63) -0.25 0.04 ( 0.15) 0.04 6 / 4.75 1.75) -0.25 0.04 ( 0-14) 0.04 / 5.00 0.00 0.02 ( 0.15) 0.02 13 0 / 13.00 5.88) 0.00 0.00 ( 0.22) 0.00 / 13.00 ; 5.88) 0.00 0.00 ( 0.22) 0.00 2 / 12.25 6.00) -0.75 0.04 ( 0.22) 0.04 / 12.25 I 5.75) -0.75 0.04 ( 0.23) 0.04 4 / 11.25 6.00) -1.75 0.10 ( 0.22) 0.10 / 5.50 ; 6.i3) -7.50 0.34 ( 0.23) 0.34 6 / 10.00 6.00) -3.00 0.16 ( 0.22) 0.16 / 5.25 ; 6.13) -7.75 0.36 ( 0.23) 0.36 - E X P 0.5 0 D B 0.06 0.33) 0.00 -0.01 ( 0.43) 0.00 D B 0.05 [ 0.34) 0.00 0.11 ( 0.41) 0.00 \i = 0.3 2 D B 0.64 0.54) 0.58 -0.82 ( 0.19) -0.82 D B 0.64 [ 0.58) 0.59 -0.59 ( 0.65) -0.70 T e x p = 13 G y r 4 D B 0.57 0.19) 0.51 -0.76 ( 0.10) -0.75 D B 0.64 ; 0.55) 0.59 -0.54 ( 0.28) -0.65 6 D B 0.57 0.19) 0.51 -0.76 ( 0.08) -0.75 D B 0.64 [ 0.55) 0.59 -0.54 ( 0.27) -0.65 8 D B 0.26 0.39) 0.20 -0.70 ( 0.09) -0.69 D B 0.64 [ 0.55) 0.59 -0.54 ( 0.28) -0.65 5 0 / 1.02 ; 0.00) 0.00 -0.22 ( 0.19) 0.00 / 1.02 { 0.00) 0.00 -0.27 (o.ii) 0.00 2 / 0.90 , 0.11) -0.11 0.04 ( 0.22) 0.26 / 1.02 ; o.ii) 0.00 -0.27 ( 0.13) 0.00 4 / 0.90 ' 0.10) -0.11 0.04 ( 0.22) 0.26 / 1.02 [ o-11) 0.00 -0.27 ( 0.15) 0.00 6 / 0.90 ; o.io) -0.11 0.04 ( 0.21) 0.26 / 1.02 { 0.11) 0.00 -0.25 ( ° - 1 4 ) 0.02 8 / 0.90 ; o.io) -0.11 0.04 ( 0.22) 0.26 / 1.02 I o.n) 0.00 -0.25 ( 0.15) 0.02 13 0 / 1.02 ; 0.12) 0.00 0.20 ( 0.18) 0.00 / 1.14 { 0.12) 0.00 -0.12 ( 0.15) 0.00 2 / 0.90 -0.11 0.32 ( 0.07) 0.12 / 1.02 ; 0.12) -0.12 0.18 ( 0.22) 0.30 4 / 0.90 ; o.io) -0.11 0.32 ( 0.07) 0.12 / 1.02 [ ° - 1 2 ) -0.12 0.18 ( 0.21) 0.30 6 / 0.90 ; o.io) -0.11 0.32 ( 0.06) 0.12 / 1.02 -0.12 0.18 ( 0.20) 0.30 8 / 0.90 [ 0.10) -0.11 0.32 ( 0.06) 0.12 / 1.02 [ 0.12) -0.12 0.18 ( 0.21) 0.30 Note. - A l l ages are in G y r . T h e numbers in brackets are la confidence l imits based on Monte C a r l o sampling of the errors, t A check mark , / , indicates a reliable fit. Unrel iable fits are coded as: " O G ' if the point lies off the model grids, " D B ' if it lies in the young double-valued B a l m e r line region, and " D M ' if it lies in a degenerate metall icity region. Table 3.7: Age & Metallicity Fits and Errors as a function of fv: H/3 vs. G4300 and H/3 vs. T i 0 2 Model Parameters H/3 vs. G4300 H/3 vs. TiC-2 S F H Age TV fitt Age A A l o g i 0 ( Z / Z o ) AZ fitt Age A A logio(£/Z©) AZ SSP 0.5 0 D M 0.51 ( 0.10) 0.00 0.00 0.34) 0.00 D B 0.51 0.33) 0.00 0.00 ; 0.46) 0.00 / i = 1.0 2 D M 0.45 ( 0.05) -0.06 -0.44 0.33) -0.44 D B 0.51 0.31) 0.00 -0.06 ; 0.44) -0.06 4 D M 0.45 0.11) -0.06 -0.44 0.30) -0.44 D B 0.51 0.33) 0.00 -0.06 ; 0.43) -0.06 6 D M 0.45 ( 0.20) -0.06 -0.46 0.28) -0.46 D B 0.51 0.31) 0.00 -0.06 ; 0.43) -0.06 5 0 D M 5.00 ( 0.87) 0.00 0.00 0.14) 0.00 / 5.00 1.00) 0.00 0.00 [ 0-04) 0.00 2 D M 5.00 0.87) 0.00 0.00 0.12) 0.00 / 5.00 1.25) 0.00 -0.01 ; 0.05) -0.01 4 D M 5.25 0.87) 0.25 0.00 0.10) 0.00 / 5.00 1.12) 0.00 0.00 ; 0.04) 0.00 6 D M 5.25 0.75) 0.25 0.00 0.08) 0.00 / 4.75 ' 1.12) -0.25 0.02 ; 0.04) 0.02 13 0 D M 13.00 4.12) 0.00 0.00 0.31) 0.00 / 13.00 2.88) 0.00 0.00 ; 0.04) 0.00 2 D M 13.50 4.37) 0.50 -0.07 0.29) -0.07 / 13.00 2.88) 0.00 0.00 ( 0.04) 0.00 4 D M 14.75 4.88) 1.75 -0.22 0.29) -0.22 / 13.25 3.00) 0.25 0.00 ; o.o4) 0.00 6 D M 16.25 4.63) 3.25 -0.36 0.26) -0.36 / 12.75 3.00) -0.25 0.02 ; o.o4) 0.02 E X P 0.5 0 D B 0.13 0.08) 0.00 -1.82 k 1.92) 0.00 D B 0.57 ; 0.51) 0.00 0.40 ; 0.49) 0.00 fj, = 0.3 2 D B 0.09 0.06) -0.04 -0.14 k 0.81) 1.68 D B 0.07 ' 0.02) -0.50 0.00 ( 0.04) -0.40 Texp = 13 Gyr 4 D B 0.18 0.07) 0.05 -1.35 ; 1.52) 0.48 D B 0.08 ; o.oi) -0.49 0.04 [ 0.05) -0.36 6 D B 0.18 0.08) 0.05 -1.35 I 1-51) 0.48 D B 0.08 ; 0.37) -0.49 0.02 ( 0.17) -0.38 8 D B 0.18 0.08) 0.05 -1.35 I 1-54) 0.48 D B 0.08 ; o.i9) -0.49 0.02 ( 0.15) -0.38 5 0 D B 1.14 0.12) 0.00 -1.22 { 0.25) 0.00 D B 0.04 ; o.oo) 0.00 -0.18 ; o.o4) 0.00 2 D B 1.02 0.00) -0.12 -0.81 ; cos) 0.41 D B 0.05 ; o.oo) 0.00 -0.17 ( 0.03) 0.01 4 D B 1.02 0.11) -0.12 -0.80 [ 0.07) 0.43 D B 0.05 ; o.oo) 0.00 -0.17 ( 0.02) 0.01 6 D B 1.02 ' 0.00) -0.12 -0.80 ; 0.07) 0.43 D B 0.05 ; o.oo) 0.00 -0.17 ( 0.03) 0.01 8 D B 1.02 0.11) -0.12 -0.81 ; 0.07) 0.41 D B 0.05 ; o.oo) 0.00 -0.17 [ 0.03) 0.01 13 0 D B 1.43 ; 0.27) 0.00 -0.87 [ 0.10) 0.00 D B 0.03 [ o-oi) 0.00 -0.24 ( 0.06) 0.00 2 D B 1.14 ; 0.29) -0.29 -0.70 ( 0.10) 0.17 D B 0.03 [ 0.01) 0.00 -0.22 ( 0.04) 0.01 4 D B 1.14 [ 0.29) -0.29 -0.70 ( 0.11) 0.17 D B 0.03 ; o.oi) 0.00 -0.22 ( 0.04) 0.01 6 D B 1.14 ; 0.29) -0.29 -0.70 ( 0.12) 0.17 D B 0.03 [ 0.01) 0.00 -0.24 ( 0.04) 0.00 8 D B 1.14 [ 0.29) -0.29 -0.70 ( 0.12) 0.17 D B 0.03 ( 0.01) 0.00 -0.24 ( 0.04) 0.00 Note. - A l l ages are in Gyr. The numbers in brackets are la confidence limits based on Monte Carlo sampling of the errors, t A check mark, / , indicates a reliable fit. Unreliable fits are coded as: " O G ' if the point lies off the model grids, " D B ' if it lies in the young double-valued Balmer line region, and " D M ' if it lies in a degenerate metallicity region. Table 3.8: Age & Metallicity Fits and Errors as a function of fv: H/3 vs. Fe4668 and Dn(4000) vs. Fe4668 Model Parameters H£ vs. Fe4668 D„(4000) vs. Fe4668 S F H Age fv fitt Ag e A A l o g i 0 ( Z / Z o ) AZ fitt Age A A log10(Z/ZQ) AZ SSP 0.5 0 / 0.51 ( 0.10) 0.00 0.00 0.31) 0.00 / 0.51 0.12) 0.00 0.00 0.34) 0.00 fi = 1.0 2 / 0.51 ( 0.10) 0.00 -0.03 0.34) -0.03 / 0.90 0.48) 0.40 -0.44 0.27) -0.44 4 / 0.51 ( 0.10) 0.00 -0.03 0.34) -0.03 / 2.60 2.80) 2.09 -0.74 0.10) -0.74 6 / 0.51 ( 0.10) 0.00 0.00 , 0.31) 0.00 / 10.00 3.13) 9.49 -0.80 k 0.05) -0.80 5 0 / 5.00 ( 1.93) 0.00 0.00 0.22) 0.00 / 5.00 2.03) 0.00 0.00 k 0.23) 0.00 2 / 5.00 ( 1.80) 0.00 0.00 0.22) 0.00 / 11.00 2.12) 6.00 -0.12 k 0.05) -0.12 4 / 3.50 ( 1.93) -1.50 0.18 k 0.22) 0.18 O G 20.00 2.00) 15.00 -0.15 ; 0.02) -0.15 6 / 3.00 ( 1.80) -2.00 0.29 0.22) 0.29 O G 20.00 0.00) 15.00 -0.07 [ 0.02) -0.07 13 0 / 13.00 ( 3.00) 0.00 0.00 0.12) 0.00 / 13.00 2.50) 0.00 0.00 ( 0.11) 0.00 2 / 12.75 ( 3.12) -0.25 0.02 [ 0.11) 0.02 O G 20.00 1.50) 7.00 -0.02 [ 0.04) -0.02 4 O G 12.00 ( 3.12) -1.00 0.06 ; o.io) 0.06 O G 17.00 0.00) 4.00 0.37 [ 0.03) 0.37 6 O G 12.25 ( 2.87) -0.75 0.06 ; 0.09) 0.06 O G 20.00 0.00) 7.00 0.40 ; o.oo) 0.40 E X P 0.5 0 DB 0.04 ( 0.34) 0.00 0.22 I 0-44) 0.00 D M 0.07 0.05) 0.00 -0.01 { 0.32) 0.00 fx = 0.3 2 DB 0.64 ( 0.58) 0.60 -0.60 [ 0.72) -0.82 O G 0.18 0.11) 0.11 0.40 [ 0.80) 0.41 rexp = 13 Gyr 4 DB 0.64 ( 0.57) 0.60 -0.54 [ 0.33) -0.76 / 0.20 0.04) 0.13 0.37 ; 0.32) 0.38 6 DB 0.64 ( 0.57) 0.60 -0.54 [ 0.32) -0.76 / 0.26 0.05) 0.18 0.10 ; 0.39) 0.11 8 DB 0.64 ( 0.57) 0.60 -0.54 [ 0.32) -0.76 / 0.32 ' 0.07) 0.25 -0.22 [ 0.34) -0.21 5 0 / 1.02 ( 0.00) 0.00 -0.24 [ 0.07) 0.00 O G 0.32 ' 0.03) 0.00 0.40 ; o.oo) 0.00 2 / 1.02 ( 0.11) 0.00 -0.24 I °-12) 0.00 O G 0.45 ; 0.05) 0.13 0.39 ; o.oi) -0.01 4 / 1.02 ( 0.11) 0.00 -0.24 [ 0.13) 0.00 O G 0.45 ; 0.05) 0.13 0.40 ( 0.00) 0.00 6 / 1.02 ( 0.11) 0.00 -0.24 ; 0.14) 0.00 O G 0.51 ; 0.06) 0.19 0.38 ( 0.02) -0.02 8 / 1.02 ( 0.11) 0.00 -0.24 I 0- 1 2 ) 0.00 O G 0.51 ; 0.06) 0.19 0.40 [ 0.05) 0.00 13 0 / 1.14 ( 0.12) 0.00 -0.12 ( 0.20) 0.00 O G 0.57 I 0.07) 0.00 0.40 [ 0.00) 0.00 2 / 1.02 ( 0.12) -0.12 0.18 { 0.20) 0.30 O G 1.02 [ 0.37) 0.44 0.00 [ 0.40) -0.40 4 / 1.02 ( 0.12) -0.12 0.18 ( 0.19) 0.30 O G 1.02 [ 0.28) 0.44 0.00 ( 0.34) -0.40 6 / 1.02 ( 0.12) -0.12 0.18 ( 0.19) 0.30 O G 1.02 { 0.26) 0.44 0.00 ( 0.00) -0.40 8 / 1.02 ( 0.12) -0.12 0.18 ( 0.19) 0.30 O G 1.28 [ 0.26) 0.71 0.00 ( 0.02) -0.40 Note. - A l l ages are in Gyr. The numbers in brackets are la confidence limits based on Monte Carlo sampling of the errors. tA check mark, / , indicates a reliable fit. Unreliable fits are coded as: "OG' if the point lies off the model grids, "DB' if it lies in the young double-valued Balmer line region, and " D M ' if it lies in a degenerate metallicity region. Table 3.9: Age & Metallicity Fits and Errors as a function of T V : U8a VS. (Fe) and H<$^vs. Fe4668 Model Parameters R6A V S . (Fe) vs. Fe4668 S F H Age TV fitt Age A A log10(Z/Ze) AZ fitt Age A A \ogw{Z/ZQ) AZ SSP 0.5 0 / 0.51 ( 0.15) 0.00 0.00 ( 0.15) 0.00 / 0.51 ( 0.10) 0.00 0.00 ( 0.18) 0.00 H = 1.0 2 / 0.51 ( 0.15) 0.00 0.00 ( 0.14) 0.00 / 0.51 ( 0.15) 0.00 -0.02 ( 0.15) -0.02 4 / 0.51 ( 0.15) 0.00 0.00 ( 0.26) 0.00 / 0.51 ( 0.10) 0.00 -0.02 ( 0.16) -0.02 6 / 0.51 0.15) 0.00 0.00 ( 0.28) 0.00 / 0.51 0.06) 0.00 0.00 0.09) 0.00 5 0 / 5.00 ( 2.12) 0.00 0.00 ( 0.14) 0.00 / 5.00 2.33) 0.00 0.00 0.22) 0.00 2 / 5.00 2.12) 0.00 0.00 ( 0.14) 0.00 / 5.00 2.38) 0.00 0.00 0.23) o.oo. 4 / 4.75 2.00) -0.25 0.04 ( 0.15) 0.04 / 4.00 2.05) -1.00 0.15 0.23) 0.15 6 / 4.50 1.88) -0.50 0.06 ( 0.15) 0.06 / 3.25 1.93) -1.75 0.24 0.22) 0.24 13 0 / 13.00 5.38) 0.00 0.00 ( 0.20) 0.00 / 13.00 4.00) 0.00 0.00 0.20) 0.00 ' 2 / 12.75 5.63) -0.25 0.02 ( 0.21) 0.02 / 12.75 4.13) -0.25 0.02 0.21) 0.02 4 / 11.75 4.75) -1.25 0.08 ( 0.20) 0.08 / 6.25 8.25) -6.75 0.39 0.42) 0.39 6 / 11.00 4.88) -2.00 0.13 ( 0.21) 0.13 OG 7.50 k 3.38) -5.50 0.37 , 0.21) 0.37 E X P 0.5 0 DB 0.04 k 0.43) 0.00 0.13 ( 0.43) 0.00 DB 0.04 [ 0.98) 0.00 0.28 0.99) 0.00 H = 0.3 2 DB 0.81 ; o.4i) 0.76 -0.68 ( 0.35) -0.81 DB 0.81 I °-42) 0.77 -0.72 , 0.76) -1.00 r e x p = 13 Gyr 4 / 0.81 ; 0.24) 0.76 -0.66 ( 0.20) -0.79 DB 0.81 [ 0.72) 0.77 -0.72 [ 0.74) -1.00 6 / 0.81 ; 0.24) 0.76 -0.66 ( 0.10) -0.79 DB 0.81 ; 0.72) 0.77 -0.72 ; 0.74) -1.00 8 / 0.72 ; cos) 0.68 -0.60 ( 0.09) -0.73 DB 0.81 ; 0.72) 0.77 -0.72 [ 0.72) -1.00 5 0 / 1.02 ; 0.12) 0.00 -0.28 ( 0.14) 0.00 / 1.02 I 0-12) 0.00 -0.25 I ° . i 4 ) 0.00 2 / 0.90 [ 0.11) -0.11 -0.11 ( 0.14) 0.18 / 0.90 [ 0.10) -0.11 -0.09 ; o.i6) 0.17 4 / 0.90 ; o.io) -0.11 -0.10 ( 0.15) 0.19 / 0.90 { 0.15) -0.11 -0.09 [ 0.24) 0.17 6 / 0.90 ; o.io) -0.11 -0.10 ( 0.15) 0.19 / 0.90 { 0.15) -0.11 -0.09 [ 0.23) 0.17 8 / 0.90 ; o.io) -0.11 -0.10 ( 0.16) 0.19 / 0.90 ; 0.15) -0.11 -0.09 ; 0.24) 0.17 13 0 / 1.02 ( 0.17) 0.00 0.06 ( 0.24) 0.00 / 1.02 { 0-17) 0.00 0.06 ( 0.26) 0.00 2 / 0.81 ( 0.21) -0.21 0.34 ( 0.16) 0.28 OG 0.81 ( 0.33) -0.21 0.37 ; 0.22) 0.31 4 / 0.81 ( 0.09) -0.21 0.34 ( 0.12) 0.28 OG 0.81 ( 0.33) -0.21 0.37 [ 0.21) 0.31 6 / 0.81 ( 0.09) -0.21 0.34 ( 0.10) 0.28 OG 0.81 ( 0.33) -0.21 0.37 [ 0.21) 0.31 8 / 0.81 ( 0.09) -0.21 0.34 ( 0.10) 0.28 OG 0.81 [ 0.33) -0.21 0.37 [ 0.21) 0.31 Note. - A l l ages are in Gyr. The numbers in brackets are la confidence limits based on Monte Carlo sampling of the errors, t A check mark, / , indicates a reliable fit. Unreliable fits are coded as: "OG' if the point lies off the model grids, "DB' if it lies in the young double-valued Balmer line region, and " D M ' if it lies in a degenerate metallicity region. Chapter 3. Dust Sensitivity of Absorption-Line Indices 106 7.5 7 6.5 4.5 as 4 3.5 3 8.5 2 1.5 1 1 1 1 1 i 1 1 1 1 1 1 - _»o.2 1 1 1 1 1 GALAXEV ; ; l i V SSP fj. = 1.0 " EXP = 0.3 T - r v \ A ( G y r ) 0.6 Gy r I 6 G j r -13 G j r -~ \ \ \ V - o.ooo1jf^ -v Jujl — 7 0.0004 ^ ^ ^ S s - X 5 0 L z 0 0 0 4 0.008 ' 1 , 1 , 1 , 1 . 1 , 1 10.0 20.0-0.05 " , 1 , 1 , ' 6 h 5 h X 4 0.05 0.1 0.15 0.2 0.25 0.3 0.35 0.4 GALAXEV • SSP fj. = 1.0 * EXP fj. = 0.3 A (Gyr) 0.6 G j r -j 6 G j r 13 G j r 0.2 1 1 1 1 1 1 1 1 GALAXEV - : SSP /J. = 1.0 EXP M = 0.3 I ~  l^A A ( Gy r) 1 Y \ M 0.6 G j r . 5 Gyr _ 13 G j r . - 0.0001 Y - \ ^ L 0.0004" - - ~ ^ ~^ \^ ~r^ ^ 2 0 0 0 4 0 0 0 8 ^ . 5 . 0 20.0 0.05 1 2 3 4 5 [MgFe]' <5i X 7 6.5 6 - ffl|s' 5.5 — Ira \ N. 1 5 — \ I J 4.5 : m 4 3.5 _ \ ni 3 ¥ 2.5 -0.0001 2 - 0.0004 1.5 -Figure 3.15: H/3 versus M g 2 [top left], H/3 versus [MgFe]' [top right], H/3 versus Fe4668 [bottom left], and H8A versus Fe4668 [bottom right] diagnostic plots. The grids are dust-free BC03 SSP models. Lines of constant age (dashed) are shown for ages 0.2, 0.5, 1.0, 2.0, 5.0, 10.0, and 20.0 Gyr. Lines of constant metallicity (solid) are shown for Z = 0.0001 (purple), 0.0004 (dark blue), 0.004 (light blue), 0.008 (green), 0.02 (yellow), and 0.05 (red). The circles are the indices measured from the solar metallicity SSP models with p = 1.0 for f y = 0 (largest point size) to f y = 6 (smallest point size), and the triangles are the exponential SFH model indices with r e x p = 13 Gyr and p = 0.3 for f y = 0 (largest point size) to f y — 8 (smallest point size). The red, green, and blue shades are for model ages of 0.5, 5, and 13 Gyr respectively. Chapter 3. Dust Sensitivity of Absorption-Line Indices 107 0 0 .02 0 .04 0.06 0.08 0.1 0 . 1 2 0 .14 0 .16 0 .18 TiO, 12 10 8 6 4 < X 2 0 - 2 - 4 1 1 1 i 1 i | - | i | i | . | i | i 1 • SSP fi = 1.0 -\ \ 0 5 * E X P = 0 3 -~ o.oooi 0.0004" * - , N A ! A (Gyr) 0 B ° * "A v \ 6 Gy r '_ Z \ 2 .0 0.004 • 0.008 N \ 5.0 S \ 10.0 • GALAXEV 0 . 0 2 X \ 2 0 . 0 " 0.05 , 1 . 1 . 1 , 1 , 1 , to 0 0.5 1 1.5 2 2.5 3 3.5 4 4.5 5 <Fe> _ 1 1 1 1 1 1 1 1 1 1 i 1 i 1 i 1 i 1 i_ 7 GALAXEV : - T>\& i \ \ °-5 \ \ ii^r** --\ A • SSP (ii = 1.0 ~ A EXP M = 0 . 3 -•o.oooi !f\ f--^ \ - 0 .0004 \y> A (Gyr) j T"-"~-2.0 -; 0.004 * -V A - . »\ i : 0.6 Gy r Z 7 6 Gy r L 13 G y r " . 1 . 1 , 1 , 1 , 0.008 ^ 5 ' ° 10.0 0.02^. 20 .0-0.05 -1 . 1 , 1 , 1 . 1 , " 0 1 2 3 4 5 6 7 8 Fe4668 Figure 3.16: H/3 versus G4300 [top left], H/3 versus T i 0 2 [top right], H8A versus (Fe) [bottom left], and r\5A versus Fe4668 [bottom right] diagnostic plots. Lines and symbols are as in Figure 3.13 Chapter 3. Dust Sensitivity of Absorption-Line Indices 108 3.6 Discussion As an illustration of the potential dangers of dust extinction when using absorption-line indices to determine ages and metallicities of star forming galaxies, in Figure 3.17 we present model tracks for the Lick H/3 versus (Fe) diagnostic plot for the dust-free exponential SFH with Texp = 13 Gyr for Z = 0.0004, 0.02, and 0.05 (gray curves). The T 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 [ j i i I i i i i I i i i i I i i i i I i i i i I i i i i L 0.5 1 1.5 2 2.5 3 <Fe> Figure 3.17: H/3 versus (Fe) diagnostic plot for dust-free exponential SFH models with TexP = 13 Gyr for three metallicities (gray curves), Z = 0.0004 (triangles), 0.02 (circles), and 0.05 (squares). Ages are labeled along the model tracks at 0.2, 0.5, 1.0, 2.0, 5.0, 10.0, and 20.0 Gyr. The coloured points (red, green, and blue shades for 0.5, 5, and 13 Gyr respectively) are the model indices measured for the same exponential SFH models, but with [i = 0.3 and dust extinctions from fy = 0 (largest point size) to fy = 8 (smallest point size). circles are the index values for the same SFH, but with dust extinctions ranging from fy = 0 (largest point size) to fy — 8 (smallest point size). There is virtually no error for Chapter 3. Dust Sensitivity of Absorption-Line Indices 109 old and metal-poor galaxies, but significant problems can arise at higher metallicities. For example, the location in the H/3 versus (Fe) plane for a solar metallicity stellar population at any age with fy = 1 lies closer to the Z = 0.05 model curve. The age errors are not significant in this case, but as extinction increases, the data points lie significantly outside the region covered by the model grids. However, since in reality we don't know the true SFH of a galaxy, what is often computed are "SSP-ages and metallicities", i.e., the values that would be measured if a given index-index combination is fit to SSP model grids. In order to gauge the magnitude of the errors on SSP ages and metallicities due to dust extinction, we fit the model galaxies to SSP grids of different index-index combinations. The error on the derived parameters is taken as the difference between the fit for the dusty model and the corresponding dust-free model (i.e. A A = Age(fy) — Age(fy = 0)). In some cases we encounter difficulties with the models with current SF, as these models can lie off the SSP model grids, even in the dust-free models. Additionally, due to the double-valued nature of the Balmer lines at young ages (discussed in §3.4.1), there is an additional degeneracy for points that lie in the youngest region of the SSP model grids. This was the case for all of the exponential SFH models at 0.5 Gyr, for which a large fraction of the stars are very young. Details of the fits and tables providing the resulting age and metallicity errors for a selected number of index-index diagnostic plots are provided in §3.5. The primary result from this analysis is that, in general, when dust extinction affects the SSP age and metallicity determinations, the errors on the physical parameters are of the same order as the lo confidence intervals from typical measurement errors, and thus would not likely be detected above the noise. The most notable exception is any index-index plane using the Dn(4000), which suffers significant errors even in the SSP models. Index-index grids using the longest baseline indices can also suffer significant errors on the measured physical parameters due to the effects of dust. Figure 3.18 illustrates the potential pitfalls of discounting dust extinction in an analysis that uses index diagnostics to determine SFHs. There, we plot R8A versus Chapter 3. Dust Sensitivity of Absorption-Line Indices 110 Dn(4000) for the dust-free SSP (gray curves) and exponential SFH (purple curves) models for solar and Z — 0.05 metallicity and ages up to 20 Gyr. The triangles represent the JJL = 0.3 model indices for fv = 0 (largest point size and darkest shade) to f y = 8 (smallest point size, lightest shade). The locus of galaxies in the H5^-Dn(4000) plane has been used to determine their SFHs (e.g. Kauffmann et al. 2003a). Clearly, if a galaxy contains a significant amount of dust, its location in this plane will be best fit by a SFH that is different from its true SFH. As shown, the exponential SFHs move closer to the SSP region when dust is added. 12 1 1 1 1 1 i | , | i 1 , | i | i | , | , 10 - 0.2 <f- GALAXEV -| 8 L 0.05 V 0 . 5 • Z = 0 .02 -6 ; o.oi */ i • z = o.o5 : 4 ro.oooi#* -2 0 " E X P r e x p= 13 Gyr \ \ Age (Gyr) " - 2 - 4 - 6 - 0.5 Gyr 5 Gyr " 13 Gyr \ 5 . 0 NR10.0 ^Xia.o > i 20.0 " 1 . 1 r 0.8 1 1.2 1.4 1.6 1.8 2 2.2 2.4 2.6 2.8 D n 4 0 0 0 Figure 3.18: H6A versus Dn(4000) diagnostic plot for dust-free SSPs of solar (circles) and Z = 0.05 (squares) metallicity (gray curves) and an exponential SFH with r e x p = 13 Gyr for the same two metallicities (purple curves). Ages are labeled along the model tracks at 0.0001, 0.01, 0.05, 0.2, 0.5, 1.0, 2.0, 5.0, 10.0,13.0, and 20.0 Gyr. The coloured points (red, green, and blue shades for 0.5, 5, and 13 Gyr respectively) are the model indices measured for exponential SFH models of the same two metallicities, but with dust extinctions from fv = 0 (largest point size) to f v = 8 (smallest point size). Chapter 3. Dust Sensitivity of Absorption-Line Indices 111 Fortunately, in the Rose dwarf/giant diagnostic plot (Sr II A4077/Fe IA4063 versus H5/A4063), the dust effects run along the lines of constant surface gravity. Thus, this diagnostic plot is still useful for determining the relative amounts of young versus evolved stars, even when dust is present. It must be emphasized that our analysis has not accounted for emission-line contam-ination. Dust effects aside, this could potentially be the most worrisome pollutant for all age determinations involving Balmer lines, particularly for the youngest stellar pop-ulations. While considering the higher-order Balmer lines ( H 7 and B.8) greatly reduces the problem, it does not go away altogether. To summarize, we have considered dust effects on absorption-line indices using the Bruzual & Chariot (2003) population synthesis models incorporating the multi-component model of CF00 for the line and continuum attenuation due to dust. For qui-escent stellar populations (e.g. spheroids and globular clusters), dust extinction effects are small for most indices. A notable exception is the 4000 A break, whose sensitivity to dust can translate into significant errors in the age determination of the stellar pop-ulation. For models with current SF, many of the indices are significantly modified due to dust reddening effects, and their behavior depends on age, dust distribution, and the effective optical depth. Unfortunately, no simple dust-correction can be prescribed, but future spectroscopic studies of stellar populations in dusty environments (e.g. late-type spirals) ought to consider possible effects due to dust attenuation in their measurements and its resulting effect on physical interpretations. Fortunately, there are particular in-dices that are negligibly affected by dust extinction (compared to the current level of measurement precision) at different regions in the physical parameter space, and the safest current approach consists of mapping a broad range of indices with the shortest wavelength baselines. 112 Chapter 4 Spectroscopic Age and Meta l l i c i t y Gradients in Spira l Galaxies 4.1 Introduction We have embarked on a long-term study of the large-scale chemical and evolutionary properties of nearby galaxies. The main objectives of this program are to produce a comprehensive study of the age and metallicity distributions in bulges and inner disks and their systematic behavior with galaxy type, bar structure, and environment, and to constrain basic theories of bulge formation (see §1). In Chapter 2, we presented an extensive study of spiral galaxy stellar populations using radial colour profiles. We found that bright spiral galaxies exhibit clear trends for the effective age and metallicity in spiral galaxies as a function of surface brightness, luminosity, rotational velocity, and size. Higher SB regions of galaxies formed their stars earlier than lower surface brightness ones, or at a similar epoch but on shorter timescales. Also, the SFHs at a given SB level, which lead to age gradients, are modu-lated by the overall potential of the galaxy such that brighter/higher rotational velocity galaxies formed earlier. These trends are of course at odds with hierarchical galaxy for-mation which predicts that more massive galaxies form late. We saw that feedback processes are often invoked as a possible mechanism to prevent the gas in lower mass galaxies from cooling and forming stars at early times. As appealing as this solution may appear, no viable prescription for feedback has been found thus far that agrees with all observational constraints. An important limitation of our broadband colour gradi-ents analysis is its sensitivity to dust extinction which could mimic (or hide) gradients in age and, particularly, metallicity. Turning to spectroscopy and the measurement of absorption-line strengths could Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 113 greatly alleviate this degeneracy due to the shorter measurement baselines and the stronger sensitivities of individual features to the stellar population characteristics. In order to reassure ourselves of the relative dust insensitivity of line-indices, we carried out a qualitative study in §3 which highlighted a few potential dangers in this assumption (especially concerning the 4000 A break and very young composite stellar populations). However, for the most part, this study demonstrated that dust will not play a major role in the measurement of absorption-line indices. Hence, to address the many astrophysical goals outlined in §1, we have obtained longslit spectroscopy for a pilot sample of 8 nearby field spirals from which we measure a suite of absorption line-indices in their bulges and inner disks. We then compare the line-indices with stellar population synthesis (SPS) models to derive luminosity-weighted single age stellar population equivalent ages and metallicities. Errors on these parameters will be mainly systematic, arising from model limitations. However, we note that the differential comparisons of subsamples (bulge versus disk within a galaxy, abundance variations among galaxies of a given size or luminosity) should be reliable to a much higher degree than with any absolute estimate (ages, metallicities, and abundance ratios). Previous work in this field was summarized in §1.2.2 where we pointed out that there are currently no studies that explore simultaneously the bulges and disks of late-type galaxies spectroscopically. Thus this is the first systematic survey of bulge and disk spectroscopic gradients along the Hubble sequence of spiral galaxies with a large aperture (> 8-m) telescope. This difficult theoretical and experimental program may not result in one unique solution. While degenerate solutions may exist, the model-independent comparison of line gradients as a function of galaxy type and environment is poised to guide theoretical developments and lay the foundation for modern spectral investigations of galaxy evolution. The outline of this chapter is as follows. In §4.2, we describe our observational set-up, choice of sample and standard stars, as well as observing strategy. Data reduction is summarized in §4.3 and the data analysis with the extraction of binned ID spectra Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 114 as a function of radius is presented in §4.4. The spectroscopic indices and line strength gradients are presented, analyzed for empirical trends, and compared with SPS models in §4.5. Comments on individual galaxies are presented in §4.6. 4.2 Longslit Spectroscopic Observations 4.2.1 Gemini G M O S The longslit spectroscopic data for this study were collected using the Gemini Multi-Object Spectrograph (GMOS; see http: //www. gemini. edu/sciops/instruments/gmos/ or Hook et al. 2004 for reference) on the 8-m Gemini North telescope at Mauna Kea in Hawaii. The GMOS detector consists of three 2048x4608 CCDs with 13.5 pm pixels providing a spatial resolution of 0.072 "/pix and a dispersion of 0.45A/pix with the B600-G5303 grating. The internal stability of the slit mask requires two equally spaced bridges between the slit edges. The dispersion (spectral) direction of GMOS is along rows and spatial direction (slit orientation) is along columns. The slit is parallel to the columns, thus the slit bridges show up along the spatial direction, whereas the gaps be-tween the 3 CCD detectors manifest as small gaps in the wavelength coverage (~17A). This is demonstrated in Figure 4.1, which shows a sample spectrum of a GMOS longslit observation. The efficiency curve of the B600-G5303 grating, which was used for all observations, is shown in Figure 4.2. Its high sensitivity in the blue matches well our simultaneous spectral coverage of 4050 - 6750 A. This includes most of the major atomic and molecular features to disentangle age and metallicity effects in integrated galaxy spectra (see e.g. §3.2 and Tables 3.1 & 3.2). We used the longest possible GMOS slit to allow for observations of some of the biggest local galaxies. The slit field of view (FOV) was 5' (length) x 2" (width) for all of our observations. The choice of slit-width was prescribed by the need to maximize signal-to-noise (S/N) in the disk outer parts while maintaining adequate spectral resolution throughout. The 2" slit and B600 grating give 10.75 A full width at half maximum (FWHM) resolution (as measured from the width of the narrowest sky emission lines). The absorption features are typically ~40 A wide. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 115 Thus the choice of grating B600-G5303 gives adequate resolution of each feature and measurement of the 24 spectral (Lick/IDS) indices between 4050 and 6750 A. Wavelength Figure 4.1: A sample GMOS longslit observation of a nearby galaxy indicating the slit bridges in the spatial direction (along columns) and the gaps between the 3 CCD detectors in the dispersion direction (along rows). Wavelength increases from right to left (Figure adapted from http://www.gemini.edu/sciops/data/dataIndex.html). Before we discuss our basic reduction procedures, followed by further data analysis, let us first turn to the description of our new Gemini sample. 4.2.2 Galaxy Sample Ideally, we would like to collect as large a sample of nearby late-type spiral galaxies as possible to study spectroscopically the stellar populations of spiral bulges and disks. As we have stressed numerous times in this thesis, the stellar populations of late-type disks based on absorption spectroscopy have not yet been studied in a systematic Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 116 3000 4000 5000 6000 7000 8000 MA) Figure 4.2: Efficiency curve (blue squares and line) for the B600-G5303 grating on GMOS-N. Our wavelength range (~4050-6750 A ) is delineated by the vertical red dotted lines. way and with sufficient S/N. A few studies of absorption-line spectroscopy of early-type spiral bulges (SO-Sb) do exist (Fisher, Franx, k. Illingworth 1996; Goudfrooij, Gorgas, & Jablonka 1999; Proctor & Samson 2002). These studies suggest that, while early-type bulges share at least some properties with intermediate luminosity ellipticals (correlation of Mg indices with velocity dispersion, negative metalicity gradients, both positive and negative Balmer-line gradients, in addition to the photometric properties discussed in §1.2.1), spiral bulges are generally younger (Sa-Sb: 1.5-6 Gyr, SO: 2 -7 Gyr, E: 4-13 Gyr), have smaller light-element enhancements ([Mg/Fe] ratios closer to solar), and span a wide range in [Fe/H] at a given age (i.e. metallicity does not correlate with age in early-type spiral bulges, as it does in ellipticals), and Fe indices also correlate with velocity dispersion (the correlation only exits for the Mg indices in ellipticals), strongly suggestive of a mass-metallicity relation for spiral bulges. Early-type spirals with their large bulge-to-disk ratio offer a significant observational advantage over late-type galaxies in that their bulges, which rise significantly above and below the disk, can be studied spectroscopically in the edge-on perspective free from Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 117 disk contamination and extinction from dust. Spectroscopic observations of late-type bulges cannot be made in the edge-on view because the small bulge is overwhelmed by the disk. The best alternative to separating bulge and disk populations is to view the disk face-on. While dust and gas from the disk will still thwart any pristine observation of the bulge, the face-on orientation minimizes line-of-sight integrations and rotation corrections, and keeps the light distribution of the bulge and disk free of inclination effects. The study of stellar populations in late-type spiral bulges, which rise to similar heights above the galaxy midplane as the disk, thus requires a sample of face-on galaxies. We further want our sample to include barred spirals in order to test for the expected mixing effects by a bar. Martin & Roy (1994) have shown, based on the [O/H] ratio from gaseous emission lines in bright spirals, that abundance gradients are weaker in barred galaxies. This analysis must be extended to absorption gradients as well. Inclusion of barred galaxies in our sample may also shed additional light on the nature of Freeman Type I and II galaxies (MacArthur, Courteau, & Holtzman 2003) from a spectroscopic point of view. Ultimately, the goal of our long-term study is to collect a sample which covers as wide a range of physical properties as possible: morphological type (Sa-Sd), IR luminosity (IR-quiescent to IR-luminous), and FIR/optical ratio (over a few orders of magnitude). These would be measured as much in absorption as in emission to map both the stellar and the gas phases. Other relevant parameters to study include: nuclear activity (quiescent, starburst, LINER, Seyfert), bar structure, B / D ratios, star formation rates (SFRs), gas content, and dust properties. The size of the current sample was a compromise between a reasonable telescope time request for a pilot study and the need to measure systematic variations within that sample. We have thus narrowed in on the following criteria: • Hubble type Sa-Sd (with emphasis on latest types) • Mix of barred/unbarred systems and SB profile types • Face-on (inclination < 35°) Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 118 • Blue Galaxtic extinction A B = 4 x E(B - V) < 0.5 mag (Schlegel et al. 1998) • Available optical and IR calibrated surface brightness profiles. • Blue major axis less than twice the slit length (< 10' 5). • Large enough galaxy size to resolve bulge and inner disk gradients. The selected galaxies must also have a guide star within the FOV of the GMOS peripheral wavefront sensor camera and be observable from Mauna Kea in the fall semester (RA~23-03h and DEC > - 20°). The face-on orientation and size criteria were not trivial to meet. Of the two major catalogs of optical and infrared imaging for late-type spiral galaxies available to us at the inception of this program (de Jong 1996; Courteau, Holtzman, &z MacArthur, in prep.), only the former has large enought face-on galaxies with K-band information. Five of the 8 selected galaxies were indeed observed by de Jong (1996) in BVRIK bands. The desire to include a few larger and/or later-type candidates prompted us to consider other databases which may not have full bandpass coverage. The galaxy sample and catalog information from NED are shown in Table 4.1. Bulge and disk photometric parameters, where available, are also shown in Table 4.2. Note that in this section we refer to the disk scale length as r^, the bulge effective (half-light) radius as r e, and the bulge scale length as r^ (converted from re using the Sersic n, see §4.4.6). The largest galaxy in our sample, NGC 628 (also M74), lacks extended photometric coverage but was observed to demonstrate our ability to separate bulge and inner disk spectroscopic gradients in well-resolved galaxies. That galaxy is also part of the SINGS Legacy project (http://sings.stsci.edu/indexl.shtml) which ensures availability of other useful photometric and spectral data products for future studies. 4.2.3 Lick Standard Star Sample Standard star observations with the same observational set-up are required for flux calibration of the galaxy spectra. Standards from the Lick catalogue (Faber et al. 1985) Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 119 Names Hubble Diameter R A D E C Vhelio M B Af N G C / I C U G C Type a ( ' )xb( ' ) 2000 (kms - 1 ) (mag) (mag) NO 173 369 SA(rs)c 3.2x2.6 00 h 37 m 12?2 +01°56'31'. '5 4366 13.70 0.110 N0628 1149 SA(s)c 10.5x9.5 01 h 36 m 41?7 +15°46'59'. '4 657 9.95 0.301 N1015 U2124 SB(r)a 2.6x2.6 0 2 h 3 8 m l l ? 6 -01°19 '06 ; '6 2631 12.98 0.139 N 7 4 9 0 12379 Sbc 2.8x2.6 23 h 07 m 25?2 +32°22'30: '2 6213 13.05 0.362 N 7 4 9 5 12391 SAB(s)c 1.8x1.7 23 h 08 m 57?2 + 12o02'52!'9 4887 13.73 0.371 N7610 12511 Scd 2.5x1.9 23 h 19 m 41?3 +10°11 '05: '7 3554 13.44 0.171 N7741 12754 SB(s)cd 4.4x3.0 23 h 43 m 54?0 +26°04'33! '8 751 11.84 0.324 10239 2080 SAB(rs)cd 4.6x4.2 02 h 36 m 27?9 + 3 8 ° 5 8 ' l i : ' 7 903 11.80 0.307 ^Galactic extinction values are from Schlegel, Finkbeiner, & Davis (1998) Table 4.1: Galaxy Sample (catalog information). Names highlighted in bold are those used subsequently to refer to the particular galaxy. are also required for calibration to the Lick/IDS system, which is plagued by variable resolution (degrading towards the blue) and poor flux calibration. This remains a crucial step for direct comparison with other similar studies and SPS models calibrated to the Lick system. The integration time per standard star was 2 seconds (but with Gemini overhead per exposure exceeding 20 minutes). Table 4.3 lists the 7 Lick/IDS standards (Worthey et al. 1994) that were selected for our project, along with their relevant physical parameters (spectral type, effective temperature, surface gravity, and metallicity). We will carry out the calibration of our galaxy indices to the Lick system in a future work. For the current study, however, we compare our galaxy indices to the high reso-lution models of Bruzual & Chariot (2003) which provide flux calibrated spectra. We can thus compare our galaxy indices with the G A L A X E V model predictions measured directly from the observed spectra provided our data are accurately flux calibrated (rel-ative, but not absolute). Flux calibration was achieved by comparison of our standard star observations with flux calibrated spectra of stars in common with the ELODIE archive (Moultaka et al. 2004; see http://atlas.obs-hp.fr/elodie). Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 120 NGC 0173 - SA(rs)c M NGC 628 - SA(s)c Figure 4.3: Observational set-up for the eight galaxies in our sample. The background images are from the Canadian Astronomy Data Centre's Digitized Sky Server (CADC; http://cadcwww.dao.nrc.ca/). The blue line represents the slit, the red (dashed) box and long arm represent the FOV of the G M O S wavefront sensor camera, with the box at the end of the arm centered on the guide star. The panels for large galaxies (NGG 628, NGC 7741, and IC 239) also show, as yellow lines, the sky offset positions. The FOV for the CADC pictures differ in all the panels but the slit length is everywhere the same (5'). Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 121 NGC 7741 - SB(s)cd I . IC 239 - SAB(rs)cd •. Figure 4.4: Figure 4.3 continued. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 122 Name Td l^e,bulge re Sersic Srct (mag/D") (") (mag/D") (") n N0173 20.20 19.2 N / A N / A N / A 3 N0628 18.30 73.1 18.32 19.5 1.35 2 U2124 18.11 21.3 15.89 2.7 1 1 N7490 17.44 15.7 15.33 1.9 1 1 N7495 17.79 13.8 17.36 1.2 1 1 N7610 18.35 13.2 18.09 2.1 1 1 N7741 N/P 49.2 N/P 7.6 1 1 10239 N / A N / A N / A N / A N / A tSources: 1 = de Jong (1996) K-band ( N / P = non-photometric) 2 = Mollenhoff k Heidt (2001) K-band 3 = Gr0sbol (1985) PSS-R N / A = not available Table 4.2: Galaxy Structural Parameters 4.2.4 Integration Times and Observing Strategy Our integration times of 45-72 minutes give S / N /A~ 5 at typically half a disk scale length (i.e. within the bulge/disk transition). The need for large (> 8-m) apertures to resolve most of the line features with S/N/A>20 beyond the inner disk, and correct for nebular emission and the presence of disk stars into the bulge is therefore obvious. An estimate of the sky background was measured from the slit edges for galaxies with major axis diameters < 4'. For galaxies wider than the slit, we divided our exposures Name Type RA DEC Teff log(s) [Fe/H] Mv B — V HD 000319 A1V 00h07m46? 48 -22°30'33!'3 8140 3.8 -0.7 5.933 0.14 HD 222451 F1V 23h40m39? 78 +36°43'13'/3 6632 4.29 0.07 6.245 0.395 HD 225239 G2V 00h04m50? 98 +34°39'3l!'2 5775 4.0 -0.45 6.109 0.626 HD 004628 K2V 00h48m20? 77 +05°17'43!'2 4960 4.6 -0.29 5.742 0.89 HD 221756 A1III 23h34m37? 71 +40°14'li:'l 9020 3.91 -0.5 5.589 0.093 HD 010975 K0III 01h48m38? 90 +37°57'10:'0 4786 2.4 -0.33 5.936 0.974 HD 218804 F5IV 23h10m27? 96 +43°32'48!'l 6308 4.06 -0.24 5.93 0.437 Table 4.3: Table of Lick/IDS standard star sample parameters. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 123 into sky/galaxy/sky sequences to filter out cosmic rays and confirm weak features. Sky exposures (before, middle, and after) were binned to achieve the same S/N as the galaxy. Chopping provides a sky subtraction accuracy of 1% (yielding continuum counts with S/N>40/A/" in the red)1. The typical on-target integrations of 3 x 15 min for MB < 10.5, 3 x 18 min for 10.5 < MB < 12, and 4 x 18 min for MB > 12 galaxies per major-axis spectrum was sufficient to resolve the higher-order Balmer lines out to 1 disk scale length. Integration times for sky (off target) measurements for the large galaxies were ~ l / 5 of the on-target integrations. (In hindsight, our adopted off-target integrations were a bit too short to provide the expected S/N in the inner disk. Integrations about 1/3 to 1/2 of the on-target exposures would have been more adequate.) 4.3 Data Reduction While the author spearheaded the Gemini Phase I and Phase II proposals, the data were chiefly reduced by Prof. Jesus Gonzalez (UNAM), in consultation with the author and Prof. Stephane Courteau. It would thus be inappropriate to include extensive dis-cussions of the data reduction process in this thesis. The detailed data reductions will be described in a forthcoming paper by Gonzalez, MacArthur, & Courteau. In the fol-lowing section, we summarize the major data reduction steps for a basic understanding of the data quality and limitations. The data reduction and analysis for this study were carried out using the XVISTA image processing package developed at Lick Observatory2. Similar to IRAF (Tody 1993), XVISTA is a versatile program that provides a large set of tools for performing complex data reductions and analyses. It is also highly adaptable/configurable through user-written procedures, or direct modification of the source code. In the following description of the data reduction process for our longslit spectroscopic data, specific 1 Note that the "Nod & Shuffle" observing technique available at Gemini would not benefit this program of extended sources due to significantly reduced slit coverage. 2 X V I S T A is maintained and distributed by Jon Holtzman at New Mexico State University and can be downloaded from http://ganymede.nmsu.edu/holtz/xvista/ as of writing. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 124 reference to the XVISTA commands and jargon will be kept to a minimum. The relevant measurements for this study (absorption-line indices, rotations) do not require absolute spectrophotometric calibration of the data. However, they do require extreme care in removing instrumental effects in order to achieve accurate relative calibration along both the dispersion and spatial directions. Proper subtraction of the sky and instrumental and far-field scattered light is also crucial. Finally, accurate flux calibration is also required for the galaxy and standard star observations, particularly for long-baseline measurements such as the molecular indices. The basic steps involved in CCD data reduction can be broadly separated into two categories: 1) Small-scale (high-frequency) effects that are associated to the detector itself (pixel-to-pixel sensitivity variations, detector defects, cosmic ray hits, etc.); and 2) Large-scale (low-frequency) effects that depend on the specific image being projected on the detector (fluxing, geometric distortions, etc.). The small-scale effects in (1) must be removed before performing the reduction steps in (2) that involve interpolations or image shifts. In the following sections we concentrate on the description of the instrumental setup, the special observational techniques, and the reduction steps up to line-strength extrac-tion. Readers only interested in extracted line-strengths and results therefrom should skip ahead to §4.4. 4.3.1 Characterization of CCD device(s) In astronomical image analysis, we must correct for features that belong to the detector and not to the image it records. This is especially important when combining images from different devices or images from the same device taken at different times. Impor-tant characteristics include readout noise and sensitivities that change among CCDs and/or as a function of time. We use simple but robust techniques to characterize the CCDs used and compare with the nominal (manufacturer) values. We must then Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 125 decide which of the nominal or measured values will be used for the ongoing analysis. For the full data set, each of the 3 CCDs was reduced as a separate image for the basic reduction steps. The images were later stitched together for the absorption-line index measurements. Modern CCDs are very linear over a large dynamic range, but the linearity is not conserved over the full pixel well range, in particular at very low count levels as well as at high count levels where it can set in before the A / D saturation limit, thus it must be carefully monitored. All of our images were found to have levels within the linear range of the GMOS CCDs. 1. Readout Noise & Gain The readout noise associated with a CCD detector is the fluctuation in the number of electrons introduced per pixel into the final signal upon readout of the device, and is denoted aro. It results from a combination of the not perfectly repeatable conversion from an analog signal to a digital number (A/D conversion) and thermal electrons coming from the electronics. For modern CCDs, the readout noise is typically quite low (of order < 10e~s per pixel) and is often not the dominant source of noise, but it still must be taken into account. The gain of a CCD determines the conversion from the charge collected in each pixel to a digital number and is set by the output electronics. Gain is usually given in terms of the number of electrons required to produce one analog-to-digital unit (ADU). Readout noise can be determined most reliably from bias (zero exposure time with shutter closed), dark (shutter closed, but non-zero exposure time), or flat-field (either internal, dome, or twilight sky; see below) calibration frames and is measured as the variance (RMS) with respect to a smooth version of a calibration frame. Gain is mea-sured from the variance of smoothed versions of calibration frames taken over a wide range of intensities (it is the inverse slope of a plot of variance as a function of counts), a% = N/Gain + a r20, (4.1) where N is the number of counts, and the variances are in units of counts. The Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 126 most stable results were obtained with the fiat-field frames, from which we measured the following gain and readout noise for the 3 CCDs: Gain (e /ADU) Read Noise ( e /pixel) CCD Nominal Actual Nominal Actual 1 2.04 2.109 3.5 4.43 2 2.32 2.334 3.3 3.42 3 2.19 2.260 3.0 3.68 The values we measure are close to, but not exactly the same as the nominal values and we opted to use our measured values for the reductions. 2. Cosmetics: identification of bad columns, ion traps, etc. Cosmetic defects associated with the CCDs must be identified and removed before re-scaling the image to avoid spreading their effect by the re-binning process. Astronomical grade CCDs usually have good cosmetics and their bad columns and moderate-to-severe charge traps can be easily identified and corrected by simple interpolation. For each detector, a map of the non-linear pixels and columns with bad charge transfer was made. A procedure then removes those defects from all data frames by linear or quadratic interpolation across the bad areas. 4.3.2 High-Frequency CCD Effects and Defects 1. Baseline, Bias, & Dark Current Subtraction The overscan area (pixels on the CCD that are not exposed to light) is used to subtract the background shape along the slit. The overscan columns are averaged and smoothed before subtraction so as not to increase the shot noise of the image or impose on it a spurious high-frequency pattern along the rows. A smooth version of the bias frames is also subtracted after overscan subtraction. The final bias frame is an average of all reliable bias frames from one or several nights if they are consistent with each other. Reliability is assessed on the basis of shape constancy along the slit. The dark current generated over time by the detector/amplifier is usually so low (< le~/hour) that it can be neglected. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 127 2. High-Frequency Flat-Fielding In order to map pixel-to-pixel sensitivity variations that are independent of where and what type of image is projected onto the CCD, one must collect frames with accumulated counts such that the total S/N is higher than the modeled variations and that has structure only on large scales, i.e. flat-field frames. These can come from internal projector lamps, a lamp reflected at the dome, or the twilight sky. Pixel-to-pixel variations in modern CCDs are typically at levels of less than a percent. We map sensitivity variations of scales ranging from a single to tens of pixels across, by dividing all images by the ratio of the highest S/N combined flat image with a smooth version of itself (to remove the large-scale spectral and geometric shapes of the flat). Pixel-to-pixel variations are thus effectively removed (to within less than a tenth of a percent). This procedure also effectively removes response effects due to out-of-focus dust or condensations on the CCD entrance window as well. Variations on larger scales are better mapped later in the reduction process (see §4.3.3). 3. Creation of a Variance Image The images at this point in the reduction process can be used to estimate the total noise expected at each pixel since the light contributions from the sky and scattered light backgrounds have not yet been subtracted. Because our exposure levels are well within the linear response range of the CCDs, the RMS variance expected from photon and readout noise (added in quadrature) at each pixel, o\, can be estimate as, o^col, row) — I (col, row) + o2TO, (4.2) where I(col, row) = Gain x N(col, row) is the number of photons in pixel (col, row) and the <rs are in units of photons (or, equivalently, e~s since each photoelectron in the CCD is generated by a single photon). The variance frames for each CCD were created by first transforming from counts to photons according to the measured gain. An image for each CCD was then created with the expected variance from the photon noise, and the gain and readout noise expected from the CCD characteristics. The variance image is a crucial tool which is carried Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 128 along all the data reduction and analysis steps under proper error propagation to assess S/N in any detections. 4- Scattered Light & Residual Background Subtraction The next step in the reduction is to subtract the spurious sources of background light. The overscan area provided a good estimation of the zero-point of the photon intensity along the slit, but residual background and far-field scattered light along the dispersion was still present and made it difficult to properly scale images with different exposure times or with sources that illuminate the slit differently. Hence, only the overscan rows separated from the slit by more than the instrumental FWHM were combined and smoothed along the dispersion axis to construct the spectral pattern to be removed across the length of the slit. 5. Combination of 3 CCDs to a Common Frame The three GMOS CCD images (data and variance frames) were combined into a single frame accounting for the nominal gap separations. The gaps, of typical width (~ 40 pixels), are interpolated across with a low-order interpolator. At this point, the main instrumental patterns associated with the individual CCD pixels have been removed. The images are now ready for transformation from the pixel {col, row) to physical (A,Rad) scale. 4.3.3 Low-Frequency Effects In longslit spectroscopy, each source on the slit generates its own spectrum that does not follow a straight trajectory along the CCD. Not only are the individual spectra curved, but they are not necessarily parallel to each other. Additionally, the spatial and dispersion are not necessarily orthogonal. Hence, simple translation + rotation (i.e. ID) transformations are not adequate for the extraction of spectra from longslit spectroscopic data, we must combine sources with known spectral (lamps, twilights) and spatial (slit bridges) scales. The procedure followed to make the 2D transformation from the {col, row) to (A,Rad) scale is described below. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 129 Geometrical Distortion Mapping 1. Distortions Along the Slit (formerly S-distortion, but more general) Distortions along the slit can be very precisely measured with a slit consisting of a series of holes along the slit length (e.g. Courteau 1992; Gonzalez 1993). In our case, the closest configuration to mask along the GMOS slit is the obscuration pattern left by the two slit-bridges on the slit-mask (see Fig. 4.1). The centroid of each obscuration was mapped using the flat-field images taken before each observation. This enabled us to refer each observation to the same reference points on the slit and effectively removes any residual flexure shifts along the slit. Since only two reference points are available for us on the slit, we map and correct the relative image-scale along the slit to within its linear term only (second order scale variations are small though, at least for our purposes). 2. Physical Scale Along the Slit The absolute image-scale along the slit is normally measured by placing two or more objects on the slit with know separation in the sky (such as the stars in an open cluster). Because Gemini overhead time is a premium, we did not perform this measurement and assumed the nominal scale of 0.072"/pix. 3. Distortions Along the Dispersion While the dispersion scale does not abruptly change across CCDs, it will show a discontinuity if the assumed gaps differ from the actual (real) separations between CCDs. The actual gap separations and tilts were mapped (to ~ a tenth of a pixel) from the calibration lamp and twilight frames. The gap-widths (39 & 41 pixels) were slightly different from the nominal value (37 pixels) and the second gap showed a tilt of 0.07 degrees which corresponds to an added separation of 5.2 pixels at the top of the CCD. Gap widths and tilts in all images were constant to within a tenth of a pixel and were assumed fixed for all our data. For each observation series, the wavelength - pixel relation at the center of the field was derived with a low-order polynomial, taking into account the discontinuous jump to correct for the actual gap separations (and see 4-Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 130 below for off-center correction). The fields along the slit reach the grating with different angles producing a line-curvature that varies with wavelength. For our grating, the line-curvature produces a shift of the order of 20 pixels at the end of the slit, and varies a good fraction of a pixel along wavelength. These variations were calibrated with a low-order 2D function (a polynomial for every row, with coefficients that vary smoothly along the slit), mapped from the calibration lamps (taken before and after) for each observation series. 4- 2D Distortion Removal The full 2D transformation Rad(coZ, row), X(col, row) between CCD rows and columns to spatial and wavelength bins was mapped and carried out in two steps, the first involv-ing an interpolation across rows (along the slit) such that the spectra run along straight parallel lines, and the second involving a transformation across columns (wavelength) to remove the line curvature. a) Along the slit: Rad(coZ, row) = SRad(coZ) x row + ARad(coZ), where the scale SRad(coZ) and zero-point ARad(coZ) factors are constants (along each column) that vary smoothly from column to column. This transformation was applied to all images before mapping the 2D wavelength calibration. b) Along wavelength: A(coZ,Rad) = A(coZ,Rad=0) + AA(coZ,Rad), where A and AA are polynomials that take into account the discontinuities due to the gaps and tilts among CCDs. 5. Resolution and Sampling The images were re-binned to linear scales along the slit and along wavelength corresponding to 0.2" x 2.0 A bins. The effective spatial resolution (combination of seeing, instrument and guiding errors) was estimated from each frame using point-like sources (such as the Ha image of H II regions, or the continuum of stars or H II regions along the slit). It varied from about 0.7" to 0.9". Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 131 Low-Frequency Response Calibration 1. Response along the slit The flux variation along the slit of bright night-sky emission lines was used to map the slit illumination of each galaxy or sky frame. The line flux was estimated after subtracting the sky or galaxy continuum close to the emission line. 2. Response along wavelength Each image was divided by a smoothed flat-field to correct, to first order, for intermediate-to-low frequency instrumental response. We did not have flux standards observed with the same set-up close to the dates of our observations. Thus, to re-move the spectral shape of the flat fields we compared our stellar data with previous observations of the same stars from the ELODIE archive (Moultaka et al. 2004). 3. Cosmic Ray Hits Interpolation Cosmic ray hits were identified as high-deviant (> 4 or) pixels with respect to a 2D polynomial fit in a small area centered at each pixel. They were then interpolated across in frames where a smooth version of the image was previously subtracted. This allows for a much easier identification of the ion hits and reduces the interpolation order in areas close to emission lines, the galaxy center, H II regions or unresolved sources. The smoothed frame was added back after the procedure. 4.3.4 Sky Subtraction Accurate sky subtraction is crucial for the estimation of line strengths and their gra-dients. The measurement of line indices well into the galaxy disk requires an exquisite handling of the night sky intensity that is typically many orders of magnitude brighter than the galaxy beyond one or two disk scale lengths, even in the darkest conditions. Additional contributions to the natural (galactic and extragalactic) night sky from var-ious astronomical or artificial sources can also have a dramatic effect on line indices, especially at low signal-to-noise. Incorrect sky subtraction might result in spectral features that could erroneously be ascribed to the galaxy profile; over-subtraction or Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 132 under-subtraction of the sky brightness yields absorption lines that appear respectively stronger and broader or fainter and narrower than they really are. The most straight-forward and commonly used procedure to subtract sky in longslit spectroscopy is to interpolate the sky spectrum from areas near the edges of the slit, if the object does not fill the slit, or measured away from the object in separate dithered observations. The first method removes smooth patterns of scattered light as well but, strictly speaking, it may also subtract light from the galaxy that may lie underneath the sky areas. Sky subtraction from slit edges is especially uncertain if the sky background areas are selected automatically. For large galaxies that fill the slit, the approach of bracketing galaxy observations with a pair of sky frames is more costly and practical only when short exposures are involved since the sky level and galaxy spectrum can change rapidly. The sky-galaxy-sky time sequence is a compromise between the need to collect sufficient sky counts for accurate background subtraction and minimizing sky fluctuations, especially at redder wavelengths and when observing close to twilight or dusk. The choice of dithering separation is also a compromise between measurements taken far enough away from the main target to minimize source contamination yet still close enough to map the same galaxy-sky background variations. For galaxies larger than the slit, we bracketed the 15 min galaxy exposures with 3 min sky spectra taken typically 5' away from the galaxy3. The sky measurements off the science frame or from an individual sky frame should yield reliable sky level estimates so long as the observing conditions are completely sta-ble, i.e. a dark and photometric sky, and minimal sky fluctuations. However, some of our spectra, for small and big galaxies, were taken (in queue mode) in non-photometric conditions, close to dusk or twilight, and/or in the presence of a bright Moon. Thus the sky brightness and spectrum changes considerably from observation to observation, depending on temperature, clouds, Moon relative position, etc. Table 4.4 lists obser-vation sequences and conditions for all of our pointings with column entries as follows; 3 For galaxies that fill the slit, we were able to measure and subtract the scattered light from one edge of the C C D that was not illuminated by the slit but that still collected photons scattered by the grating. For galaxies much smaller than the slit, scattered light was automatically removed in the sky background subtraction. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 133 Column(l): Galaxy name; Column(2): date of observation; Column(S): Universal Time of observation; Column(4): Indicates whether the observation is of a large galaxy (G) for which sepa-rate sky positions (S) were necessary. For smaller galaxies, the sky could be measured at the ends of the slit so no separate sky exposures were necessary (G+S); Column(5): Exposure time of observation in seconds; Column(6)\ Rough estimate of the attenuation due to cloud cover in magnitudes0.; Column(7): Airmass for each observation. To illustrate the sky level variations for each set of observations, Figure 4.5 shows the sky background in units of log 1 0 (photons/A/D ) versus the time to twilight or dusk for all of our galaxy observations. For all observations, the sky level is an average of the background intensity measured on either end of the slit, over the spatial range de-scribed above, in the 4200-5500 A range (a region free from sky lines). Each point type corresponds to a different galaxy (see legend); IC 239 was observed twice on two sepa-rate days, in both cases very close to twilight. The alternating solid(on sky)/open(on galaxy) symbols are for the largest galaxies with dithered sky pointings. The dotted lines represent fits to the dark (~ constant) and local (changing) sky levels. These fits do not include the galaxy measurements (open symbols). The two numbers next to each galaxy name are the fraction of Moon illumination and the Moon-object angu-lar separation in degrees, pMoon, respectively, close to the mean time of each galaxy observation sequence. The former was obtained from the US Naval Observatory As-tronomical Applications Department4. The altitude and azimuth of the Moon at the mean time of the observation, used to compute the Moon-object angles, were obtained from the sunmoonposn program (vl.l) of the Geoscience Australia's National Mapping Division5. The dotted lines represent fits to the constant sky level + linear Moon + exponential twilight contributions. Special care must be taken to account for the possibility of clouds and scattered 4http: //aa.usno.navy.mil/ 5 http: / / www.ga.gov.au / nmd / geodesy / astro / smpos.jsp Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 134 Galaxy Date meanUT Obj ExpT Cloud A t t a Airmass name (yr-mm-dd) (hh:mm:ss) G / S (sec) (mag) IC 239 2002-08-30 14:23:51 S 180 -0.55 1.059 14:36:07 G 900 -0.55 1.062 14:47:36 S 180 -0.57 1.064 14:59:06 G 900 -0.57 1.074 15:10:35 S 180 -0.57 1.078 15:22:06 G 900 -0.62 1.094 15:33:42 S 180 -0.63 1.101 IC 239 2002-09-01 14:30:44 S 180 -0.55 1.061 14:42:17 G 900 -0.45 1.068 14:53:48 S 180 -0.55 1.071 NGC7741 2002-09-02 12:10:08 S 180 -0.60 1.029 12:21:37 G 900 -0.60 1.048 12:33:05 S 180 -0.50 1.055 12:44:35 G 900 -0.59 1.082 12:56:11 S 180 -0.57 1.091 13:07:40 G 900 -0.57 1.128 13:19:07 S 180 -0.57 1.141 N G C 628 13:51:20 S 180 -0.58 1.018 14:02:53 G 900 -0.43 1.034 14:14:22 S 180 -0.50 1.040 14:25:52 G 900 -0.45 1.064 14:37:28 S 180 -0.52 1.073 14:49:00 G 900 -0.50 1.107 15:00:28 S 180 -0.49 1.119 NGC7490 2002-09-05 11:06:08 G+S 1080 -0.50 1.041 11:26:31 G+S 1080 -0.60 1.059 11:46:54 G+S 1080 -0.40 1.085 12:07:16 G+S 1080 -0.58 1.120 N G C 173 12:55:11 G+S 1080 -0.58 1.093 13:15:34 G+S 1080 -0.60 1.126 13:35:55 G+S 1080 -0.53 1.170 UGC2124 14:19:31 G+S 900 -0.55 1.079 14:36:57 G+S 900 -0.00 1.112 14:54:17 G+S 900 -0.25 1.118 NGC7495 2002-09-14 06:49:32 G+S 1080 -0.73 1.376 07:10:35 G+S 1080 -0.75 1.277 07:30:54 G+S 1080 -0.78 1.202 07:51:17 G+S 1080 -0.76 1.143 NGC7610 2002-09-15 07:45:50 G+S 900 -0.61 1.189 08:03:10 G+S 900 -0.60 1.141 08:20:30 G+S 900 -0.59 1.103 "derived from SkyProbe attenuations (mean magnitude differences of ~100-200 stars in the C H F T F O V measured in each image) for each night (http://www.cfht.hawaii.edu/Instruments/Elixir/skyprobe/). Table 4.4: Observation sequences and conditions. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 135 2.5 GO P o o 1. o 0.5 h • x. N7490 NO 173 U2124 • N7741 • N0628 x N7495 * N7610 A • 10239 10239 (4%, 124.5°) (4%, 132.2°) (4%, 106.3°) (26%, 90.7°) (26%, 71.7°) (58%, 86.5°) (68%, 75.2°) (36%, 42.3°) (56%, 26.4°) 6 X-. X •x -200 -100 0 100 Time to twilight (min) 200 Figure 4.5: Sky brightness in units of log10(photons/A/D ) as a function of the time to twilight or dusk in minutes. Each point type represents a measurement of the sky brightness for a given telescope pointing. Open symbols represent "on galaxy" pointings for the large galaxies whose sky levels were measured 5' away from the galaxy (solid symbols). For the smaller galaxies that did not fill the slit, the sky level was interpolated from the slit edges for each integration. The dotted lines are a two-component fit to the independent sky values (with a linear contribution from the ambient sky and an exponential contribution from scattering of Sun and/or Moon light (the open "galaxy" weights were set to zero in those fits). Each galaxy name is followed by the percent Moon illumination and the Moon-object angle. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 136 light by the atmosphere6. For a few of our galaxies (IC 239, NGC 628, UGC 2124) we attempted to model the contribution of the Moon and Sun light to the rapidly changing sky levels in order to recover to true galaxian profile. Scattering of Sun or Moon light by the atmosphere is both wavelength and angle dependent. Rayleigh scattering, the scattering of light off of air molecules and tiny particles (e.g. atmospheric gases such as oxygen and nitrogen), occurs for particles of size < l/10th the wavelength of the light. It is strongly wavelength dependent (A - 4 ) a has a (1+cos2 9) angular dependence. Rayleigh scattering thus dominates when looking at large angles from the light source (Sun or Moon). For particle sizes > 1 wavelength (aerosols of pollen, dust, smoke, water droplets, and other particles in the lower portion of the atmosphere), Mie scattering predominates. The A dependence of Mie scattering is much weaker than that of Rayleigh; the former also produces an asymmetric scattering pattern reminiscent of an antenna lobe. Larger scatterers produce a stronger and more focused scattering beam. Because of the forward scattering efficiency, Mie scattering dominates when looking in the direction of the light source. Hence, the closer we get to dusk or twilight (or the smallest Moon angle), the more dominated the sky background light is by the gray Mie scattering of sunlight (or moonlight). The extinction coefficient, A;(A), takes into account the scattering attenuation from the direct line-of-sight of the observation, and there is a scattering component which depends on the separation angle of the Moon(Sun)-object. Finally, there is also contribution from Ozone which is dominant at < 3500 A with a small peak between ~ 5000 - 6000 A and was observed in one of our observations (IC 239, the one closest to twilight). Hence, the contribution of the Moon to the sky brightness is a function of the phase of the Moon, the altitude (or zenith distance/angle) of the Moon, the altitude of the sky position, the Moon-sky position angular separation, and the atmospheric extinction (airmass). Krisciunas & Schaefer (1991) derived a model of the scattered moonlight as a function of the above five variables. The same derivation is true for the Sun, except 6 We assessed the amount of cloud extinction on every science frame by adding up the light of each galaxy in a fixed aperture and comparing the galaxy intensities as well as the attenuations listed in Table 4.4. Consecutive galaxy observations could vary by ~ 5-10% . Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 137 there is no dependence on Sun phase. The properties and physics of the changes in sky brightness and spectrum were deeply investigated with our data (following an approach similar to that of Krisciunas & Schaefer 1991), and could basically be understood in terms of the scattering prop-erties of the Earth's atmosphere as described above. But a quantitative modeling and characterization could not be achieved (too few data points for too many free parame-ters) that was adequate for the level required in this project and we opted for empirical controlled sky subtraction methods described below. We developed an empirical procedure to perform the sky subtraction as follows: first, we co-added the 3-4 independent aligned spectra for each galaxy(+sky); this yields a high S/N spectrum with which to determine the intrinsic shape of the galaxy spectral profile. The sky-illumination pattern (constant along emission lines, and slightly curved for the continuum) was scaled to extract the sky intensity and spectrum from the slit edges (last arcmin and at least 4-5 disk scales away from galaxy center for the small galaxies). The small fraction of sky over-subtraction due to galaxy photons at the slit ends was estimated from the consistency with an exponential disk, a galaxy profile that varies slowly along wavelength (according to extinction) and the identification of residual galaxy features in the assumed sky. The sky spectrum of each individual observation was then derived by comparison with the sky-subtracted mean observation. In this comparison, the galaxy profile at every wavelength was allowed to scale (due to the relative differences of exposures, sky transparencies and extinctions) but not to change in shape relative to the mean template, (Gal)(Rad,A), by fitting at every wavelength a linear combination of the mean-galaxy template and the sky-illumination pattern, Skyj//(Rad,A): (Gal+Sky)(Rad,A) = (Gal)(Rad,A) * C(A) + Sk y i / /(Rad,A) * S(A) The shape of the "cloud-factor" C(A) was in all cases consistent with the shape of the extinction curve, as expected from changes in sky transparency and extinction among sequential observations. These variations in the "continuum" sky spectrum show very good consistency with variations due to scattered moonlight and extinction (see below). Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 138 Very bright lines from different ions show large and random intensity variations and are thus very difficult to subtract accurately (we take this into account in our further analysis, see §4.4.3). This method of sky subtraction, as opposed to polynomial fitting at the slit edges for individual observations, maximizes the S/N of the subtracted sky. The exposure times for off-target observations were 1/5 of on-target integrations. For gray time, these exposures resulted in tens of counts in the sky continuum, good enough for a proper sky subtraction and S/N. For dark time, these 3 min sky exposures lead to count values of ~ 5 counts in the disk, much too low and CCD features could even be seen in the sky profiles. In future observations, we will opt for longer sky exposure times, closer to 1/3 or even 1/2 of on-target to ensure more than 20 counts in the sky. Because accurate sky subtraction has such a fundamental effect on the final line in-dices measurements, we have endeavored to keep any level of subjectivity, inherent to all sky subtraction procedures, to a strict minimum. The sole assumption of our approach above is that the bright part of the disk profile is closely exponential. Sky subtractions from within the slit are good to roughly 0.5% and the nodded sky subtractions were good to ~ l - 2 % . We are now ready to start making measurements from the fully calibrated and sky-subtracted 2D spectra. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 139 4.4 Data Analysis 4.4.1 Velocities Measurement of indices requires that all galaxies are in the same zero-recessional veloc-ity scale. Radial velocities were measured by comparing the "mashed" (i.e. coadded) central 3" spectrum of each galaxy in several wavelength intervals, chosen because they enclose prominent absorption lines (Mgb, H/3, & Gband), using a twilight observation as the reference. Use of the twilight spectrum as a template is beneficial for two reasons; (1) except for a few night sky emission features, the twilight spectrum is essentially that of the Sun (in reflected light). The Sun's spectral classification is that of a G dwarf star (G2V). The spiral galaxies in this study have a large contribution to their luminosities from G dwarf stars, and therefore have spectra very similar to that of the Sun; and (2) using the twilight reference as a zero point effectively puts our velocity measurements in the Geocentric frame. We thus can expect our measurements to differ from Heliocentric values by ~ ±30 kms - 1 , accounting for the relative Sun-Earth motion at the time of observation. The least-squares fits were done in an iterative and interactive manner (some of our galaxies show emission in H/3, thus the H/3 region was not included in the fits for these galaxies). Our measured geocentric radial velocities are listed in Table 4.6, Vy s (kms _ 1 ) , along with the literature (heliocentric) values taken from NED. Indeed, our radial velocity measurements show differences from the literature values. Most are within the ~ ±30 k m s - 1 margin mentioned above, but a few have larger differences (e.g. U2124, N7490); given the inhomogeneous sources and measurement methods for in the NED database, we are not concerned with these differences. Another issue related to velocity measurements if that of the intrinsic galaxy rota-tion. We deliberately selected mostly face-on galaxies, but to assume no rotation would require exactly face-on inclinations (which are determined from photometry), and an assumption that all spiral disks are perfectly circular (which is not always the case, see Andersen et al. 2001). We therefore measured "pseudo-rotation curves" for each of our galaxies and removed this rotation from the spectra. The procedure followed was to take one side of the galaxy as a template for the other side (i.e. use it as a mirror) and Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 140 measure the relative velocity shifts between the two sides. In this sense the rotation curve is symmetric with respect to the center (determined as the peak of the spectro-scopic light profile) by construction. Our measured "pseudo-rotation curves" are shown in Figure 4.6. The rotation curves tend to rise sharply near the center, followed by a -60 -40 -20 0 20 40 60 Radius (") Figure 4.6: Measured pseudo-rotation curves for each galaxy (see text). These velocities were removed from the spectra before measurement of the indices. subtle drop until leveling off to a constant value. Since the surface brightness of spiral disks drops off exponentially, it is impossible to measure velocity shifts very far out into the disk with the current data, thus the velocities were extrapolated to the outer disk at the constant velocity where it was seen to level off (typically beyond or around 1.5 to 2 exponential disk scales). The peak rotations are in the range from 20-130 k m s - 1 while the extrapolated rotations are in the range 10-60 kms - 1 . It should be noted that at A = 5000 A, a AA = 1A shift translates into a velocity shift of 60 kms - 1 . Therefore, the shifts due to the galaxy rotation are indeed small, and typically at a Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 141 level that would not seriously affect line-index measurements. It is interesting to note that Bottinelli et al. (1990) quote H I line width measurements for all 8 of our galaxies. While we do not particularly trust such measurements for galaxies so close to face-on (see e.g. Courteau 1997), it is interesting to compare their values (which are measured from the gas) and our own (measured primarily from the stars), so we list the Bottinelli et al. 1990 rotational velocities (taken as W50/2, i.e. half the line width at 50% of the peak H I flux) in Table 4.5. The Bottinelli et al. values agree best with our peak mea-sured velocities, Vpeak, but large deviations exist. Again, we do not worry about such differences as we don't feel the low inclination H I line width measurements should be trusted. Additionally, radio H I line widths measure the rotation of the gas only, and our concern here is to remove relative velocities observed in the stars. Also note that, in most cases, we attempted to align our slit with the apparent photometric major axis of the galaxy, precisely to minimize rotation effects. Name Vext VPeak W50/2t (NED) (kms"1) (kms"1) (kms"1) N0173 35.7 24.3 128.4 149 N0628 25.2 7.7 18.4 27 U2124 0.0 11.7 24.5 81 N7490 21.8 30.8 34.1 147 N7495 19.2 58.9 75.4 95 N7610 40.5 52.8 89.8 120 N7741 47.0 12.8 58.8 96 10239 24.0 14.5 49.2 60 t Source: Bot t ine l l i et al. 1990 Table 4.5: Measured galaxy rotation velocities (see Figure 4.6 and text for measurement details). The inclinations were estimated as i — cos_1(6/a), where a and b are the galaxy semi-major and minor axes respectively taken from NED (given in Table 4.1). Vext is the extrapolated velocity and \fpeak is the maximum velocity measured (see Figure 4.6). The last column lists the H I line widths from Bottinelli et al. 1990 for comparison. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 142 4.4.2 CCD Gap Locations As mentioned in §4.2.1, there are two ~17A gaps between the 3 CCDs. The rest-frame wavelength range that these gaps span depends on the effective central wavelength of the grating, the wavelength solution and the galaxy redshift. In order to avoid having the gaps fall in regions of important line indices while staying in the optimum efficiency range of the grating (see Fig. 4.2), we requested a central wavelength of A c = 5380 A. This would place the gaps in the ranges: Gapl ~ 4909 -4926 A and Gap2 ~5835-5853 A. However, we measured the central wavelength to be between ~ 5450-5453 A, a full 70 A redder than our desired value. Hence, unfortunately, our estimates for the gaps were also off their mark, and as a result, for some galaxies, the gaps did lie in important A regions (e.g. the molecular Mg indices and, in one case, the H/3 index). This also prohibits us from measuring the age-sensitive H ^ index as we no longer have the blue wavelength coverage, and the HSp index is more likely to suffer from calibration difficulties at the CCD edges. Table 4.6 lists the galaxies, their measured redshifts and the effective wavelength regions of the 2 CCD gaps. We also must take into account the line curvature in the data, which has the effect of shifting the gaps bluewards from the center to the slit edges. The line curvature was found to be well behaved between observations and very close to parabolic in shape, thus it could be easily modeled by determining the AA from center to outer edge (r = 150"). The gap locations and the indices affected are listed for the galaxy center. If the line curvature causes a change in the list of affected indices, we list the new index list and the approximate radius at which the shift occurs. It turns out that this occurs rarely and only at a radius beyond about 120", which is the last radius we consider in the small galaxies, so the gap shifts will not have a noticeable effect on our index measurements. Also note that, for the broader molecular indices (e.g. Mgi & Mg 2), these relatively narrow gaps (which are interpolated in the data) should not cause serious problems with the index measurement, but care must still be taken (see §4.5.3). Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 143 Name yNED vy* (±) Gapl Gap2 Indices Rad (kms"1) (A) (A) affected (") N0173 4366 4306.60 (10.9) 4895.44-4913.33 5840.74-5858.26 Mg 1,Mg 2,Ti0 1 0 - - HftMgi.Mga.TiOi 120 N0628 657 629.12 ( 7.9) 4954.75-4972.92 5911.82-5929.42 Fe5015,Mgi,Mg2,NaD 0 U2124 2631 2508.81 (11.9) 4924.80-4942.80 5875.77-5893.39 Mgi,Mg2,NaD 0 N7490 6213 6147.90 ( 0.4) 4865.93-4883.71 5805.45-5822.87 R7?,Fe5782,TiOi 0 - - H£,Fe5782 120 N7495 4887 4833.22 ( 0.8) 4887.71-4905.64 5831.35-5848.80 H)8,Mg1>Mg2,TiOi 0 N7610 3554 3564.84 ( 0.1) 4905.14-4923.10 5852.78-5870.28 Mgi,Mg2,NaD 0 - Mgi,Mg2,NaDlTiOi 120 N7741 751 755.66 ( 0.1) 4952.49-4970.65 5909.12-5926.71 Fe5015,Mgi,Mg2,NaD 0 10239 903 893.26 (11.0) 4950.91-4969.06 5907.10-5924.80 Fe5015,Mgi,Mg2,NaD 0 Table 4.6: CCD Gaps: effective wavelength ranges and indices affected. Due to line cur-vature, the gap ranges move bluewards along the slit. If the list of indices changes due to these shifts, we list the new index list and the approximate radius at which the shift occurs. 4.4.3 De-Redshifted Sky Lines . There are a number of sky emission lines that are very bright and variable, and thus are extremely difficult to accurately subtract from the data. When any of these lines falls within one of the passbands of a particular index, the measurement of this index should be treated with caution, particularly in the outer, low SB, regions of the galaxy where the sky dominates. Given that the sky lines effectively have zero redshift, they get "deredshifted" when the galaxy's recessional velocity is removed from the data, and hence will lie at different wavelengths for each galaxy, so affecting different indices. In Table 4.7 we list the four most prominent sky lines in our wavelength interval, their effective deredshifted wavelength, Xeff, and the affected indices (where an index is considered to be compromised if {Xeff ± 2 A) falls within any of the three passbands), for each galaxy after deredshifting. The clearly compromised indices are removed from further analysis (see §4.5.3). Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 144 Galaxy Galaxy Galaxy Galaxy N O 173 0.01437 N 0 6 2 8 0.00210 U 2 1 2 4 0.00837 N 7 4 9 0 0.02051 Sky Line A e / / affected Ae// affected Ae// affected Ae// affected [0 I] 5577 5496.88 5565.30 5530.33 5462.63 Na 15893 5808.35 Fe5782 5880.63 NaD 5843.68 T i O i 5772.15 Fe5782 [0 I] 6300 6209.50 T i 0 2 6286.78 6247.28 T i 0 2 6170.80 [0 I]6364 6272.58 T i 0 2 6350.65 6310.74 6233.49 T i 0 2 N 7 4 9 5 0.01612 N 7 6 1 0 0.01189 N 7 7 4 1 0.00252 10239 0.00298 Sky Line affected Ae// affected Ae// affected Ae// affected [0 I]5577 5487.09 5510.68 5562.94 5560.38 Na 15893 5797.99 Fe5782 5822.93 T i O i 5878.15 NaD 5875.44 NaD [0 I]6300 6198.43 T i 0 2 6225.09 T i 0 2 6284.12 6281.23 [0 I]6364 6261.40 T i 0 2 6288.33 6347.96 6345.04 Table 4.7: List of the most prominent sky lines in our wavelength region [left column]. For each galaxy we list its measured redshift, zUs, the de-redshifted sky-lines, Xeff< and the indices potentially affected by the sky line (at Xeff)- Indices are considered to be compromised if (sky line ± 2 A) lies within any of the three passbands. 4.4.4 H II Region Masking The spiral galaxies in our sample are currently forming stars. The star formation in general traces the spiral arms. Thus, our longslit observations will undoubtedly inter-sect regions of current, active star formation. While most of the gas in galactic disks is in an un-ionized (neutral) state, in regions of intense UV radiation, such as near hot young (O, B-type) stars or after a supernova explosion, the gas can become ionized. In particular, ionized regions around young hot stars are referred to as H II regions. H II regions are dominated by emission of hydrogen recombination lines, but also recombi-nation lines of helium, and several collisionally-excited "forbidden" metal lines (where "metal" refers to any element heavier than He) such as oxygen, nitrogen, sulphur, etc... Because of the dominant emission, H II regions stand out clearly in our CCD images, particularly at the Ha line (A = 6563 A) which has the strongest line intensity. Which emission lines are present and their relative intensities depend on the physical condi-tions of the H II region (temperature, electron density, composition). Typical H II regions are extremely tenuous compared to terrestrial standards, with particle densities of order 10 to 106cm~3, allowing for the "forbidden" line transitions (i.e. long-lived Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 145 meta-stable states only seen in very low density conditions allowing enough time for the radiative transition to occur before getting collisionally de-excited). Typical H II region sizes range from 1 - 100s of pc and have temperatures of < 104 K . In Table 4.8 we list the dominant H II region emission lines and their empirical strengths relative to the H/3 line. The I(A)/I(H/3) ratios are mean values from Lee, Salzer & Melbourne (2004) observed line ratios of 14 emission-line galaxies from the K P N O International Spectroscopic Survey (KISS). Note that the physical conditions (metallicity in partic-ular) in the H II regions of these emission-line galaxies are likely different than those found in the normal star forming galaxies of our sample, so these line ratios are simply meant to be loosely indicative. Line A (A) I(A)/I(H/3) Line A (A) I(A)/I(H0) ES 4102 (0.18) He I 5876 (0.10) H 7 4340 (0.43) [Ol] 6300 (0.04) [0 III] 4363 (0.07) [N II] 6548 (0.03) He I 4472 (0.03) H a 6563 (2.81) H/3 4861 (1.00) [N II] 6583 (0.10) [O III] 4959 (1.53) [O III] 5007 (4.59) Table 4.8: Dominant H II region emission lines and typical, empirical (see text), strengths relative to H/3. Square brackets, [ ], indicate "forbidden" lines. An example of an H II region from our data is shown in Figure 4.7. Note that most of the dominant emission lines listed in Table 4.8 (and labeled on the figure) are clearly visible. A notable omission is the [O III] A4959 line as this is precicely where one of the C C D gaps lies for this galaxy. H II regions are not accounted for in SPS models, and since we are mostly interested (here) in the underlying stellar population of the galaxy, these emission line regions must not enter our stellar population analysis. Given the immense complexity of the H II regions, any attempt to model and remove their signatures (which will differ for each H II region) would be futile. The only reasonable option is to mask out these regions altogether. The procedure we developed for H II region masking is as follows (and is shown Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 146 A °< o o m l o o I LO "ST V 16 14 12 10 h 8 *>• CO o *»• (V) O a: 06 : : cn co: : O cn : : o cn CD ' '' ' J I I I • > • I i i i i I i i i j I i i i i I ' i ° m J on CO OD CO 4500 5000 5500 6000 6500 X(A) N0628 Figure 4.7: Example of an H II region in NGC 628. The spectrum was obtained by averaging together 5 pixels (1") within a prominent H II region located at ~ 100" from the galaxy center. The flux is normalized to the average flux in the A = 5100-5400 A interval (which avoids strong emission lines). The most prominent emission lines (see Table 4.8) are indicated by the dotted vertical lines. graphically in Figs 4.8-4.11) : (1) Read in 2D spectrum and its variance (2) Fit (at each row) a 3 r d order polynomial (clipping all points that deviate more than 3cr from the fit) to the spectrum on either side of the Ha line (while avoiding [N II] lines, see Table 4.8), between: (Ha - 100) < A(A) < (Ha - 30) (6462.8-6532.8) and (Ha + 30) < A(A) < (Ha + 100) (6592.8-6662.8) (3) Subtract the fit from the spectrum at each radius. (4) Mash (sum) the spectrum from (3) between (Ha ± 30 A) and average to get the Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 147 total flux (in counts) in the Ho;+[N Iljemission (see top panels of Figs 4.8-4.11). (5) Mash (sum) the variance between (Ha ± 30 A) and average (6) Take the square root of mashed & averaged variance spectrum. (7) Divide (4)/(6) to get the S/N in the (Ha ± 30 A) region (see middle panels of Figs. 4.8-4.11). (8) Identify peaks with S/N > MASK, smooth them with a "seeing" box (using the measured FWHM's given in Table 4.11) and set the region to zero in the mask (see bottom panels of Figs 4.8-4.11). However, we do not filter out the bulge region (within lrfc or the seeing FWHM in ", whichever is bigger), even if it has emission. (9) Multiply the 2D spectrum with the mask before extracting ID stpectra. After several tests using different S/N thresholds above which the region gets masked, we adopted a nominal mask of S/N(Ha+[N II]) = 2.0. This threshold, which was a compromise between masking most of the prominant emission-line regions and preserv-ing enough signal in the disk, is indicated by the dotted blue line in the middle panels of Figures 4.8-4.11. All subseqent data (spectra and index measurements) presented in this thesis have been filtered by this masking procedure. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 148 _ 500 3 400 h 2 «-» 300 C 200 h B wit-0 L > V S ^ I M H * M + c ae 2 \ W d 10 -8 -6 -4 h 2 0 h - i — i — i — i — i — i — i — i — i — i — i — i — i — i — i — i — i — i — J -NO 173 -h-r-+ o i f y ^ ^ h ^ -i I i i i i I i i i i L -100 -50 50 100 Radius (") 2^  + Q + s 2 Ui S :—i—i—i—i -u 1 1 1 1 —i i—i—i—p—r- -Li. 1 i i i i 1 i ' 1 1 1 M " L kJ i i i i —i—i—i—i—|—i—i—i—i—: N0628 j 1 1 I 1 | 1 1 1 1 | 1 A i l . 1 ' J1..L i i i kJ-1 1 1 1 | A -i i i i 1 I I 1 M i l l IV \ i i iV' M i l — I I i i -1 j 1 1 i i i i 1 i i • i i i • i i i i 1 i . i i i i -150 -100 -50 0 50 Radius (") 100 150 Figure 4.8: Graphical representation of our H a masking procedure (see text) for NGC 173 & 628. Three frames are plotted for each galaxy: the flux (in counts) in the Ha+[N Ifjemission lines [top panel], the S/N of the flux in the Ha+[N II] lines [middle panel], and the adopted mask [bottom panel]. The dotted line in the middle panel represents the S/N threshold above which the region is masked (set to zero in our mask in the bottom panel). Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 149 + a x + 6 X z " m co S 500 400 300 200 100 0 10 8 6 4 2 0 1 0 h - i 1 1 — i 1 — i 1 1 1 1 1 — | 1 1 — i 1 1 1 1 1 — i 1 r U2124 —i r r A i n i I I j i i i i-|—i-. i j i i i i i i i i i i i i i i i i i i i i_ - 1 0 0 - 5 0 0 50 Radius (") 100 + X s x as \ Ui cc 2 . I 1 I . 1 . 1 - i — i — T—1—n—* 1 1 1 i — ' — ' — 1 — — i — — -N7490 1 . i f ' 1 1 H -: i H k, i - '1 ' -ii.... 1.1 —1—1—r— l l J —1—i— \—1—\—1— 1 • w -1—H-|—I-—i—i—i— i i fl 1 1 III , . 1 , 1 1 1 1 1 1 1 - 1 0 0 - 5 0 0 50 Radius (") 100 Figure 4.9: Same as Fig. 4.8, but for UGC 2124 & 7490. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 150 -100 -50 0 50 Radius (") 100 ft 5 0 0 =3 400 300 3 200 X P 100 tim 0 TT • • • i + a 35 m CO 2 1  8 6 4 2 0 0 h n 1 ; r-- 1 1 1 1 1 1 1 1 1 1 r-N7610 -100 -50 0 50 Radius (") 100 Figure 4.10: Same as Fig. 4.8, but for NGC 7495 & NGC 7610. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 151 Figure 4.11: Same as Fig. 4.8, but for NGC 7741 & IC 239. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 152 4.4.5 ID Spectra In Figures 4.12-4.19 we present the fully calibrated and velocity (redshift + rotation) subtracted spectra at a number of radii for each galaxy. The spectra are plotted on a logarithmic scale and adjusted with an arbitrary constant to offset the spectra from each other (the constant is the same for all figures except for N7495 where it is doubled). Not all radial bins are shown for the well sampled galaxies, but the first bin and last two bins are always plotted. The spectra were mashed (radially) to a minimum S/N/A >40 in the 5050-5450 A wavelength interval (selected to avoid prominant emission lines such as [O III] AA5007, 4959, HaA6563, [N II] AA6583, 6548, and the [O I] sky line at A5577). The minimum S/N/A of 40 was selected as a compromise between accurate index measurements for reliable age and metallicity determinations, and sufficient ra-dial extent for gradient measurements. The light weighted radius and the spectrum's corresponding S/N/A are labelled at the right edge of the plot (note that the r = 0"0 bin is in reality closer to r = 0'.'04 due to our binning of O'.'2/pixel and assuming a Gaussian light profile in the galaxy center). Many of the Lick indices are marked as coloured vertical lines - solid lines delineate the central passband and dashed lines the pseudo-continua. A number of pertinent observations can be made from a visual examination of the spectra: • Improper subtraction of prominant sky emission lines (falling in and around the NaD and TiO indices, see Table 4.7) are evident at the largest radial bins for all galaxies. • There is a significant overall blueing of the SED with radius for most galaxies. The three notable exceptions are N7495, N7741, & 10239. The former two show significant emission in all 4 Balmer lines, [N II] AA6583, 6548, and [O III] A5007 in the center (where H II regions do not get masked) which could hide an under-lying redder stellar SED. N7610 also shows significant central emission and bluer SED compared to those galaxies with less obvious central emission (e.g. N0628 & U2124), but a significant (further) radial blueing is evident. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 153 • For galaxies with significant emission in the Ha+[N II] region, in some cases an emission spike superimposed on an underlying absorption can be seen in the other Balmer lines (most noticeably in H/3, but a spike can be seen even in H 7 & H.5 in extreme cases, e.g. N7495 & N7741). Given that the Balmer lines are the most sensitive age indicators, it is likely that age estimates of the underlying stellar population will be difficult to constrain in these galaxies. The most likely source of the emission is due to normal, but intense, current star formation. We can rule out significant contributions from active galactic nuclei (AGN) as the source of emission based on the low [N II] A6583/Ha and [O III] A5007/H/3 emission line ratios (see, e.g., Hao et al. 2005), but small contributions in our most severe cases, N7495, N7610, & especially N7741, cannot be entirely ruled out (but these would be restricted to Seyfert 2 or LINER AGN as the broad lines characteristic of Seyfert 1 spectra are not seen). • Clear radial variations in many Balmer and metal-line absorption line strengths can be seen in many of our galaxies, but these are best examined by looking di-rectly at the measured indices, which we turn to in §4.5.1, after a description of potential resolution effects (internal and/or instrumental) on the index measure-ments. 4.4.6 Resolution and Velocity Dispersion Effects In order to compare the measured indices to the models, the models must be degraded to the resolution of the measurement (or, alternatively, the measured indices could be corrected for resolution effects, e.g. PS02). The G A L A X E V models used for this analysis have resolution F W H M G a u s s = 3 A. Our pixel scale is 0.1" and our slit width is 2". From the twilight observations, the slit width was measured to be F W H M 6 o x c a r = 10.75 A, so we have 5.375 A/". Therefore, if the seeing is any worse than about FWHM = 0.56"(= 3 A), which turns out to be true for all cases here, the minimum resolution will be that of the seeing convolved with the instrumental profile ( F W H M G q u s s = 0.5 A). Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 154 Figure 4 .12: Spectra as a function of radius for NGC 173. See text for figu re description. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 155 Figure 4.13: Same as Fig. 4.12, but for NGC 628. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 156 Figure 4.14: Same as Fig. 4.12, but for UGC 2124. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 157 Figure 4.15: Same as Fig. 4.12, but for NGC 7490. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 158 Figure 4.16: Same as Fig. 4.12, but for NGC 7495. The constant was doubled in this plot with respect to that in all the other spectra plots. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 159 -lsuoo + ( (Y) j ) 0 I §o i Figure 4.17: Same as Fig. 4.12, but for NGC 7610. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 160 CO O OS w in o CO T—1 1—4 O d OS d d d d lO o CO CO C\} q CO d CO ^ " d CO co 1—1 CD o 1 2: • • • • • • • • • • C\2 C\2 i—i ^—i T- I i—i O O O O -jsuoo + ( (y ) i ) 0 l §o i Figure 4.18: Same as Fig. 4.12, but for NGC 7741. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 161 Figure 4.19: Same as Fig. 4.12, but for IC 239. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 162 The two extreme regimes that control resolution are either due to the spectrograph slit or the seeing, depending on whether the studied object is extended or point-like, respectively. Our measurement dispersion, a m e a s , is given in terms of the slit dispersion, asut, the instrumental dispersion, G i n s t , and the object profile dispersion, o o b j in eq. 4.3, { Fsiit * Finst for aobj » o - s l i t , ( F o b j x F s H t ) * F i n s t for a o b j ~ a s i i t , (4.3) Fobj * F i n s t for aobj <C o-suti where the x and * symbols represent multiplication and convolution, respectively, of the functions, F . Since the three sources of dispersion have different profiles (Gaussian, boxcar, exponential), we must convert from the typically measured unit for each func-tion (<j, FWHM, h) to a common dispersion unit. The relationship between FWHM and dispersion, in terms of the unit of measurement representing the function of interest are provided in Table 4.9 (e.g. F W H M G m s s = 2.3548crGauss). Function Unit Dispersion FWHM e Fraction (") (1/2-light) (win Unit) Boxcar W (12)"1/2 = 0.2887 1.0 0.25 1.0 Gaussian a 1.0 y/Sln{2) = 2.3548 0.6745 0.6826 Exp h y/2 = 1.4142 21n(2) = 1.3863 ln(2) = 0.6931 0.6321 Table 4.9: Dispersion relations for boxcar, Gaussian, and exponential functions. The first column lists the function, followed by the typical measurement unit and then the conversions, in terms of the typical unit, to o , F W H M , and the one dimensional 1/2-light radius (r 1 D). The last column lists the fraction of the total light contained within the typical unit. Literature values were not available for all of our galaxies (see Table 4.2), thus we had to perform bulge-to-disk (B/D) decompositions on the spectroscopic SB profiles. Before fitting, the spectra were mashed in the 5100-5500 A wavelength interval (avoiding emission lines) and H II regions in the disk were masked (as it is the exponential profile of the underlying stellar disk that we seek). The fits included a sum of 2 or 3 Sersic profiles (see below): bulge+disk, bulg+bar+disk, or bulge-t-disk+nucleus. All functions and derivatives were degraded by the measured seeing before fitting. The sub-arcsecond Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 163 seeing and fine sampling allowed for resolution of a small, exponential, central nucleus in two of our galaxies, with scale length rnuc < 0.2". Such nucleii have been observed in HST data of earlier type galaxies (e.g. Seigar et al. 2002; Balcells, Graham, & Peletier 2005). In performing our own B/D decompositions, we noted large differences from the literature values (only small differences were expected given the different wavelength intervals of the profiles). In particular, several of our late-type bulges were found to be non-exponential. In Paper I we found late-type bulges to be best-fit with (close to) exponential profiles, which translates to an index of n = 1 in the generalized power law profile referred to as the Sersic profile (Sersic 1968). Parameterized in terms of effective quantities, the Sersic profile is given as, l / n Ib(r) = 7e xp < -bn 1 (4.4) where re, the effective radius, encloses half the total extrapolated luminosity. Ie is the intensity at this radius and bn is chosen to ensure that F Jo Ib 2nr dr = 2 Ib 2irr dr, F Jo (4.5) so the fitted effective radii for the bulge are two-dimensional (i.e. weighted by radius). In magnitudes Eq. 4.4 translates to l / n jjLb(r) = Lie + 2.51og(e) bn - 1 - 1 (4.6) where /j,e is the effective surface brightness. It is trivial to convert from effective radius to scale length, ro, and central surface brightness, LIQ, noting that, re = (bn)nr0 He = no + 2.51og(e)6„. (4.7) (4.8) Most notably, the bulge of N7495 was best fit with a n = 4 Sersic index (i.e. a de Vaucouleurs profile), usually only found for Es, SO bulges, and only the earliest-type spiral bulges. This suggests perhaps that the Hubble type of this galaxy is much earlier than Sc. In fact, as noted in NED, the original UGC catalogue does refer to this as an Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 164 "early spiral" (Nilson 1973). Our only Sa bulge, U2124, is indeed better fit wi th n = 2.5 than an exponential profile, consistent with being an early-type bulge. The only other bulge with a high Sersic index for its Hubble type is N7490 with n = 2.35, the remaining bulges are al l fully consistent with being late-type nearly-exponential bulges. Whi le we cannot explain the discrepancies between our values and those found in the literature (although difference in scale lengths between different authors of order 10% are typical; see Paper I), we opt to use our measured values for consistency in the remainder of this analysis. Bo th the literature and our measured values are listed in Table 4.10. Literature values Spectroscopic values Name Hubble Seeing re Sersic r<t re Sersic rd Notes NGC Type (") (") n (") (") n (") N0173 SA(rs)c 0.70 — 1.00 19.2 3.3(0.6) 1.65 18.9 non-exp. bulge N0628 SA(s)c 0.85 19.5 1.35 73.1 11.3(0.3) 1.20 70.6 + a < 0.1" nucleus U2124 SB(r)a 0.81 2.7 1.00 21.3 12.5(0.2) 2.50 34.0 + a 21" bar N7490 Sbc 0.74 1.9 1.00 15.7 5.3(0.2) 2.35 21.6 non-exp. bulge N7495 SAB(s)c 0.71 1.2 1.00 13.8 4.8(1.0) 4.00 18.4 deVauc. bulge N7610 Scd 0.70 2.1 1.00 13.2 1.5(0.1) 1.00 19.9 normal exp. bulge N7741 SB(s)cd 0.80 7.6 1.00 49.2 7.2(0.3) 1.00 50.0 + a broad bckgnd 10239 SAB(rs)cd 0.92 — 1.00 — 11.9(0.5) 1.00 37.6 + a < 0.1" nucleus Table 4.10: Bulge and disk scale parameters: literature K-band values compared with our own spectroscopic profile fits. Fits were performed on the spectroscopic profiles between A = 5100-5500 A (avoiding emission lines) and masking H II regions outside of the bulge region. Errors on the disk scale lengths are 10% and bulge effective radius errors are in brackets next to the measured re. See Table 4.2 for literature references. Table 4.11 lists our galaxy r e ' s , their respective scale lengths (converted using eq. 4.7), and their resulting F W H M values (in " and A). We also list the seeing values of the combined galaxy observations, measured by fitting Gaussian profiles to resolved H II regions in each galaxy. The last column indicates the "resolution regime" for the central ~ 2 r e for each galaxy (outside this region the slit width wi l l dominate the resolution). One other broadening effect, independent of the instrumental setup, that must be considered when measuring narrow spectral features is the galaxy velocity dispersion, ovd. Galaxy disks are dynamically cold systems, flattened and supported by rotation, Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 165 Name Sersic r 2 D e hb F W H M F W H M seeing seeing Res" nt t (") (") (") (A) (") (A) N0173 1.65 6.04 3.3 0.55 0.76 4.07 0.70 3.8 obj N0628 1.2 2.40 11.3 4.70 6.52 35.05 0.82 4.4 slit U2124 2.5 47.15 12.5 0.26 0.37 1.98 0.81 4.4 seeing N7490 . 2.35 32.02 5.3 0.17 0.23 1.23 0.74 4.0 seeing N7495 4.0 3459.48 4.8 0.001 0.002 0.01 0.71 3.9 seeing N7610 1.0 1.678 1.5 0.89 1.24 6.66 0.70 3.8 obj N7741 1.0 1.678 7.2 4.29 5.95 31.97 0.80 4.3 slit 10239 1.0 1.678 11.9 7.09 9.83 52.83 0.92 4.9 slit t To convert between re and hb for the Sersic profile, use hb = re/(bn)n. Table 4.11: Bulge FWHMs and seeing values for composite galaxy observations. and thus have velocity dispersions too small (of order ~ 10-20 k m s - 1 ) to cause any effect on our index measurements. However, the spheroidal bulges of spiral galaxies can have significant support from random motions. Proctor &: Sansom (2002) measured velocity dispersions for late-type spiral bulges in the range 50-200 k m s - 1 for Hubble types Sa-Sbc. Because our galaxies are mostly of later type (Sbc-Scd, with one Sa bulge), we can expect bulge velocity dispersions closer the small end of this range. Measurement of such small velocity dispersions is a difficult task with the current data (compromised by our 2" slit width) that will be attempted in a future study, but for the current study, we are principally concerned with any effect the velocity dispersion can have on the index measurements. To this end, we considered using the ratio of two indices that measure the same spectral feature, with one index having significant sensitivity to resolution (i.e. a narrow index) while the other has virtually no dependence on resolution (a broad index). The ratio of these two indices would then give an indication of any resolution gradient present in the data. The Mgb and Mg2 are the only two indices that, on first look, appear to satisfy the above criterion. In order to gauge the magnitude of this ratio for velocity dispersions expected in our spiral bulges, new models were constructed applying a velocity dispersion avd = 150 km s - 1 smoothing (which is at the high end of the range expected for our late-type bulges) and compared Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 166 them to the model resolution of 3 A. We show the comparison in Figure 4.20. The complete degeneracy between the two indices confirms that they are indeed measuring the same spectral feature. Surprisingly, however, neither the avd = 150 k m s - 1 nor the boxcar FWHM = 10.75 A smoothings have a significant effect on the narrow Mg6 index (no effect was expected for the broad Mg2 index). A significant effect is seen only at high velocity dispersions of ovd > 250 k m s - 1 (not shown). As a result, we cannot use this index pairing to reveal resolution gradients in the inner regions of our galaxies (but we note that this technique could potentially be applied to elliptical galaxies with their high velocity dispersions). We can, however, be confident that Mg indices are not sensitive to bulge velocity dispersion. Figure 4.20 shows resolution effects for a number of other indices (H/3, (Fe), and H^F), plotted against the unaffected Mg6 to isolate any effects. The indices most severely affected by resolution are the Fe, Ca (not shown), and H/3 indices. Unfortunately, there are currently no broad indices defined that are primarily sensitive to Fe abundance, and a mixture of broad Mg and narrow Fe indices would be dangerous due to the possibility of abundance ratio gradients. The higher order Balmer-line indices only show weak sensitivity at the oldest age/highest metallicity region of the grids, and none of our galaxies lie in this region of the grids. Of particular note, resolution effects could mimic an [cc/Fe] enhancement over a broad range of ages and metallicities (see bottom left panel in Fig. 4.20). The overall conclusion is that resolution (whether it be internal to the galaxy or external from instrumental and/or seeing) will not play a major role in our index mea-surements. Thus we compare our indices using Only the boxcar-convolved models. The only possible exception would be in the case of exquisite seeing, a small bulge (pos-sibility for U2124, N7490, & N7495), and a small velocity dispersion, and this would only affect the central points. We do however consider age and metallicity fits that use only the indices least affected by resolution effects, so even this extreme case will be accounted for. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 167 0.3 |P 0.2 0.1 I 1 I ' I ' I 1 GALAX V 3A G A L A X E V b o x . i o . 7 5 J l 1 i 1 i : i 1 i 1 i 1 r 20.0 13.0 GALAXEV vd-150 km/s SSP A (Gyr) 0.05 0.008 0.004 Z 0.0004 0.0001 I . I i I . I . I 0.5 1 1.5 2 2.5 3 3.5 4 4.5 5 5.5 Mgb 4.5 4 3.5 3 A 0) 2.5 fc, V 2 1.5 1 0.5 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 I 1 I 1 I 20.0 -- GALAX V3A 13.0 1 0 . 0 ^ 0 . 0 5 ' G A L A X E V box-io75A 5.0 - GALAXEV v d. 1 5 0 k m /„ 2 o ^ SSP A (Gyr) ^^Z?^^ -z J jJBftVOOOl "r . I . I . I . I . I . I . 1 . 1 , 1 , 1 " 0.5 1 1.5 2 2.5 3 3.5 4 4.5 5 5.5 Mgb 5.5 5 4.5 4 3.5 GALAX V3X G A L A X E V D o x - i o . 7 5 A GALAXEVvd_1501tm/. 20.0 13.1 H/3 5 4.5 4 3.5 I 1 1 ' I 1 1 ' I ' 1 1 I GALAX V 3 i G A L A X E V b o x . i o . 7 5 A GALAXEV„ 0.05H _L_ J _ 0 -1 Figure 4.20: Comparison of selected index-index model grids at three different resolutions: (1) the intrinsic model resolution of 3 A with iso-Z tracks as solid lines with colour scheme: Z = 0.05 (red), 0.02 ( Z 0 ) (yellow), 0.008 (green), 0.004 (light blue), 0.0004 (blue), & 0.0001 (purple), and iso-age tracks as dashed black lines at ages of 1, 2, 5, 10, 13, & 20 Gyr (from left to right); (2) a boxcar smoothing with FWHiM ;^,; = 10.75 A to mimic our slit (pink lines - lighter and solid for iso-Z tracks, darker and dashed for iso-age tracks); and (3) velocity dispersion broadening with ovi = 150 k m s - 1 to mimic the largest velocity dispersion expected for our late-type bulges (gray curves - darker and solid for iso-Z tracks, lighter and dashed for iso-age tracks). Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 168 4.5 Results from Indices 4.5.1 Line-Strength Profiles In Figures 4.21-4.28 we present the radial index profiles for all measured indices for our 8 galaxies. The indices are measured using eqs. 3.1-3.4. The integrals are computed by adding all the spectral bins (or pixels) that fit completely within the bandpass and fractional pixels are interpolated (taking into account that the spectra are binned, not sampled). Errors on the indices expected from random noise (photon + read-out noise) are computed from the variance images that were propagated through the reduction process (§4.3.2). Variance spectra, <J^ (A), are extracted from the variance images in an analogous fashion as the galaxian spectra are extracted from the data images (radial binning and H II masking). This variance image is then used to estimate the error on atomic, o~n(lEw), a n d molecular indices, crn(Imag), expected from the random noise as: 0n(lEw) — ~nT~ and where + °FH cO A rO V O + 0~FR F cO 2.5 log10(e) (7n{-Lmag) ~ A A l n _ n i r _ „ ^n^EW), AA C lQ-O-tlmag A c 0 = (Ac2 - Aci)/2, Fco = F(Xc0), />A C2 = / F(X) dX, J\cl Fr = Fpb Abo — A c 0 A r0 — ^ 60 1/2 (4.9) (4.10) (4.11) JXpbl c r ( A ) 2 where pb stands for "passband" (c, b, or r) and all other variables are defined in §3.2. The error in any given bandpass is uncorrelated with the errors in the other two band-passes, thus no covariance terms appear in the above equations. All atomic indices with a measured index < 0 (except for Balmer indices) were set to -99.99 for plotting in the log (except for the higher order Balmer indices, a negative index Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 169 value is most likely caused by emission fill in or poorly subtracted sky lines). These plots are meant to serve as guides primarily for relative trends (i.e. gradients) in the indices as a function of radius (the indices will be compared to the SPS models in the following section). To facilitate the visual interpretation, we have plotted all indices in a loga-rithmic scale (straight logio for the atomic indices, magnitudes for molecular indices). All indices are normalized to their average value between r = 0-0.5" (for molecular in-dices we normalized the lQmdex values), thus we plot logi0(index) — C for atomic indices and (index) — C for molecular indices, where C = logi0((index(r = 0.0-0.5"))). This places the central values close to zero, so rough estimates of radial variations are easily obtained. If desired, to convert to the measured log10(index), simply add the value of C printed (in gray) at the bottom left corner of each panel to left axis values. The axis on the right represents the measured index value, thus the measured index (in A for atomic indices and mags for the molecular indices), can be read directly from the right axis label. The axis scale is kept the same in all figures (1.0 dex for atomic indices and 0.4 for molecular indices to account for the factor of 2.5 in the magnitude definition) to facilitate observation of relative trends. The exception to the above scheme is for the Balmer line indices. Because of the complexity of the pseudo-continua in the blue region of the spectrum, the higher-order Balmer-line indices can have negative values within the model SSPs. They also straddle zero, hence a logarithmic scale is not suitable. Thus for all Balmer line indices (including H/3 for comparison with the higher-order Balmer lines) we plot them in their measured linear scale (in A). A number of general observations can be made from the index versus radius plots in Figures 4.21-4.28. For the galaxies with good radial sampling and without significant emission (N0173, N0628, U2124, & N7490), the gradients are very symmetric between the two sides of the galaxy. There are small features that do not precicely agree on both sides, but this would not be expected as they are not mirror images of each other (i.e. galaxies are not azimuthally symmetric). These features tend to coincide with visible spiral arms crossing our slit at the same radius (see optical images and slit placement in Figures 4.3 & 4.4). Global negative gradients from the central region to ~ l r d are Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 170 evident in most metal lines of order Aindex ~ —0.1 to —0.4dex. On the other hand, the Balmer lines tend to increase with radius, with A(IIJA) ~ +4 to +7 A. Note that any gradient will be smeared out by the seeing within the seeing disk, so the "flatness" of the inner few points (the measured seeing for each galaxy is denoted by the light blue dashed vertical lines) should not be taken as the true inner profile. Also of note is that our only barred galaxy, U2124, for which we aligned the slit along the bar, shows the shallowest gradients in essentially all of the indices out to the edge of the bar. This would be consistent with strong radial mixing induced by the bar which would weaken any pre-existing gradients. In general, the trends in indices which measure the same spectral features (e.g. Mg! vs. Mg 2 vs. Mgb, the Fe indices vs. each other) are consistent with each other. There are glaring inconsistencies in a few cases, e.g. TiOi & T i 0 2 in N0173, but these can often be explained by systematic issues with the index measurements, e.g. for N0173, TiOi is affected by the gap in wavelength coverage and T i 0 2 is affected by a prominant sky line that is very difficult to model and subtract accurately. For all galaxies, the H/3 index does not appear to follow the higher-order Balmer indices. It tends to remain relatively flat while the others are increasing, and for the strong emission-line galaxies, it dissapears off the plot limits altogether. This is a clear sign of the more prominant emission-line filling in the H/3 line compared with the higher-order Balmer lines (we return to this issue in §4.5.3). Naively, these gradients suggest that, as you go out from the center, galaxies become more metal-poor and younger, as one might expect for an "inside-out" type of formation. These trends are quite different than those observed by Fisher, Franx, & Illingworth (1996) for SO galaxies, where the Mg 2 gradients were seen to be steep in the bulge, but flatten outside the bulge region where the disk begins to predominate. They do, however, also find very shallow H/3 gradients (as is also observed in elliptical galaxies, e.g. Gonzalez 1993), but the higer-order Balmer lines in our study suggest that the H/3 gradient gets flattened by emission fill-in. Note, however, that in the context of the SPS models, a constant H/3 profile in the presence of falling metal-line indices is not Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 171 consistent with a constant-age (this would require a rising H/3 profile, see, e.g. Fig. 4.29). Rather, it implies that the central regions are younger compared to the outer parts. General trends are not as easily discernable for the strong emission line and/or poorly radially sampled galaxies (N7495, N7610, N7741, 10239), but it is clear that some of the indices will be unusable for further analysis of the ages and metallicities of these galaxies (i.e. those suffering from strong emission-line contamination). This point is further explored in the next section where we compare the index measurements to those predicted by SPS models. Given the overall symmetry in index measurements between both sides of the slit for our galaxies, for the remainder of the analysis we coadd both sides of the slit in efforts to maximize our S/N. A discussion of trends in individual galaxies is reserved for section §4.6 where we synthesize all of the different analyses (spectra, index gradients, age and Z fits from different combinations of indices and, where applicable, from BVRIK photometry). Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 172 i z ~ I 01 z I 01 CM fi o o I o CO o I "n CO CO a> O U I -—* in fi o g - 0 . 2 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 ||i 111 11 in 11111111111 p| I I I hC=0.513 I I I -q l l l | l l l l | l hC=0.750 l l l l l l l l l l l l l I— ^ A hC=0.654 i | i i i i | i i | i i | i i i i | i i n | t i i i i i i H i | h i | i i l i i | i i n | I I II' " J o I I I I !! 11! 11 i I I— <-&— I I I JC=0.070 I I I -l l l | l l l l | l l l l | l l l l | | l l l | l l l l l | l l l l | l l l l [ ~ I I I " I " ™ % ^ I I i -C=0.208 I I I i i | m i | i i i i | i i i i | . m | N J n l m i | M . i j i I i -I I I C=0.559 I I I "• .I....I....I....IL...I..I..I — i — 4*: 4 2 4.8 0.6 6 6 4 H 2 ilfa H o.i 0 -0.1 0.3 0.2 0.1 0 6 4 2 0.8 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S /N / A £ 40 Mask=2.0 l o I io( R / R d 8 ) - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 NO 1 7 3 - SA(rs)c Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 173 i o" CM lO 4> &• . O u I CO 00 m o 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 IP 1 1 1 1 1 qi 1 11 111 111111 : .s - I I i -14 Q — ^ A- * » * -0.2 --0.4 -•°-6 K , -H-H 0.2 -0 - * * — - * • -0.2 --0.4 -.2 -0.6 h o i »" o lO ii O U I o r-tQ g o I 00 t>-V o o I n (fl . z O 0.2 -0 --0.2 -0.4 -0.6 0.2 -0.6 C=0.460 iiii| i in|ii|i i | i i i I T i f C=0.432 iiii|iiii|iiii||iii|ii|ii|iiii|iiTi"|: C=0.236 I 4 I . 0.8 I I I 0.6 | 0.4 I \ I r.i hdi = 0.8 | 0.6 - 4 - 2 H2 8 0.6 i i i I i i i i I i i i i I i i i i I p i i i I i IIi i I i i i i I i i i i u 0.1 h I I g-o.i h -0.2 h 0.1 0 h < - A — A g-o -o h -0.2 h <^  - 2 -= _ 4 -3 -2 1 1 -0 --1 -- 2 L 3 -2 -1 -0 -i -I o.i i i •j C=0.029 I I iiii|iiii|iiii|tiii|ii|ii|iiii|irTrti Ho -o.i C=0.073 Mi|i i i i | i i i i | i i i i | i i i i | i i l i i | i in|i , i i i | 4 -3 -2 -1 1 0 r _ 1 Ji i lni i l 2 -0 -A J i l i i i i ltii i l i i l i i l i i i i lni A -i -i i i | i i i i |n|i i | i i i i|ii i iJ1 ^...•.-*.xi ^ -1 Miiiiii|iiiiiiiiiitiiiiii[iiiiiii|inir I V I A I I I M I I I I • IL . . . I . . I . . I . I ' • • I • • • -T -0.2 0.2 0.1 0 -0.1 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S / N / A * 4 0 Mask=2.0 l o g i o ( R / R d " ) - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 N0173- SA(rs)c Figure 4.21: All measured indices as a function of logi 0(R/R^ s) , where R U s is our measured spectroscopic disk scale length, for NGC 173. All atomic indices are plotted in the log minus a normalization constant, C, indicated in the lower left corner of each panel (see text), except for the Balmer lines (which straddle zero). Axis labels on the right (in gray) are the measured scale (in A for atomic and mags for molecular indices). The blue upright and pink downward triangles are for opposite sides of the galaxy. Vertical dashed lines indicate the location of the seeing disk (light blue), bulge effective radius (red), and disk scale length (green). Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 174 -3-2.5-2-1.5-1-0.5 0 0.5 S / N / J U 4 0 Mask=2.0 l o g i o ( R / R d S ) -3-2.5-2-1.5-1-0.5 0 0.5 N0628- SA(s)c Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 175 u i CM lO tu M O U I S" CO co m • <D Et. . o I CO o m u O a i aT o in ' U I .—^ CM CO f-m CO E L . . ac O o I 3* a . z O Fi I I | I I I I | I I 11 |E I M | I I I I | I I I | M I 111 I I I J . -0 .6 rxul 0.1 h Y 0 1-$-0.1 \ 0.1 Y O K 01 p-0.1 -0 .2 4 3 I I C=0.061 I [lll|llll|llllHl I 2 le*.. 1 0 - 1 2 0 £ -21?* -4 - 6 it 3 2 1 0 -1 - 2 E 1111111111 |P 1111 I 1 1 1 1 1 1 1 1 1 1 1 1 U, 1 r i 7- o.i i C=0.014 I I I i i llll|l|ll I 4 i l i i i i h l n l Li I A I i & -r -P"II|IIII|IIIIHIII|IIII|I|II|IIII|^I|-I -- 1 1 0 -0.1 -0 .2 0.2 0.1 0 -0.1 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S /N / J U 4 0 Mask=2.0 l o g i o ( R / R d 8 ) - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 N0628- SA(s)c Figure 4.22: Same as Fig. 4.21, but for NGC 628. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 176 u i SK u i N 2 O u I £» CO CM Ti-cs o I — O o n •<f o u I CO CO CO 3 lO es CJ o 0 -0.1 -0.2 0 -0.1 -0.2 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 TTTTTT l " " l " " l l " " I N ' l . l " " . q 0 1 i T i I i C=0.038 I l l | l l l l | l l l l | l l l | | l I leL H -o.i 5 r e l h / " r : _ 0 . 2 o*"0-6 A . I I C=0.063 I g f a f t H o.i HcJ I I. I i i H -o.i i i - o hC=0.066 I H4^ i i i i | i i i | | i i i i | i i i i , !Hi| i i ir j g A i iiili<iwiiiiTtfri' ' V ~ i 111 i i\ hC=0.757 I I I l l l | l l l l | l l l l | l l l i | l l l l | l l l l | |M hC=0.100 I 1 1 1 1 ' ' ' • 1 1 1 1 1 1 • • • JL. 0 2-0-2 - 0 .2 * o i —^ o m 4> o o i s o I •c o -0 .4 X o CO S -0 -2 3^ o 111 j • • • 11 • 111 [ 111 r (• • • • | ii' i [ P 111 j 1111 j-^ i 11 hC=0.521 I I I I -j l l l | l l l l | l l l l | l l l i | l l l l | l l l l | | l l | l | l l l ( t U ^ * - * I -J"" -0.4 -0.6 0 -0.2 -0.4 -0.6 0.1 0 -0.1 -0.2 0.1 0 -0.1 -0.2 0.2 0 -0.2 -0.4 -0.6 I hC=0.781 I ***** I I I I I I I I I : H 2 | i i i i | i i i i ; i i iHii i i | ini| i i i i i l i i i i t 4 6 4 I I I I I I I - | 2 -C=0.733 I I I I I i 11111111111111 11111111111 |i 111J1111 I < r & r - Ho.i II H 0.2 i II Ho C=0.113 I I I I -0.1 I l l l l | l l l l | l l l i | l l l l | l l l l | t l | l | l l l l | o.4 I I- <rA-111 11 ^C=0.260 l l l l l l l l l l l l l j I -J I I I "I 4 I '! 3 i i l U C=0.630 I I I • 7 i i l , i i , l , , , i i , , , U , , . . i . . i , [ | , I 4 2 A.8 0.6 8 6 4 H 0.3 0.2 0.1 8 6 4 2 1 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S / N / J U 4 0 Mask=2.0 l o g i o ( R / R d S ) - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 U2124- SB(r)a Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 177 CJ i o" N o O CJ I in co co & CJ i «" o to v o CJ I .—* OS o CJ I N -CO V M o CJ 1 s S3 :z O 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 I I I | I I I I | I I I I | I I I | | I I I I | I I I I | P I | I | I I I I M I I I I H C=0.468 Ml l l l l l l l l l l ff>.f i | |n i i | in i | | i i l i | i i i i ^ i M I * a * u ^ ^ i i i ^ T ^ t j i i i KA C=0.418 I Ml | l l l l l l l l l | l l l t | l l l l | l l l l | | l I C=0.251 I i i i | i i i i | i i i i |ni i | i |i 1111111 if h C=-0.020 Tll|llll|llll|lllt|llll|llll|lll|l|lll|-[- 0.2 - 0.6 - 0.4 C=-0.096 W I I I I "J I t 4 2 4.8 0.6 4 4.8 0.6 4.8 0.6 0.4 = 4.8 H 0.6 - 0.4 0.1 -0 -u l P-C . -- 0 .2 0.1 0 o i o" £ - 0 . 1 -0 .2 4 3 2 1 0 -1 2 0 - 4 3 2 1 0 -1 - 2 b. X <5l ac T 11111 [ I 11 p 11 I 1111 I 11 I |l l M 11 I I I M I I I I -J I I V l C=0.030 I I I * l l l l | l l l l | l l l t | l l l l | l l l l | | l l | l I I I I I I I I j i i | i i i i | i i i i | i i i t | i i i i | i in|hl i f i i i iT 11 I \]u I KI -! | * H IIII|IIIHIIII|IIII||II1I|IIMJ I I l l I I •rf i i i i i : i 111 • I — I... • I • • .t I — I — IL. I • I —7 0.1 0 -0.1 -0 .2 0.2 0.1 0 -0.1 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S/N/A * 40 Mask=2.0 l o g i o ( R / R d " ) - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 U2124- SB(r)a Figure 4.23: Same as Fig. 4.21, but for UGC 2124. The major axis of the bar is marked by the purple vertical dashed line. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 178 i z~ u o I w z o o I 'p CO N <# CCS O o I o" o CO o B C o u I CO CO CO t BO o u I CO o 0 -0.1 -0.2 0 -0.1 -0.2 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 C=0.113 L <-A-. C=0.156 rrrrrrj i i 1 1 1 1 1 1 Mi 1111 I I I I I 0.2 0 I 0.1 CT - 0 h ^ v - -0.1 « - 0 . 6 l[llll|lll||llll|llllf - 0.2 - 0.1 - 0 i i i i | i i i i | i i i i j i m | i i i | | i i i i | f i t 4 ] g 0 " 1 j— K A- - - -h-C=0.156 lll|llll|llll|llll|llll|lll I I I h-C=0.?62 I I I lll|llll|llll|llll[llll|lll||llll|ll fN 1 - 2 S-o-2 u i in .—« o CD b. *M O i si" 2 u I 01 B0 2 I I 'Si o l l l | l l l l | l l l l | H I I | | l l l l | l l l | | l l l l | I I I I M h ^ - A - . - . - . ^ - S J ^ . ^ A M L I TI - ±1- i i i Pc=0.571 I I I lll|llll|llll|llll)llll|lll||llll[llll[j o i S" CS CO I -0 .4 0 -0.2 g>-0.6 0 -0.2 -0.4 -0.6 0.1 0 -0.1 -0.2 0.1 0 -0.1 -0.2 0.2 0 -0.2 -0.4 -0.6 i i i •C=0.949 I I • • | i m | i i i i | i i i i | i i i i | i i i | | i m l i i i i H H 6 4 • - A I "^ - f A -I 6 4 H2 hC=0.735 1111111111111 I -H-f+f I - 0.2 I i . - 0.1 I -I C=0.123 lllllllllllll I I— • •|Mii|im|Mii|.mi...a] A i i 1 I I C=0.272 [Mi|iiii|iiii|iiii||iiii|iii||iiii|[iiirj C=0.641 I I I I I n.... i... 11 : 3 I 4.8 0.6 8 ° 8 1 -0.1 0.4 0.3 0.2 0.1 8 6 4 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S / N / A £ 40 Mask=2.0 l o g i o ( R / R d 8 ) - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 N7490- Sbc Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 179 o i o CVJ l O V fc. ' oo o I '— i d CO CO lO £ o I co o l O V o u I — s 03 o I CO m o i i |i i i i | ; i n i | m ] i n i | I I I h <-A-i S" ta o -0 .6 ho. C=0.508 l l l l l l l l l l l l l • C=0.478 I I l l l | l l l l | l l i l | l l l l [ l l l l | l l l | | l l l 0.2 0 -0 .2 -0 .4 -0 .6 0.2 0 -0 .2 -0 .4 -0 .6 0.2 0 -0 .2 -0 .4 -0 .6 0.2 0 -0 .2 -0 .4 -0 .6 0.2 0 -0 .2 -0 .4 -0 .6 0.2 0 1--0 .2 h I r e ' h dt | l | l l l l | | l l l l | l l l l | l l l l | : l l l l " [ - 4 m^iA 2 i l l i ^ ^ - . J i <* ni 4. h C=0.263 I I l | l l l l | l l l l | l l l l | l l l l | l l l | | l l l h ^ - • ^ ^ I I C= -0 .000 I I lllllMIMMIIIIIIIlllHlnl I C=-0.277 I l i l l | l l l l | l l l l | l l l l | | l I • III, [MM ° - 4 l~C=0.659 - 4 - 2 r£ 0.6 mm mM 4 8 H 0.6 0.4 -H 0.6 0.4 i ts 0.6 0.4 0.2 8 8 4 2 0.1 V o P-o.i - 0 .2 0.1 V o • | I M I | I I " | " M - <-A p-0.1 -0 .2 4 3 2 1 0 -1 2 0 <v - 2 = - 4 3 2 1 0 -1 - 2 ac 3 2 1 h r 111 II 11111111111 i i C=0.042 l l l l l l l l l l l l l H o.i o -0.1 C=0.083 II|IIII1IIII|IIII||IIII|III11IIII|IIIIT I I I -•*T -I -I J | i i i i | i n | | i i i i | i i i i { i i | l l l l | l l l l | l l l l | l l l l | l l l | | l l l l | j | | | [ I I | J 0 i i i i i i i i i i i i i i i i i i II I -0 .2 0.2 0.1 0 -0.1 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S / N / A S 4 0 Mask=2.0 l o g i o ( R / R d ' ) - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 N7490- Sbc Figure 4.24: Same as Fig. 4.21, but for NGC 7490. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 180 i z ~ o I z o I CM CM T f fi o_ "So o o I o CO Tf o u I CO CO CO Tf ii EL. o o I to Tf Tf e o_ "M o 0 -0.1 -0.2 0 -0.1 -0.2 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 1 1 1 1 1 ' 1111 • 111 • • | | ••>>]>•>> 111 •A I I ~C=0.004 I i i i l i in | i i i i | i i i i | l i i i | i i i l l in i " * ' f i r e l hdl I I* -I -0.1 I -I -0 .2 +4 1 i "C=-0.013 I GO l l l l l l l l | l l l l l l l l l | i l l l | l l l t | l l l l j l l l l+ . oft- 3 |-C=-0.26(1 4 11111111111 |l 11.1 4-•4 4 H H+f 111111 I LLJjj 11 hC=0.049 I -I -I -I 1 I I -C=0.160 I i i i | iMi | i i i i | imUi i i | in h-C=-0.232 • » — • • I -J • • • • J I *••] I I I I — 11 . . . I...Ji i I I* d 0.1 co T f O o I 0.6 0.4 0.2 0.1 0.8 0.6 0.4 0.2 2 J .8 0.6 0.4 0.8 0.6 0.4 o i 1—I o « o o I s u I oj oo 3E O I ,—. 2 60 O - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S / N / A * 4 0 Mask=2.0 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 0.1 0 -0.1 -0.2 0.1 0 -0.1 -0.2 0.2 0 -0.2 -0.4 -0.6 hC=0.233 l l l l l l l l l l l hC=0.297 I • • • • I • • • •! • • • 11| ^'••.-.f^ -4' 'V. ••| l f ' ! . ' l " i | l lhii|iii{(.ii[il -C=0.341 l l l l l l l l l l l l l I— <-&----•-«' C=0.030 l l l l | l l l l | l l l l | | l l l | l l l | | l l l l | l l l l | 1C=0.088 I H 2 I 2 I -I 1 - 2 - 0.6 T 0.4 | l l l | l l l | | l l l l | l l l l | i - 0 0.2 0.1 log 1 0 (R/RJf) C=0.238 I I I I i I i i I i i i i | | i l l I I I I 11II11111111 0.8 0.6 0.4 8 -I 0.6 0.4 H -0.1 -0 .2 0.2 0 -0.1 0.8 0.6 0.4 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 N7495- SAB(s)c Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 181 i o -eu m co o u I to CO co l O g o CJ I o to CD o o I a" o ID V fa. 1 1 BO o CJ a" as Z o 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 C=0.137 00 l O cu fa. 1111111111111 I I C=0.133 i •. i i \ i | l m | ' " l | i M | M . . | ] H 2 I I H 2 I * -= 0.8 - 0.6 - 0.4 i i i i | i i n |n i i |h i i | i in | i i i i | i i i i | -C=-0.056 •I-.. I I - f t I h C=-0.339 m |hii|MH|m»Jiii<t l'-liiii||||||i'"l'"H""l""i jM -= 0.8 - 0.6 - 0.4 1 I C=-0.227 ) i I [ l l l | l l l l | l l l l | l l l l l t l l l | l l l l | l l l l | l I I l l C=0.285 h"i i I i i i i I i i i i I | I I J 0.8 0.6 0.4 0.8 0.6 0.4 0.2 0.8 H 0.6 - 0.4 - 0.2 H 0.2 H 2 ^ 0 . 8 0.6 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S /N /A . £ 40 Mask=2.0 0.1 V o P-0.1 -0 .2 0.1 ? 0 o" p-0.1 -0 .2 4 3 2 1 0 -1 2 0 -2 -4 -6 C=0.046 ac ac 11111111111 n 11111111 >>1111 C=0.013 I I .. I . 1 I I ~ I -i i i <-A----+-+-4i [Jll|llll|llll|llll|| I _ 4 L «*.,,.XXV-H+H-H I I II I I i I I i -H+H i --H+H I I i l J I -4 • • i — • — ' • 1 4 .ii...i W l log 1 0(R/Rlf) - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 N7495- SAB(s)c Figure 4.25: Same as Fig. 4.21, but for NGC 7495. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 182 CJ i 2 CJ CJ I M 2 CJ CJ I ?» CO CO CJ *<3> o 0 -0.1 -0.2 0 -0.1 -0.2 0 -0.2 -0.4 -0.6 h- <-aV : 4; r-C=-0.183 111111111111 j 1111 ||i 11111111 j 11111111 f-4-,-"v-" T i i i i i i i c=-0.090 r e i h d i -111 • • • • 111111111111 T i M i i ) ) 1)) 111 • • • <X ft i i i i C—0.065 I I I -J 111,, H , H 0.2 ^ 0 -0.1 -0 .2 - 0 .3 0 -0.1 -0 .2 - 0 .3 B 0.6 0.4 I co m 1 v EL. o CJ I CO CO CO rf 1 O CJ I lO o I D 6 EL. 4.8 0.6 0.4 4.8 0.6 0.4 0.8 0.6 0.4 0.2 0.1 O CJ I w 2 CJ i oo 2 CJ I i 2 o 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 0 -0.2 -0.4 -0.6 0.1 0 -0.1 -0.2 0.1 0 -0.1 -0.2 0.2 0 -0.2 -0.4 -0.6 I I I | I I I I | I I I I 1^,1 I I > | l | | I I I I | I I I I | I I 1 L U g -C=0.444 " l " " l " " l " " H | M i i | i i i i | i i n t ] C=0.036 .4,, I I I I I I I I I llH|MM|llll|lM I I I I I I I _C=0.103 I I I M l | l l l l | l l l l | l l l l | l l l l ) | l l l l | l l l l | l l l l | J C=0.271 7i 1111 • • l •• •, 1 • I I 1 1 1 1 1 1 1 I — I -- 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S / N / A £ 40 Mask=2.0 l o g i o ( R / R d 8 ) - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 N 7 6 1 0 - Scd Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 183 111111111111111111 pi i >p 1111111111 " ' ' -A I I C=0.013 C=0.247 r e i h d I I •Mil •tj H2 H+h H 2 * -J C=0.185 I I l l l | l l l l | l l l l | l l l l | | l l l t | l l l l | l l l l I I 0.8 0.6 0.4 0.8 0.6 0.4 2 6.8 0.6 0.4 0.1 T* 0 g-o.i - 0 .2 0.1 V o p-0.1 -0 .2 4 3 2 1 iX 3C u. C Q . as 0 -1 2 0 - 2 - 4 - 6 3 2 1 0 -1 - 2 3 2 1 0 •-•2.. I I I 111111111111111 ||l 11 f 1111 I | I I I 11 I I 11 ( J I I -I I I I I JC=0.021 I I I i i i | i i i i | i i i i | i i i i l l i i i i l i in , [ .n i i | i i i i f l C=0.055 I I •<*•*+ I I I I I I I I I -I Mn n i l I l " • t T : : " ^ I I I I K 4 3 + B + ••••::.l:t-::.-s: i i i i i i I +++ I . I ! I m l / 1 I I I I I r i i I I I I I III n i ' l • I I I I I I I I 4 0.1 0 -0.1 5ft8 0.1 0 -0.1 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S/N /A £ 40 Mask=2.0 l o g i o ( R / R d S ) N7610- Scd Figure 4.26: Same as Fig. 4.21, but for NGC 7610. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 184 -3-2.5-2-1.5-1-0.5 0 0.5 S/N/JU40 Mask=2.0 l o g i o ( R / R d 8 ) -3-2.5-2-1.5-N7741 -1-0.5 0 0.5 • SB(s)cd Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 185 i o" ;> cu o CJ i CO co uo & 8? CD O T UO CU -~ o u I a" o CJ I CO r-cu o cj I S" z o 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 • • 11111 j 11 • • I • C=0.139 C=-0.067 IIIIIIIIIIIIII h C=-0.220 — i B C=-0.425 | l l l | l l l l | l l l l | l rc=0.125 I I I I I i I i i i i I i II | H I I | 1,1 I I l ^ ' W 1 1 1 | A •. A J - - * * J* ' I r e ' h d l I f +•••1' I N. - 4.8 iii - 0 6 II: - 0.4 IrUf ? ^i i i l i l i i i l l f i ' l l l i i i i j i 2 11111111 ill 1111111|| MI ij. lr JUL Titr i i l i in l i i i i l i i i ' i i i i in j j 11 i i • •' — i .i... i... .II 4.8 0.6 0.4 0.2 4.8 0.6 0.4 H 0.2 0.6 0.4 H 0.2 8:4» 2 1 0.8 0.6 0.4 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S / N / A £ 40 Mask=2.0 0.1 0 CJ i-o.1 - 0 .2 0.1 0 cj i o" p-0.1 -0 .2 4 3 2 1 0 -1 2 0 - 4 - 6 3 2 £ 1 X 0 -1 - 2 x * 1 1111111111111 j n 11111111 I I I I | | I I i " I I I I l I II C=0.007 I I II l l | l l l l l l l l l l l | l l | l l l l | l | l l l | l l l l l | I I C=0.042 I I II I II I II I II Mil MM I l i Mi l l I'll l i l . <-A-. •Al 1 A I I 7 II V J mn 11j1111IIIiH11 . & - A A . . Vl A I * 3 1 I 1111 1111 I i 11 I i i I I i 11 4 r<-A-• A l A A . I A l l I \ / i*A- || I I I • i | i i i i | i i i i | ih i | i i i i | i | i i i [ i i i i l l i i i i f l i i i i A-1 I • I — I. • • • I • l ^ l • I ' " : ' I 'I" 1 1 A _J 0.1 0 -0.1 -0 .2 0.1 0 -0.1 -0 .2 log 1 0 (R/R d u ») - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 N7741- SB(s)cd Figure 4.27: Same as Fig. 4.21, but for NGC 7741. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 186 i z " o u I N z X o £-<4 0 •0.1 -0.2 0 -0.1 -0.2 111 j 1111 j 1111 j 1111 j i II 111 m j11p CM CM t ti "So o o I o" o co T f o I 0 -0.2 -0.4 -0.6 hC=0.031 X o CO oo CO Tf V o 0 -0.2 -0.4 -0.6 X o -<-\. io lO Tf Tf e o 0 •0.2 •0.4 -0.6 C=-0.057 l l l l l l l l l l l l l l l l I r e ' h d i "C=-0.015 I I 11111111111111111111111111111 y i o -0.2 -0.4 -0.6 hC: hC=0.607 l l l l l l l l l hc=o.oio • I — I — I • H-rf l|in4t¥'P I T h 17 . i i " " i 1 1" i • • •• o I -o.i I I -0 .2 | 111 i~H mm JL 0 -0.1 n -0 .2 -0 .2 | - 0 . 3 o - 0 .3 = 4.8 - 0.6 - 0.4 0.2 6 4 H 2 4.8 6 4 H 2 4.8 = 4.8 -j 0.6 A 0.4 0.2 o i to BO O I 2 I « Q0 a i •a 00 o -0.4 -0.6 11111111111111111111 hC=0.528 I o •0.2 -0.4 -0.6 0 -0.2 hC=0.366 I l | l l l l | l l l l | l l l | | l l l l | l l l l o to fi -0 .4 -0.6 0.1 0 -0.1 -0.2 0.1 0 -0.1 -0.2 0.2 0 -0.2 -0.4 -0.6 -I, £'l""r| •4 i 3 : I 4 •1-4 h-C=0.623 l l l l l l l l l l l l l 4.8 0.6 4 2 4. i J i -j l l i i i i l i i i i l i i i i l n i i l C=0.047 l l l l l l l l l l l l l l I I I I I I l i | l l l l | l l l l | l iu | l l l l+ I I / 1 1 C=0.143 l l l l l l l l l l l l l I'v I I I 4 | H > i I C=0.442 1 I i i i 3 111111111111111111 • • — i 0.8 0.6 8 0.1 0 -0.1 -0 .2 0.2 0.1 0 -0.1 4 0.8 I —-A 0.6 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S / N / A * 40 Mask=2.0 l o g i o ( R / R d " ) - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 10239 - SAB(rs)cd Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 187 CJ i o" ?>-CM CU Cu Q^O O o I S" co co iO 0> o rf m cu Si ac o u I J—s ca o if CJ I CO r-CU o CJ I a S3 o 0.2 0 •0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 0.2 0 -0.2 -0.4 -0.6 I C=0.291 I r . \ C=0.326 I M | M M | l l M l l M H M M | m . 111111111111111111 j 11111111111111 I rfjrH+l I C=-0.039 I i - • • i t -i C=0.003 I l l | l l l l | l l l l | l l l | | l l l 1 i C=-0.230 I LL i i | i i i i | i n i | i i i l l i C=0.250 | I h-.. I — I — I . . . 11.... I.... I • H 2 • •• |—j—I Q tti-m if H 2 = 4.8 - 0.6 • ii-H 6 111 ITI & * 3 4 4.8 0.6 1-4 0.8 0.6 0.4 ^4.8 0.6 H 0.4 4.8 0.6 0.4 0.2 H 2 - 1 - 0.8 - 0.6 0.4 0.1 Y o i-o.i - 0 .2 0.1 V o p-0.1 -0 .2 4 3 2 1 0 -1 2 0 = - 4 - 6 3 2 £ 1 X o -1 - 2 3 2 1 0 I C=0.015 I l l | l l l l | l l l l | l l l | | l I h- <-A I*a M l l l l l l l l l l l l l l l l l l l l l l l l l l ^ l l l l l l l l 0 1 X i I • i • • I • • • • I • • • i 1 • • • • I • • • • , • • • i j i i • • I-I I I -* J3g&.-.VA^. ^  k 1 >v 1 I I I I i / i i i J Iv I C=0.060 i i i * 7 ^  i - J t . - f v i i i i i 7 Ji i | l l i l | l l l l | l l l l | l l l l | l l l l [ . lHl| l l l l"D ..4 I " I " I -I -I /I I 4 I I I I I I l l l | l l l l | l l l l | l l l | | l l l l | l l l l 1 I I l l I • I — I.... I... LI — I — I • • I J • I — T 0.1 0 -0.1 -0 .2 0.2 0.1 0 -0.1 - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 S / N / A £ 40 Mask=2.0 l o g i o ( R / R d " ) - 3 - 2 . 5 - 2 - 1 . 5 - 1 - 0 . 5 0 0.5 10239 - SAB(rs)cd Figure 4.28: Same as Fig. 4.21, but for IC 239. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 188 4.5.2 Measured Indices Compared to Models Having examined the radial variations in the individual indices, we now want to deter-mine if the observed trends translate into trends in the physical parameters of ages and metallicities. For this we need models to compare with and, as in the first two chapters, we appeal to the G A L A X E V models of BC03. In §2 we convolved the model colours with an exponential SFH to better mimic the SFHs of spiral disk galaxies. The same can be done with the indices in the G A L A X E V models because they provide the flux measurements in each index passband for the SSPs. Figure 4.29 shows a comparison of SSP versus exponentail SFH Mgb versus H/3 grids. Note that we plot the H/3 axis "backwards" in anticipation of comparison with colour-colour grids, which have age increasing from left to right. The "jagged terrain" of the SSPs are nicely smoothed out in the exponential SFH grids, but the general location of the grids in the index-index plane are the same, except that higher Mg6 values at H/3 < 4 are still consistent with the grids due to the mixing of old with the young stellar populations. This turns over at H/3 > 4 when only smaller values of Mgb are consistent with the model grids due to the increasing contribution of young relative to old stars. We take this opportunity to emphasize the labels on the SSP grids, as these will be omitted in most of the subsequent grids (for space considerations): Iso-.Z' tracks are solid lines with colour scheme: Z = 0.05 (red), 0.02 (Z 0 ; yellow), 0.008 (green), 0.004 (light blue), 0.0004 (blue), & 0.0001 (purple), and iso-age tracks are dashed gray lines at ages of 1, 2, 5, 10, 13, & 20 Gyr (from left to right and lightest to darkest gray shade). However, many other commonly-used SPS (e.g. Thomas, Maraston, & Bender; hereafter TMB03) models only provide the index measurements themselves, so convolutions with arbitrary SFHs are not possible (as of this writing). Thus, to facilitate comparisons with other authors & models, we compute light-weighted SSP ages and Zs for our line-index measurements. We also note here a degeneracy in the models that could potentially cause a problem in our age/Z determinations. At low metallicities (Z < 0.004) and older ages (A> 8-Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 189 4.5 4 3.5 3 2.5 2 1.5 1 4.5 4 3.5 3 2.5 2 1.5 1 Figure 4.29: Comparison of SSP [left panel] versus exponential [right panel] SFH Mg6-H/5 grids for the GALAXEV SPS models. \so-Z tracks are solid lines with colour scheme: Z = 0.05 (red), 0.02 (ZG) (yellow), 0.008 (green), 0.004 (light blue), 0.0004 (blue), & 0.0001 (purple), and iso-age tracks are dashed gray lines at ages of 1, 2, 5, 10, 13, & 20 Gyr (from left to right and lightest to darkest gray shade) for the SSPs. For the exponential SFH grids, the average age (eq. 2.11 with A = 18 Gyr) is labeled. 10 Gyr) the models turn over on themselves (i.e. become degenerate) such that the data could be consistent with being either very old (up to 20 Gyr) or relatively young (< 8-10 Gyr), where the turnover depends on the metallicity. If our data lie in this region of the grid, we will only be able to set lower limits on the age of the stellar population. In Figures 4.30-4.37 we present a selected sample of metal-line versus Balmer-line indices for our 8 galaxies, with the SSP model grids convolved with a FWUMboxcar over-laid. The metal-indices are roughly ordered in increasing wavelength from top to bot-tom. Balmer-line indices are ordered from left to right with decreasing wavelegth (and decreasing degree of emission-line contamination). The line-types and colour scheme of the grids are as in Figure 4.29 [left panel}. These plots highlight a number of issues related to the stellar populations of our Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 190 late-type spirals. First, the data span a large range of locations in the index-index plots, both within and among galaxies, implying a large range in ages and metallicities. As expected, the galaxies with significant emission in their spectra lie off the some of the SPS grids. The Ff/? index is clearly more affected than the higher order lines for most galaxies and, in the case of N7741, the data at many radii lie off the grids for all three Balmer lines. For those galaxies with data mostly within the grid limits (N0173, N0628, U2124, & N7490), there seems to be a general trend of decreasing age and metallicity from the centre to the outer regions of the galaxy, as expected from the line gradients observed in §4.5.1, but interpretaions are difficult by simple visual expection. This is in large part due to the fact that different index-index plots can imply significantly different ages and metallicities. For example, the Mgb versus H/3 diagnostic for U2124 implies higher metallicities than does the (Fe) versus H/? diagnostic. If the models can be trusted, such differences can be illuminating as they may imply abundance ratios in the galaxies that are different from those used to construct the models. Indeed, U2124 appears to have an enhanced Mg/Fe ratio in its central regions which is evidences in Figure 4.38 where we plot the (Fe) versus Mgb diagnostic for all 8 galaxies. Another obvious inconsistency with the models for many of the galaxies is that the C N 2 index lies well above the grids, particularly in the central regions (see Fig. 4.31 for N0628, for example). As mentioned in §1.2.2, a high Mg/Fe ratio is indicative of a short formation timescale and is often observed in elliptical galaxies. Short formation timescales are most consistent with monolithic collapse, but could be accommodated within the con-text of hierarchical merging. High CN indices have been observed in globular clusters in our galaxy (e.g. TMB03) and external galaxies (Burstein et al. 1984; Kuntschner et al. 2002; Schiavon et al. 2002a,b). The most often invoked explanation for the strong CN absorption in GCs is that, because of the high densities, stars may accrete car-bon and/or nitrogen-enriched ejecta from surrounding asymptotic giant branch (AGB) stars (Renzini 1983; Kraft 1994). TMB03 found that the CN strong GCs could only be explained by nitrogen-enhanced models, as carbon-enhanced models lead to serious Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 191 inconsistencies with other indices (Mgx & Fe4668 in particular). However, such CN-enhanced populations have also been observed outside of GC (e.g. elliptical galaxies, as mentioned in §1) and finding this enhancement in other environments would argue against the above mentioned explanation for the enhancement. We also note that CN-strong galaxies, Es and SOs in particular but also in some spiral bulges, are observed in Figure 9 of PS02, but the authors do not comment on this feature, except to state that carbon is poorly understood due to its many production sites. Trends are not easily defined for the strongly emission contaminated galaxies as many of them lie entirely off the model grids. However, the (Fe) versus Mgb diagnostic in Figure 4.38 demontrates that the Balmer-line indices are indeed the cause of the mismatch since, for all three (N7495, N7610, & N7741), the data are well within the grid limits. However, also of note is that these galaxies all seem to lie at lower values of (Fe) and Mg6 (i.e. towards the lower left corner) than the less emission-contaminated galaxies. This could be caused by the continuum filling from the strong continuum of the source of the emission (which is most likely star-formation). A strong continuum source would have the effect of diluting all of the indices not directly affected by absorption or emission features in the source. Additionally, if the source has a blue SED, the bluer indices will be more affected. This scenario is extremely plausible given the bluer SEDs seen in the spectra of these galaxies (and noted in §4.4.5). In future work we will attempt to model this effect using a pure H II template (from the galaxies themselves by isolating H II regions as was done in Fig. 4.7) and subtracting off incremental fractional contributions to attempt to uncover the underlying stellar population. Again, see §4.6 for comments on the individual galaxies. But first, in the following section we present our procedure for finding best-fit ages and metallicities using by comparing different combinations of the indices to the SPS models. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 192 2.5 2 1.5 1 2 0 - 2 - 4 - 6 - 8 3 2 1 0 H0 HyA Hc5F NO 173 Figure 4.30: Selected metal-line versus Balmer-line indices for NGC 173. The light blue star marks the center, red asterisks the bulge effective radius, and green crosses the disk scale length. The data points (black squares) have sizes proportional to log(r_ 1) and are connected by dotted black lines. Error bars are based on photon noise. Overlaid are the GALAXEV SSP models. Age (gray dashed lines) and Z (coloured solid lines) increase towards the top right corner in each panel. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 193 2.5 2 1.5 1 2 0 - 2 - 4 - 6 - 8 3 2 1 0 H/S H 7 A H<5P N 0 6 2 8 Figure 4.31: Same as Figure 4.30 except for NGC 628. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 194 2.5 2 1.5 1 2 0 - 2 - 4 - 6 - 8 3 2 1 0 H0 H7A H(5P U2124 Figure 4.32: Same as Figure 4.30 except for UGC 2124. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 195 Figure 4.33: Same as Figure 4.30 except for NGC 7490. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 196 2.5 2 1.5 1 2 0 -2 - 4 - 6 - 8 3 2 1 0 H/ff Hy A H<5F N7495 Figure 4.34: Same as Figure 4.30 except for NGC 7495. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 197 Figure 4.35: Same as Figure 4.30 except for NGC 7610. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 198 Hjff H y A H<5r N7741 Figure 4.36: Same as Figure 4.30 except for NGC 7741. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 199 2.5 2 1.5 1 2 0 -2 - 4 - 6 - 8 3 2 1 0 H 0 H 7 A H<5 P 1 0 2 3 9 Figure 4.37: Same as Figure 4.30 except for IC 239. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 200 L' 1 I 1 1 1 I 1 1 1 I 1 1 1 I ' 1 1 I J . 1 1 I 1 1 1 I 1 1 1 I ' 1 1 I 1 1 ' I I I | I I I | I I I | I I I | I I I | J I | I I 1 | I I 1 | I I I | I I I Figure 4.38: (Fe)-Mgtb diagnostic plot for all 8 galaxies. Symbols and grids are as in Figs 4 .30-4 .37 . Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 201 4.5.3 Age and Metallicity Fits The age and metallicity fits are done in a manner analogous to those performed in §3.5 except, as we mentioned there, the method has been extended to include more than just two indices in the fit. So, in eq. 3.11, N is now an arbitrary number of indices (at least 2 and up to the 24 we measure here). Note that, while the errors on the measured indices are independent (see eqs. 4.9-4.12), the model tracks are not orthogonal (e.g. Fig. 4.29). This produces non-orthogonal errors in age and metallicity, which can lead to spurious correlations if not understood. Additionally, the degree of non-orthogonality of the age and metallicity tracks changes with position on the grid. For this reason we have used Monte Carlo methods to model the "effective ranges" of the fitted ages and metallicities, taking into account a normal distribution about the measured errors (see description in §2.7). We have seen in §4.4.5 that many of our galaxies have significant amounts of emis-sion, particularly in the central regions where no H II region masking was done. Many authors attempt to correct for emission by fitting emission-free templates to their galaxy spectra. The templates can either be from linear combinations of stellar templates (e.g. Gonzalez 1993), or from models (such as the G A L A X E V models used here). Neither method is ideal. The former requires a library of stellar templates, ideally taken in the same run with identical observing conditions, that match the galaxian spectra ex-tremely well. The latter imposes a model dependence as, if a model with a given age is used to make the correction, the same model age will be returned (you get out what you put in). Additionally, if there is a significant amount of dust in the galaxy, causing a reddening of the spectrum, this must be included as another model-dependent free parameter in the template fits (for both the empirical and model methods). Gonzalez (1993), in a study of 40 elliptical galaxies, created individual templates from stellar spectra obtained with the same observational set-up for each galaxy spec-trum (at each radius) for the purpose of measuring accurate velocity dispersions. His templates fit the observed spectra to within ~ 1%. Division of the galaxy spectrum by the best fit template enabled detection of faint levels of emission (seen predominantly Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 202 in [O III], and H/3) in over 60% of his sample. From the galaxy/template spectrum, he computed pseudo-EWs for the [O III] A 5007 & 4959 A and H/3 emission and found them to be strongly correlated as EW(H/3) = 0.7EW([O III]). As such, the H/3 in-dex in absorption could be corrected for emission line fill-in by adding the correction 0.7EWQO III]) to the measured index. He also made a correction to the Fe5015 index as the [O III] A5007 line is contained within the limits of the central bandpass. Gonzalez (1993) also strongly emphasized the importance of having suitable templates for reliable measurements of both velocity dispersion and emission corrections. The tight correlation between [O III] and H/3 emission was tested for later-type bulges (SO-Sb bulges) by PS02 using a method similar to that of Gonzalez (1993). While they found the same [O III]-H/3 relation for their earliest-type bulges, the six late-type bulges did not follow the Gonzalez (1993) correlation for ellipticals, with some bulges showing a stronger H/3/[0 III] ratio, while others showed clear [O III] emission, but none in H/3. This is likely partially due to the stronger absorption in the younger stellar populations, but could also be due to template mismatches, or true physical differences in the emission line regions in these galaxies. Thus, no straight-forward correction seems to apply to spiral bulges and, by extension, their disks as well. The higher-order Balmer lines are much less affected by emission line fill-in. For a broad range of physical conditions expected in galactic emission-line regions, the line-strength ratios of H5 & H 7 relative to H/3 are 0.25 and 0.5, respectively (Osterbrock 1989), providing some relief from emission line fill-in. Indeed, in many of our galaxies, the H/3 index lies far off of the model grids (from which ages ^ 20 Gyr would be inferred), where as the higher-order indices imply much younger ages (~5 Gyr). Emission fill-in has the effect of weakening the absorption and thus increasing the apparent age of the galaxy (and is often invoked as a main factor in the measurement of elliptical galaxy ages greater than the age of the Universe). We can therefore consider our age determinations as upper limits. Given the above mentioned issues with making emission-line corrections to the mea-sured Balmer-line indices, we have instead opted to use a scheme in which the emission Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 203 line affected indices are systematically elliminated from the fit, and the age determina-tions are compared. We prefer this method as it is less model-dependent, and takes full advantage of the short measurement baselines of the indices (significantly weakening any effects due to dust redenning, see §3), while still using information from the entire observed SED. Our fitting scheme is as follows: Case A: Fit all well-measured indices By "well-measured" we mean the index (all 3 bands) is not severely affected by systematic problems such as the wavelength gaps between the 3 CCDs, listed in Ta-ble 4.6 for each galaxy, or strong sky lines (which are extremely difficult to subtract accurately), listed in Table 4.7. The indices that were systematically eliminated from all fits are highlighted in red in the deviation plots (left panels in Figs. 4.39-4.46, see below for plot description). Note that we never include the NaD index in any of the fits as this index is affected by sky lines as well as poorly understood absorption from the interstellar medium of the galaxy. The TiO indices were also excluded from all fits except for N0628 due to sky line subtraction issues. Case B: Case A — (all Balmer-line indices, Fe5015, & CNi) Here we attempt to measure ages and metallicities by eliminating all indices severly affected by emission. This is feasible because of the significant age dependence of G4300 and the weak age dependence of the metal-line indices. It is clear that all Balmer-lines could suffer from emission line fill-in. The Fe5015 index has the [O III] A5007 line in its central bandpass thus could also suffer fill-in. While not often mentioned in the literature, we note that the CNi index could also be compromised by a strong source of emission due to the fact that it contains the H<5 line in its blue continuum (which forces the index to larger values). This is clearly the case for at least two of our galaxiess, N7495 & N7741, where H<5 is seen in emission (see their respective spectra in Figs. 4.16 & 4.18). As an example of the strong influence emission can have on CNi , including CNx in the fits, having removed all Balmer-line indices and Fe5015, the age is forced Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 204 to the maximum model age of 20 Gyr to accommodate the large C N i . When C N i is removed, the best-fit age is closer to 1 - 2 Gyr. We also remark, however, that because of our resolution limitations, if there is a strong emission spike in H 7 , its tails can creep into the red passband of the G4300 index (e.g. see Fig. 4.18 for N7741 at r = 3.3"), but G4300 was not elliminated form the fits. Case C: Case A - (H/3 & Fe5015) Elliminate only the indices most severely affected by emission. This case should be appropriate for galaxies with only small amounts of emission, which could be un-detectable in the spectra, because the higher-order Balmer-lines will be only weakly affected. If the age estimates here are younger than in Case A, the implication is that emission is likely present, whether or not it is clear form the spectra. Case D: Fit only H.fF, [MgFe]', & G4300 This group of indices are the least affected by abundance ratios (TMB03; Thomas, Maraston, & Korn 2004). While only a few of our galaxies showed evidence for signifi-cant [a/Fe] enhancement (N0173, U2124, & N7490), many showed C N 2 values in large excess of the model grids (N0628, N7490,10239), which this combination of indices also avoids. •Results from the age and Z fits for all 4 cases are displayed in Figures 4.39-4.46. In the left panels of each plot, we show the deviation in units of error for each index in all four fits, quantified as Xi = (data* — fitj)/<5j. Each case is denoted by a different colour: Case A: green, Case B: blue, Case C: pink, & Case D: gray-brown. Three-pronged point-types mark the indices included in the fit for each case, while open squares mark indices omitted from the fit. Point sizes and rotations were adjusted for visibility when points are overlapping. Thus, if a given colour is not visible for a given index (except (Fe), which is never explicitly fit, and [MgFe]' for which we only show the Case D results), it is because the x f ° r that index is off scale (for example, the Case D deviation for TiC>2 at r = 0" in N0173 is x = —27, so no gray-brown box is visible). The dotted black Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 205 vertical line marks the zero deviation case. The indices are labelled on the y-axis (the dashed horizontal lines are there to guide the eye to the index labels); those marked in red are the poorly-measured indices that are never included in the fits (note they are different for each galaxy). For each galaxy we plot deviations at 3 different radii (labelled in green above each panel): the central point [left panel], the point closest to the bulge re [middle panel], and the point closest to the disk scale length rd [right panel}. Below the radius label for each panel are the best fit ages and metallicities for each case. The right panels in Figures 4.39-4.46 show the age gradients in log [top panel], metallicity gradients in log [middle panel] and the x2 figure of merit as given in eq. 3.11 [bottom panel], as a funciton of the logarithmic radius scaled to the disk scale length (with labels in arcsec on the top panel) for all four fit cases (the colour scheme is the same as in the deviation plots and is labeled in the bottom panel). The data points for each fit are connected by dotted lines to guide the eye. Labels on the right axes show the linear values. For each galaxy, the seeing FWHM (blue), re (red), and rd (green) are indicated with arrows at the top. The dashed lines in the age and metallicity plots denote the model limits, so fits with these values should be treated with caution (recall that we do not attempt extrapolation). The error bars for ages and metallicities are taken as half the interval containing 68% of the 100 Monte Carlo realizations (i.e. the lo confidence interval). Error bars for saturated (grid limits) fits should also not be trusted. Examining the fits for the different cases in Figures 4.39-4.39 we see that indeed, in the different cases can find very different best-fit parameters. The reason for this in some cases is clear, for example, when there is significant emission and the fits are different depending on which of the Balmer-lines (if any) are included in the fit (see, for example, N7495 fits in Fig. 4.43). From the deviation plots we also see that, in general, the indices affected by prominent sky-lines are poorly fit (they were excluded from the fits, but the deviation parameter was still computed). For the galaxies with very high central CN indices (e.g. central region of N7490) those indices were poorly Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 206 fit in all cases. There are even examples where the four cases find different fits, but all indices are equally well fit in each case, such that there is no obviously deviant index with enough weight to force the fit to a different value (e.g. the r = 14.5" point for U2124, Fig. 4.41 [left panel]). These are likely fits in the age-degenerate, low-Z regions of the grid. Fits with large uncertainties (according to a Monte Carlo.sampling of the measurement errors) are also likely in the more degenerate regions of the grids. Trends in age and metallicity from the fits are discussed for each galaxy in the next section. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 207 Age (Gyr) to ( 3 § V ) § ° I ( s Z / Z ) b l § o i 1 1 1 ! I I I 1 : i i i i i m - i - i - i U i " • L i i . ^.. r . . . . . . . : i i 1 T i i i " i i i i i i i L i i i i i i i • I I 1 11 | ! 1 1 I I 1 11 1 1! 1 1 l | -» ii in i t ii III f ft ii ii r .ii ii ii |i JI ii i : p ii ii f .( ii in r t II HI 1 .II ii 1 ii j n i -B ) II I. li ! 1 i: i! I X li II i II ' * t li it . : I 1 ill t li 1 ill 1 li II T t 1! t II li 1 II I 1 I I I 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 : 1 1 1 1 1 1 1 1 1 1 1 1 1 -11 I 1 1 1 1 1 : i i i i i i i - ill 1 1 l 11! [ g 1 1 : 1 1 i i i 1 1 1 1 r j i it II ii i it i jk " 1 i i I i i 1 0 T i, in i t i. in J id t II - i i i i i i i i i  I   1 1 1 1 I   f    j i i i i i i i i i i i i i i , i : i 1 li il ill 1 '1! II ill .11 II II. t l! i ill -i i i i i f i i I * i i i i x i ' : 1 1: 1 1 ill 1 'li II it 1 "I 1 it li I , J : l! t Il II it t II 1  ill II II it 1; II ill li 1 li II it t II II III ' li II It 1! II il III j 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 " L 1 1 1 1 | 1 1 : i i i i _ i I - i i i i . i i ; i i 9 IP f i m ri--y-^--$—j~-r~l— : II l i li j I I c ii i II l i i IB j . 1 i i i ii i ii ' • i i i i i i i c 1 1 1 1 3 • ' 1 ?!   ! I  !  II I  1 1 1 1 1 FJ I I I I I I I I I I I I I I I I I : i i i i i i i i i i i i i i i i i J * II i i i m e ? " i I T I i i II : i l m l l l I l 1 I .1 l i ( JI il i ~ I i II i it t li i I I li II it it j II HI j • CO o z N o o 2 S1 1 \ , ° ' " i -i« *r -^i eo" TJ o v — 2 I i , ' g " t o o o I o > 0! Q BOv l O l Q l O l O l Q — ( II « I I d . U . EL , E*. lO CQ. 5 = l O -co « to CO co in rf * V V -— -— uo n o LO CO o rf CO CO rf rf rf • U O N "5 iC W " CJ rf CD CJ Figure 4.39: Results from age and metallicity fits for NGC 173. See text for plot description. O P c fD o to CD 3 n> CD c —\ C O t o X h X I CD O (V) oo Age (Gyr) log .^Z /Z . ) [MgFe]' <Fe> H7, H«, H7A H<5» TiO, TiO .r-NaD Fe57B2 Fe5709 Fe5406 Fe5335 Fe5270 Mgb Mg, Mg, Fe5015 H/S Fe4668 Fe4631 Ca4455 Fe43B3 G4300 Ca4227 CN S CN, r =0.0" 3.310.0 1.3 2.8 -0.2-0.4 0.4-0.2 r = 17.2" 0.9 0.9 0.9 1.7 0.4 0.4 0.4-0.5 r = 78.4" 3.7 0.8 2.3 1.6 -0.8 0.2-0.7-0.4 I | I I 1 | 1 I I | I I I | I I 1 | - y j - ( - -- *K-1' - - ;*-<-- 4 - * --Hti j —dBD-; -H>i --••-*k---- - H * . -)H-(r -B-J M - e --E3+-l-i - -D - -x*l --•4» — M i - -* i -X- -^A 1 HI " -_ __j r0yq. . . • I • ' ' I ' • • I • ' • I • i 1 •111 1111 i f i !#=a - -** --4--• - J U — j - * --*H jji 4 91 * 4 HQ . I • • • 1 j • • I • • • I • • • I • 11 • • I • 1111 • 11 • i • I ^ -B — . — •Q— --H! Q-BJ Q- -jM * f 4* - f f l - - ; *•< =4< <*K 1 *< :0H -SHfv - Q * * C -•EM • - D - -«&<: « * ! — *( ;<* -— * H * k — • 4 — BiB - • i • • • i • • • i • • • i • LL, -20-10 0 10 20 -20-10 0 10 20 -20-10 0 10 20 Deviation (x = (datarUtl)/6l) N0628 Rad (") 0.06 0.1 CD 1 < o M 0.5 O O h 0.5 O o tSJ -0.5 &0 o -1 h 80 60 %< 40 30 A: all well-measured indices ii: remove Balmer, Fe5015, & CN, C: remove H0 & Fe5015 D ^ n c g i j a / F e ] affected 8 -1 1 I 1 L _E_L o 0.04 0.008 N 0.004 0.001 -3.2 -2.8 -2.4 - 2 -1.6 -1.2 -0.8 -0.4 0 l o g 1 0 ( R a d / R ^ ) N0628 I •SP X) n o a o e 0> P5 CO f I 1—1 I rt> c» LO o OO OP c ro CO QJ 3 QJ c rt> oo CO <T> X n IT) CD O ro ro Age (Gyr) log 1 0 (Z/Z. ) [MgFe]' <Fe> Hrr H<5, HrA H<5» TiO, TiO, NaD Fe5782 Fe5709 Fe5406 Fe5335 Fe5270 Mgb Mg, Mg, Fe5015 HP Fe4666 Fe4531 Ca4455 Fe4383 G4300 Ca4227 CN, CN, r =0.0" 6.217.5 6.5 5.3 0.2-0.1 0.2 0.3 r = 14.5" 6.2 5.0 5.020.0 -0.3-0.2-0.8-0.7 r = 23.9" 9.0 3.7 5.7 10.0 -0.8-0.6-0.6-0.9 111111 • 111111 4. -o-i —XJB - -HOQ - -jfrB- - -- - i B N- Q - -•-! —cs>i et> * * | - I - B H - -; H D - - 04-i - B -%>~f{-9. 4-4-af.-l -• i • • • I • > | M I | I I I B J - i • - -m -B- 43-.«ja-. - -efc | * o --4 --4 - -m --4 — i * - - 4 3 -- * B 1 I 1 • • I I i • I • • i I I i i • | i i i | i • i Ji -fj - - O H -9X~ _ ^ -4u --xe -»# ! ^ S B D -- 4EB — --a* 4 -4 -, . ^ i i i i i i i i i i n i i - 20 -10 0 10 20 - 2 0 - 1 0 0 10 20 - 2 0 - 1 0 0 10 20 Deviation (x = (data,-fit1)/6i) U2124 a> 1 oo < , o OO 0.5 O 0 0.5 P o £-0.5 OO o " -1 80 60 %< 40 20 0 0 . 0 0 0 . 1 I 1 1 1 1 I I I I I I I I I I I I I I I I I I I I I I I I I all wel l -measured indices remove Balmer. Fe5015. & CN, remove HB & Fe5015 non [a /Fe] affected (M) . ^ l . . . 1 .e, * ' • . ° I Q , a 2 0 10 B 6 4 S 1 0 . 0 4 o.ooa M 0 . 0 0 4 -2.8 -2.4 -2 -1.6 -1.2 log10(Rad/R«") -0.8 -0.4 U2124 s c-f-x) o •1 I 1 C L o ft ft C L Cb' 13 CO f I CD as I CD to lO o CO OP c T co m m 3 ru QJ cn Lo n> x n CD CD CD CO O r =0.0" r = 5.6" r = 20.9" Age (Gyr) 10.712.210? 4.5 3.7 4.0 3.7 5.0 3.0 5.0 3.020.0 log10(Z/Z.) 0.3 0.3 0.3 0.4 0.2 0.2 0.2-0.1 -0.6-0.7-0.6-1.0 10 20 - 2 0 - 1 0 0 10 20 - 2 0 - 1 0 0 10 20 Deviation (x = (data1-fit l)/(51) N7490 0.06 o.i Rad (") S 0.6 i 1 < O CO ISI O 0.5 0 0.5 0 0.5 -1 80 60 40 20 0 4 h 4 * ' H—| I I I |—I I I—|—I—I—I—|—h-t—1 i—h I | I I I | I I I | I I I | I I I | I I I | I I A: all well-measured indices B: remove Balmer, Fe5015, & CN. C : r e m o v e x H f ? & F e 5 0 1 5 D : n o n [ot /Fej affected _J I I I L 0 ° o • i I i i i I i ' ' j " 1 V 20 OB 9 a 8:8AB tsi 0.004 o.ooi -2.4 -2 -1.6 -1.2 -0.8 log ) 0 (Rad/R«0 -0.4 0 N7490 9 1 CD X) o o Cl 3 ft) to a, CD tf o N cx § Co i Q fu B. I—1 o Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 211 Age (Gyr) cd K ,- CD CO © - C J I D Tp cc - d o o 6 I'I'I ' I ' I I'l' I ' I J I' 1 1 i ' ate oo S do d m i -e—Dp-45 I i i i l i i i l i i i l i i i 2 t B * c ° • - l O _ «> •O b. e -a 4) cn « > o s a> 9a x > o 5 0) X < CO o o lO 05 z He, in o iq iq 6 d d I o o as co o o CM 3-D \ ca oc o ( 3 § V ) 0 l 3 ° i ( e Z / Z ) b l § o i o cvi to d II o d II r-o d i d o I d oi d i i i i i * T f ! m T i i i i i i i i i i i i i i i i i I I: III ll II II | II II |l I g Q ' I 1 1 1 I I I ±9 I I I I I I I I I * T I I I I I I I I I I i I t I I 0 I li ll ii ll Hi X i i i i i i \ Y m iA i i I I I T i • ? i i l l l l l i i i i i i i i m • X II 1  .1 II II L ^ o r-to 2 IS •a o '—' 2 L 2 > 0) _ Q Figure 4.43: Same as Figure 4.39 except for NGC 7495. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 212 Age (Gyr) -o oo co o o ««> •» CM HOC 01 —OO O " 1 I'l' 1 1 I 1 I IFF! ' I ' I T 1 If I: i f oi q I CD 3 SS s do o I'l" 1 I a o i 9 ljn n XOI f 2 O + e 2 O <B £ l S3- R»T BJ CJ Q d © d lO — d I o CO (3§V)° Soi ( 0Z/Z) b ,So t 1 , , . 1 , , , 1 a o o CO CJ CD <> 2 3 « K \ £0 Di SO o ro 6 CM d 6 o n e . d -< C\! | co co c\i d <o co d d o — L 1 II 1 1 1 1 1 1 : i i i i i i i 1 i - i i i i i i i I i t l I 1 I H II: li. IU J , r t - - V - ^ - ^ " $ T " H , : 1 P X ts T il i i i " i i r m i ^ i i t II 1 ii II. j in T n i • II 1 1 1 1 1 1 1 1 1 II 1 1 II: I I .1 ii 1 .CD J Hi Is Ill .11 HI 1. * < III |. .:; II ill 1 1 1 <|, i f Hi I III 1  III 1 ,11 | ill 1 -J i T i i i i i * i i i I i i i i : li II' .J II 1 1 I ill § ill 1 III 1 HI 1 II 1 ~ 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 J •- 1 1 1 1 1 | 1 1 V II 1 1 1 1 II 1 1 II 1 1 1 1 1 i • M in I il i| | i in i in i: j i HI li A ii nl i; m ii ii i * :i III r 'T 1 1 1 1 1 j 1 I 1 1 1 1 1 1 1 ¥ 1 1 Hi || [|| I l| il i| |. A 9^ i • i i » { 4 i f i44 r'vi41< i l l M ^ i II t li If n II II II. 1 ill 1 Hi 1 ill f III I ill T T 1 III T 1 1 Hi I i " l i C i l T i l l l l l l l i l i l i I D i i i i i i i i " 1. II 1! 1 9 II II li II 1 ill 1 It 1 ill 1 II 11 1! |! Ill il III li. iii .1 iii 1. j • I l l • 1 1 I 1 1 J ! LU 1 1 : i i ii ii i, ii flj I i * - i i i I i I i i 0 : * v l i * 1 1 & 1 -| I 1 1 1 1 | 1 1 1 1 1 | 1 1 1 1-I 1 1 III II 1 'I III 1 1 .'I III |' ill 1 ill |i '• f II H III 1 T JL III 1 1: il II! f ill II Hi ijj.-n 1! Il ,1 li II ll H 1 1 II T III li •SJf II III II J3 II 1  J II II II 1 III II II I III I Hi' 1 1 l!' • ii I I i I o I i 1 H I I i t il li i j i i i i i i i i i i i -• T - r ^ f i - 7 - T - g - f -• 1 1 1 1 , 1 1 1 1 f i l l 1. II! 1 II II III II 1 '•_ ill 1 ill • A 1. Ill 1 II • 1 1 1 1 o CO O CD o i> Z o o O o CM ' o | CM o 03 T3 a o II X 1 a CM C o O CM ati O evi — o o o CM N 0) a; £ 2 £ ? 2 2 E o o co r» LD lO iO iO lO V V 4> 0) V b Gi. Gb Eh Eh 2 s CQ. CO co 0) co to •<(• • * V o CJ CO o CO o co co c o EL, M Z 2 CM CJ CJ o o Figure 4.44: Same as Figure 4.39 except for NGC 7610. g i * <T> a. ^ L O Ol 3 m QJ in c —^ 4^  Lo lO o> X n ro •a O 4^ . O r+ CO CD O--t-i - s ct> r+ QJ X, i n ' in r> QJ_ ro in r = 0.0" r = 17.2" r = 63.3" Age (Gyr) 20.0 1.0 3.0 0.0 20.0 0.8 2.5 0.0 20.0 0.820.0 0.0 log.jCZ/Z,)-1.3-0.6-1.0 0.0 -1.5-0.4-0.9 0.0 -1.5-0.4-1.6 0.0 NaD Fe5782 Fe5709 Fe5406 Fe5335 Fe5270 Mgb Mgz Mg, Fe5015 HB Fe4668 Fe4531 Ca4455 Fe43B3 G4300 Ca4227 CNg CN, 10 20 -20-10 0 10 20 -20-10 0 10 20 Deviation (x = (datal-fit,)/<5|) N7741 0.06 0.1 Rad (") s r. 0.6 X 1 6 4- 10 sb too M 1 0.5 0 -0.5 -1 -1.5 -2 -2.5 0.5 F 0 tsi -t—t- i i i i i i T I T 1 n l l j •ia..mM H B t f l I I •' I I I I i . -t—t- I I I 0 h-e. 0.5 -1 1.5 80 M 60 40 20 0 I 9 * - . !T1 I ' ' ' I ' ' V • • i - i - i I I I I I I l'T'4- I l°l I I I I I I I I I I I I I I I I A: all well-measured indices B: remove Balmer. Fe5015. & CN, C: remove HB & Fe5015 D: non [a/Fe] affected a Poo"1 _i i I i i i_ Ea_a_ o x S • _l 1 L I B I • 0-2 O S is 0.0004 0.0001 -2.8 -2.4 -2 -1.6 -1.2 -0.8 log10(Rad/R«») -0.4 0.4 N7741 OP c CD w cn CD 3 CD cn Oq CO CD CD x o CD XJ O CO CD Age (Gyr) log 1 0(Z/ZJ [MgFe]1 <Fe> H7r lift, Hc5A TiO, TiO, NaD Fe5782 Fe5709 Fe5406 Fe5335 Fe5270 Mgb Mg, Mg, Fe5015 HP Fe4668 Fe4531 Ca4455 Fe4383 G4300 Ca4227 CN, CN, r =0.0" r = 10.8" r = 21.8" 7.0 8.3 7.0 3.5 1.7 1.8 1.6 1.8 20.019.5 1.3 1.3 -0.5-0.5-0.5-0.3 -0.1-0.0-0.0-0.5 -0.9-0.9 0.1-0.7 11111111111111 • • 11 B; 4s -ki-h -i<* - 4<e 4 4 c*. -« i I ' ' • I • ' • I • ' • I • ' • I I I I I I I I I I I I I I I I I I I I ,1 - -s-i D-- ---*! = ' - -* 3 -- -4-- r * - i -- L - - - 4 L 4 -14) _ , . _j fgj i , , , i , •, i • , , i , , , i  I I I I I I I I I I I I I I I I J , -O- -- - £ ] > © - -- -jJO—0-i > *<" •>-Jfo-i-- >4-- -+# 4 - ->Mo \M»--oo- -j 4 - i * -H 4 1 1 go i * • • i * • • i • • • i • • • i -20-10 0 10 20 -20-10 0 10 20 -20-10 0 10 20 Deviation (x = (dat^-fit,)/^) 10239 0.06 0.1 <D 1 <_ o 00 0.5 O 0 0.5 n N u \ N "o-0.5 M O ^ -1 80 60 % 40 20 0 A: a l l wel l -measured indices B: remove Balmer, Fe5015, & CN, C: remove H/S & Fe5015 D: non [a /Fe] affected ** | a S • . a !-«> I . • • I , , , I , • p. I • .o . I 9 • Q 1 q • • I x 10 > 8 TO e » 4 f 0.04 0.008 tSl 0.004 -2.8 -2.4 - 2 -1.6 -1.2 log 1 0 (Rad /R»" ) -0.8 -0.4 10239 9 I Ct> ct o M N O O i cx CD tf ex CD* a cn Fx &> i—* | i LO I—4 Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 215 4.5.4 Colours Compared to Models We now turn to a comparison of the ages and metallicities implied by our index mea-surements with those obtained from broadband photometry. As was emphasized in §2, an infrared (e.g. H or K) passband is needed in order to break the age-metallicity degeneracy in the broadband colour plane. Five of our galaxies are in common with the de Jong (1996) sample (U2124, N7490, N7495, N7610, N7741) for which BVRIK SB profiles are available (unfortunately, adequate photometry could not be found in the literature for the remaining 3 galaxies). Since our index measurements never ex-tend beyond ~ 2 r d , we cut the profiles off at a 73-band SB above 24 mag arcsec -2 so that sky subraction should not be a problem. Note that the observing conditions for the if-band observation of N7741 were not perfectly photometric, thus a small (< 0.2 mag) systematic offset could be present in the if-band profile, fn Figures 4.47-4.51 we present colour-colour plots for the five galaxies. The grids are the same G A L A X E V SSP models used on Figures 4.30-4.38 where coloured solid lines are constant metallicity tracks (yellow is solar) and dashed gray lines are iso-age tracks (darkest gray is 20 Gyr, and progressing to lighter shades as 13, 10, 5, 2, & 1 Gyr). 4.5.5 Age 8z Metallicity Fits from Colours Using the same G A L A X E V SSP models, we fit age and metallicities for each radial bin in the SB profiles using all five passbands: BVRIK. The fits from the colour gradients are presented in Figures 4.52-4.56 along with the fits from all four cases for the indices. A discussion of the colour grids and fits is provided for each individual galaxy in the following section. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 216 H 3 5 " GALAXEV > ^ i 3 h 0.4 0.6 0.8 1 0.8 1 1.2 1.4 1.6 1.8 0.6 0.8 1 1.2 1.4 B - V B - R V - I U2124 Figure 4.47: Colour-colour plots for UGC 2124. The data are from de Jong (1996). Lines and point types are as in the left panel of Figure 4.29. > 2.5 2 1.5 3 OS 2.5 1.5 1.5 1 7 GALAXEV /r^Sgfki W1 •• •^"^r^TTTTT 11 • • • i • 1 1 1 1 • 11 • ~ nffrt... i . . , i , , , i r 3.5 < 3 I 2.5 K 2 1.5 3 2.5 , 2 1.5 1 : 0.4 0.6 0.8 1 0.8 1 1.2 1.4 1.6 1.8 0.6 0.8 1 1.2 1.4 B - V B - R V - I N7490 Figure 4.48: Colour-colour plots for NGC 7490. The data are from de Jong (1996). Lines and point types are as in Figure 4.30. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 217 B - V B - R V - I N7495 Figure 4.49: Colour-colour plots for NGC 7495. The data are from de Jong (1996). Lines and point types are as in Figure 4.30. Figure 4.50: Colour-colour plots for NGC 7610. The data are from de Jong (1996). Lines and point types are as in Figure 4.30. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 218 K 3 5 I 3 > 2.5 2 1.5 3 , 2.5 K 2 1.5 * 2 I " 1.5 1 | I 1 I | I I I | I I I" G A L A X E V ^fg" 0.4 0.6 0.8 1 B - V 0.8 1 1.2 1.4 1.6 1.8 0.8 0.8 1 1.2 1.4 B - R V - I N7741 Figure 4.51: Colour-colour plots for NGC 7741. The data are from de Jong (1996). Lines and point types are as in Figure 4.30. A 0.001 - 1 . 6 - 1 . 2 - 0 . 8 - 0 . 4 0 log 1 0(Rad/R«») c o l o u r s : BVRIK U2124 Figure 4.52: Age and metallicity fits from indices and colours for UGC 2124. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 219 -1.6 i o.ooi -1.2 -0.8 -0.4 0 log10(Rad/R^) colours: BVRIK N7490 Figure 4.53: Age and metallicity fits from indices and colours for NGC 7490. Rad (") r„ 0.5 I 1 e 46 10 ^ 1 CD O X < 0.5 O W 0 -0.5 h 0.5 [S3 \ -0.5 h on o -1 h T i r r • i T T T ^ : ? . j f | T TTTTfT - I — I - •4 1 h - I — I — I - H 1 h • -1f H h H 4 ft H 2 - 6.8 - 0.8 - 0.4 _i 1 1 1 1 1 I i i i I • • • 20 O 1 0  ~ H 0.04 8:8ie N 0.004 o.ooi -1.8 -1.2 -0.8 -0.4 0 log10(Rad/R*) Colours: BVRIK N7495 Figure 4.54: Age and metallicity fits from indices and colours for NGC 7495. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 220 s 0.5 i I • OJ 1 60 < , o 60 0.5 O 0 0.5 © n tsa -0.5 H 60 O -1 •4 1 1-Rad (") 1 0 . A. x i t ] •a \-H 1 h H !-•—I- -I 1- •I—I—h 1 0 8 s > n o 1 0.04 8:848 M 0.004 0.001 -0.8 -0.4 0 log 1 0(Rad/R««) c o l o u r s : BVRIK N7610 Figure 4.55: Age and metallicity fits from indices and colours for NGC 7610. -1.6 o.oooi -1.2 -0.8 -0.4 0 log 1 0(Rad/R«») c o l o u r s : BVRIK N7741 Figure 4.56: Age and metallicity fits from indices and colours for NGC 7741. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 221 4.6 Comments on Individual Galaxies In this section we examine each galaxy individually, attempting to synthesize all of the observations and model fits. N G C 173 - SA(rs)c This late-type spiral is normal in appearance with fluffy, but well-defined spiral arms. It appears slightly inclined (see Fig. 4.3) and has the largest peak rotation velocity in our sample of galaxies (see Fig 4.6). The slit was oriented along the major axis. Figure 4.8 shows that there is a significant amount of emission throughout the disk (most of which gets masked out), as well as some emission in the center (which does not get masked). There are a few emission regions at ~ 50" on both sides that make it past the S/N threshold, but the SB at these radii is very low, so they will only have very little weight in the final radial bin. The radial spectra, Figure 4.12, show evidence of the above-mentioned central emission (in the Ho; line), and possibly some disk emission. Sky subtraction issues of strong lines become evident beyond about 5", so we do not include sky-line affected indices in any of the fits (Fe5782 & T i 0 2 for this galaxy; TiOi is also excluded due to the gap location). From the radial index profiles (Fig. 4.21), we see that the data go out to ~ 1.4 rd-The radial trends are very symmetric on either side of the galaxy. There are small features that do not precisely agree on both sides, but this would not be expected as they are not mirror images of each other. Clear negative gradients are seen in most of the metal-line indices, but with a slight "bump" between re and rd. This is likely associated with spiral arms crossing the slit (which are not perfectly symmetric on either side of the slit). Additionally, TiOi (affected by the gap) & Ti02 (affected by a strong sky line) don't agree. We do not include the TiO indices in any of the fits, (i.e. they are not considered "well-measured") and, as can be seen in the deviation plot (left panel of Fig. 4.39), they both deviate significantly from the best-fit models in most cases. The Mgi does not seem to agree with the trends displayed by Mg2, this is likely due to the fact that they are both affected by the gap. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 222 The Balmer-line gradients roughly agree, except at the outermost regions where H/3 appears to suffer from emission line fill in (it flattens whereas the higher-order Balmer lines rise to the edge), or it could be effected by the gap which, due to line curvature, is only within the H/3 index range at large radii. The G4300 index also has a strong age-sensitivity, but in the opposite sense of the Balmer-lines (it increases with age), and thus should roughly mirror the Balmer-line gradients, which indeed it does. On the index-index plots (Fig 4.30), both H/3 & B.5p extend off the model grids at intermediate radii, while R^A does not. For H/3 this could be due to line filling, but is not a likely explanation for H<5 (otherwise it would also appear in the H 7 indices). The most likely explanation is that B.5 is right at the edge of our wavelength range thus could suffer from less-secure calibration. This galaxy appears to have super-solar [a/Fe] abundance ratios (see Fig. 4.38) in a range between ~ r e to 1/2 r^, which is not expected as disks are not thought to have the short star formation timescales that lead to high [a/Fe] ratios. The C N 2 index is also high in central region, but only on the B.JA grid [top middle panel] as emission line fill-in seems to have moved the H/3 index low enough to bring the C N 2 onto the grid. As discussed in the introduction (§1), such CN anomolies have been seen in many globular clusters and some spheroids and has been associated with an enhancement in nitrogen (TMB03), but the cause of the enhancement is not understood. The fit without any Balmer lines (Case B), finds a much older (~ 6-20 Gyr) and a slightly less metal-rich (Z ~ 0.01 -0.02) central region (all Balmer lines, except perhaps H/3, are poorly fit in this case). Al l other cases (except one point for Case D) agree on a younger age for the inner region (~ 2.5-4 Gyr) and high metallicity (Z~0.03-0.04). Close to re there is a large jump in age from very young to very old, and Z from metal-rich to metal poor. All four cases agree on this jump. Examination of the grids shows that the position of the jump is located in the grid region that is very degenerate in age. The age fits here are thus very uncertain and it is possible that the data are consistent with a much smoother age gradient. The last few points around 1 rd are consistent with very old (20 Gyr) for Cases A Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 223 & B, but with young (~ 1.5-2.6 Gyr) for Cases C & D. If the latter are considered, then the inner disk and central region do have similar ages, but the metallicity is much higher in the center (Z> ZQ) than at ~ lrd (Z ~ 0.0004). Unfortunately, we could not find photometry for this galaxy in the literature. N G C 628 - SA(s)c This face-on spiral galaxy is the prototype of highly regular, two-principal-arm, grand design spiral structure. Figure 4.8, bottom plot, reveals significant emission throughout the disk (most of which gets masked out), as expected from the clear knots of SF along the spiral arms seen in the image (Fig. 4.3), but none in the very center (within ~ 5"). Accordingly, there is no Ha: emission evident in the spectra until r > 5". The radial bins extend to ~ 1.8 rd. Sky subtraction issues of strong lines become evident beyond about 70", so we do not include sky-line affected indices in any of the fits (which in this case was only NaD, which is excluded regardless). Also excluded from all fits was Fe5015 which is affected by the gap for this galaxy. Again the gradients are fairly symmetric. Negative gradients are seen in most of the metal-line indices, but again there are "features" present that are likely associated with a spiral arm crossing the slit (which are not perfectly symmetric on either side of the slit). Many of the metal-line indices turn up at the largest radii which could be from a second spiral arm crossing the slit, The Balmer line gradients roughly agree with each other, except at the outermost regions where H/3 appears to suffer from emission line fill-in (it flattens whereas the higher-order Balmer lines rise to the edge). This conjecture is further supported by the fact that H/3 extends beyond the grids (Fig. 4.31) at largish radii. This galaxy does not appear to have super-solar [a/Fe] at any radii but, again, C N 2 is very high in the inner region possibly implying a N enhanced population. Removing all Balmer-line indices from the fits (Case B) predicts larger ages at many radii. From the index deviation plot in Figure 4.40 [left panel] we see that, at r = 0", none of the Balmer lines are well-fit in this case. Since there was no evidence for emission in the central regions, the other Cases should be more realistic. In fact, even Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 224 though there is no a-enhancement evident, it is Case D that seems to have the most stable fits. If we follow Case D (because it will not suffer as much from the indices such as C N 2 that extend far off the model grids), the fit is roughly consistent with being very young (~ 1-3 Gyr) with little overall gradient, leaning to slightly younger ages at 1 rd. The metallicity is ~ ZQ in the center, decreasing to Z ~ 0.004-0.006 at rd (or roughly 0.7dex). There is one prominant spike to old age and low Z at about r~20-30". One interpretation could be that the slit is crossing an inter-arm region such that we are only seeing the underlying/thick disk population, which would be expected to be older and at lower Z than the currently SF regions. However, many of the outer points lie in the degenerate region of the grid, which could also result in these age jumps, in addition to possible emission-line fill in. U G C 2124 - SB(r)a This is a very strongly barred galaxy with our earliest-type bulge classification and the slit was aligned along the bar. There is very little emission in this galaxy. We only detect one prominent H II region at ~38" that gets masked (see top frame of Fig. 4.9). The data go out to 0.7rd (~24"), and the bar length, according to de Jong (1996) is ~20", so we reach just beyond the edge of the bar. The SB drops off very quickly beyond the bar, so there just isn't enough signal there to acquire sufficient S/N. The molecular TiO and Mg indices were excluded from all fits due to sky lines in the former, and gap location for the latter. The deviation of individual indices, Figure 4.41 [left panel], demonstrates that the TiO and molecular Mg indices are indeed poorly fit. From Figure 4.23, we see that the index profiles are all very symmetric. Metal-line indices show only a mild (< — 0.1 dex) decrease out to ~ rd where, just beyond the bar, the indices jump to lower values. The Balmer lines do not agree as well. The two H 7 indices agree and show a significant increase to the edge of the data. H/3 & B.8p, on the other hand, remain mostly flat (with a small decrease at the edge). This trend in H/3 could be due to emission-line fill in, but the Hc5p cannot. It is more likely compromised due to being at the edge of our wavelength calibration. This galaxy appears to be mildly a-enhanced in the central regions (see Fig. 4.38), Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 225 closer to elliptical galaxies which is constistent with it being our earliest-type bulge under the assumption that bigger, earlier-type, bulges are more akin to elliptical galaxies in terms of their formation. The Case B fit for which all Balmer lines are excluded predicts much older ages and lower Z than the others in the central regions, and only the H/3 Balmer line is well-fit at r = 0". Since there was no obvious emission in this galaxy, and all other fits roughly agree, we consider them to be more realistic, thus indicating that this galaxy is fairly young (~4-6 Gyr) and metal rich (z ~ 0.04) at the center. Contrary to the fiat profiles seen in the individual metal-line indices, the best-fit metallicity appears to show a significant gradient, from ~ 0.03 -0.04 at the center, to ~ 0.004 just past the bar's edge, and it is most pronounced for Case D where the a/Ye ratio is taken into account. However, this decrease is dominated by a jump at beyond which the profile is indeed quite flat. The age profile is odd in that, just beyond 1", it oscillates between ages of ~ 5 -20 Gyr for all four fits. This could again be a case of the age degeneracy of the models, and it is entirely possible that the age is roughly constant at ~ 4 - 8 Gyr throughout. The biggest mystery surrounding this galaxy is its colour gradients (Fig. 4.47). The colours all lie well off the grids (redward in optical colour). The central point could very well be affected by seeing mismatches between the photometric observations, but the rest of the profile still lies redward of the grids. Part of this reddening could be caused by dust, but most dust models will tend to move the data along lines roughly parallel to lines of constant age in this region of the grid (see dust vectors in Fig. 2.12), and thus would not help in shifting the data to within the model grids. Many other sources of error in the photometry could present themselves, but we could not identify any that would be severe enough to cause errors of the order needed. Remarkably, between roughly r~1.5 and 8", the metallicity derived from the colours (Fig. 4.52) lies intermediate between the four index fit cases, but this should be taken with care given that the colours are not on the model grids. Again, though, the metallicity profile is extremely flat. The fact that the metallicity profile appears much smoother is Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 226 likely because the SB profiles are azimuthally averaged, thus the features we see where galaxian structure crosses the slit gets effectively smoothed out. The age, as expected from the grids, is only consistent with a saturated 20 Gyr at all radii. As of this writing, we have no clear explanation for the extremely red optical colours of this galaxy. Repeat photometric observations would be extremely useful to rule out photometric errors. N G C 7490 - Sbc This is a normal late-type face-on spiral galaxy with fluffy spiral arms. Our best-fit Sersic index is n = 2.35 which is on the high end for such a late-type classification. We detect a large number of H II regions, indicating significant current SF, including a large spike right in the center which does not get masked. The spectra (Fig. 4.15) show weak signs of emission in [N II] with discernable fill-in of strong underlying Ha absorbtion. Otherwise, the SEDs appear quite red in the center, and blueing as radius increases. Again, sky subtraction issues appear beyond > 6" thus all sky-line affected indices are excluded from the fits (Ti0 2 and Fe5782). TiOi was also excluded from all fits due to a gap. Indeed these indices are poorly fit at all radii (as seen in the deviation plots in Fig. 4.42). H/3 also falls in one of the gaps, but it was still included in the Case A fit, and was equally well fit in that case as the others which exclude H/3. Once again, the radial index profiles are very symmetric (Fig. 4.24), and significant negative gradients are seen in most metal-line indices (of order —0.2 to —0.6 dex from center to Ira). Mgi does not decrease as strongly as Mg 2 or Mgb but it is not clear why this is the case. It also systematically has a greater devation in the fits for almost all cases. The age indicators H7 k. G4300 agree with each other but, again, H/3 & B.5F deviate at the outer regions The CN and Fe4668 (which measures C 2) indices are way above the grids in the central regions and are clearly not well fit in the deviation plot (Fig. 4.42) for the central r — 0" point. This galaxy shows an indication of [a/Fe]-enhancement at intermediate radii in the (Fe) -Mgb plane (Fig. 4.38), but not into the disk, which could imply shorter formation timescales for the intermediate regions. Following Case D, the age gradient is roughly Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 227 consistent with being fairly shallow, decreasing from 4 Gyr in the center to ~ 3 Gyr at rd, but there are jumps to older ages between re and rd. The metallicity is very high in the center, decreasing to very low values (Z ~ 0.002) at rd (~1.3 dex). However, it appears that the gradient is steeper in the bulge region and flattening out into the disk. Colour gradients for this galaxy are displayed in Figure 4.48. Again the central region is quite red in optical colours. Beyond ~ r e, the colour profile is within the grids, but right at the edge. As such, the age estimates from the colours (Fig. 4.53) saturate to 20 Gyr out to ~rd, in stark contrast to the index age fits around 3 Gyr. Again the metallicity fits from the colours are in much better agreement with those from the indices. The metallicity decreases less rapidly in the case of the colour fits, and seems to lie in between the "spikes" seen in some of the index fits. This again could be a consequence of the azimuthal smoothing. It would be interesting to obtain SB cuts along the same position angle as in the spectroscopic measurements to confirm if this is indeed the case. N G C 7495 - SAB(s)c This is another normal late-type face-on spiral galaxy with fluffy, poorly-defined spiral arms. Our best-fit Sersic index of n = 4 does not agree well with the late-type classification. There is so much Ha-I- [N II] emission from the center out to ~35", that most of the disk gets masked out, and the bulge points show large emission spikes in all four Balmer lines (Fig. 4.16). As a result, we only have 5 data points for this galaxy, going out to just 0.3 rd, and all of them are affected by emission line fill-in. The SEDs are quite blue suggesting a substantial contribution from a blue continuum source. All sky-line affected indices are excluded from the fits (Ti0 2 and Fe5782), as well as TiOi and both molecular Mg indices as they fall within gap regions. The grid plots (Fig. 4.34) echo the severe emission-line contamination with all Balmer-line indices extending off the grids, but Figure 4.38 of (Fe) versus Mg6 demon-strates that at least some of the metal lines lie within the model grids. Thus, age and metallicity fits can be attempted with Case B that removes all indices severely affected by emission lines, although we note that a strong continuum from the emission source Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 228 would dillute the metal-line indices leading to lower values. As expected, all Balmer-line indices are poorly fit at all radii (see deviation plots in Fig. 4.43), and those fits including Balmer-lines all predict old ages (> 10 Gyr), while the Case B finds much younger ages (~1.5 Gyr). Also as expected, the metallicities are quite low for all cases. Given that we only have 2 points outside the seeing disk, radial variations cannot be constrained. Colour gradients for this galaxy are displayed in Figure 4.49. In this case, the central regions within ~ re are quite blue in optical colours, in agreement with the blue SEDs we observed in the spectra (Fig. 4.16), but not in the optical-IR colours (except perhaps the central point which could be affected by seeing mismatch). So the continuum is coming from a source that has blue optical colours, but normal optical-IR colours. Beyond r e, the disk seems to follow the line of constant solar metallicity with decreasing age out to r^, beyond which metallicity drops, while age is roughly constant. This trend is reflected in the fits shown in Figure 4.53. Comparisson with the index fits is difficult due to the issues noted above, but it does appear that the higher Z predicted by the colours is consistent with our interpretation of continuum dillution of the metal-line indices, and recalling that it is the (seemingly normal) optical-IR colours that drive the metallicity determination in the colours. N G C 7610 - Scd This late-type spiral galaxy has two very prominant spiral arms, but appears as-symetric and is seen at an inclination of ~41°. The Ha + [N II] emission echoes this, having many more H II regions on one side of the galaxy. There is also a strong Ha: emission spike at the very center (Fig. 4.10), which does not get masked out. There are a few H II regions that survive the masking cut, thus it is not surprising that the spectra (Fig. 4.17) show signs of emission throughout, as well as having a very blue SED even in the central (unmasked) region. Sky subtraction of strong lines was problematic at all radii (due in large part to its low SB) and all affected indices were excluded from the fits (both TiO indices and NaD). The two molecular Mg indices were also excluded from all fits due to the gap. Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 229 Clear trends are hard to discern in the index gradients (Fig. 4.26), which only go out to 0.55 r^ and stronger assymetries are seen than for the previous galaxies, but only for certain indices (e.g. Ca4227 is not symmetric, but Fe5270 is). In some cases, indices that are supposedly measuring the same elements do not agree with each other (e.g. Fe5270 is roughly flat while Fe5335 has a negative gradient). From the index-index grid plots (Fig. 4.35) we see that H/3 is clearly affected by emission line fill-in, thus the Case A should not be trusted. K6F also seems to disagree with H7.4 which could be due to compromised wavelength calibration at H<5, thus Case C might not be useful either. Indeed, the different fits for this galaxy (Fig. 4.44) give very different results. Comparing the Case B & D fits implies that all of the Balmer lines were emission-line contaminated, leaving just fit B to consider. Indeed for Case B, all Balmer-line indices are poorly fit at all radii, as is the emission affected C N i , but the age-sensitive G4300 is well reproduced. Case B implies a young, slightly increasing age of ~ 0.8 -1.5 Gyr out to ~7r* e (11"), and a negative gradient in metallicity that could be as large as 2.2dex, but is also consistent with being a milder ~ 0.6 dex. Colour gradients for this galaxy are displayed in Figure 4.50. Again we have a case where the colours lie redward of the model grids in optical colours, but have normal optical—IR colours and it is not clear why this is the case as dust reddening does not affect the optical colours as much. Thus, again, comparisson with the index fits (Fig. 4.55) is difficult due to the difficulties noted above, but it does appear that the colours predict a slightly higher Z, which could easily be caused by dust extinction. The colours also predict that the negative Z gradient continues well into the disk. N G C 7741 - S c d This late-type barred spiral has a prominent bar and two short spiral arms. There are bright H II regions in the bar and arms and we have alligned the slit perpendicular to the bar (see Fig. 4.4). This galaxy has the biggest and brightest central emission region (see Fig. 4.11), and there is weaker emission throughout the disk. The central starburst could be a result of gas funnelling to the center due to the bar. The spectra (Fig. 4.18) show evidence of emission throughout, particularly at the center where strong [0 III] Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 230 emission is seen in addition to the Balmer and [N II] lines, and the SED is quite blue. The sky-line affected indices (both TiO and NaD) and the gap affected indices (both molecular Mg and Fe5015) were excluded from all fits. The index gradient trends are quite scattered for this galaxy (Fig. 4.27), particularly the bluest indices (e.g. Ca4227, G4300, Fe4668). The redder indices are slightly better measured, but show deviating trends (e.g. Fe5270 is flat on one side, rises on the other, while Fe5335 decreases on both sides), and the Balmer lines are completely asymmetric. Accordingly, the data lie well off the model grids with Balmer-lines (Fig. 4.36), but are within the grid regions in the (Fe) versus Mgtb plot (Fig. 4.38). Thus the only fits we can consider are those of Case B (Fig. 4.45). These are consistent with young and roughly constant age (save a "blip" to very young at ~2-3") of < 1 Gyr, possibly decreasing to even younger age past rd. The metallicity is sub-solar and rising gradually from a central Z ~ 0.004 to Z ~ 0.008 at rd (~ +0.3 dex). Once again reflecting the blue SED (Fig. 4.18), the central regions within < rd are quite blue in optical colours but not in the optical-IR indices (Fig. 4.51). The colour fits in this galaxy (Fig. 4.56) are remarkably consistent with those from Case B for the indices. The gradients are much smoother, likely reflecting the azimuthal averaging, and agree with the constant young age, but the metallicity appears to have a mild negative gradient (down to Z ~ 0.001, or ~—0.7dex from the center) that is more pronounced in the colour fits. IC 239 - Scd This late-type spiral is our lowest surface brightness galaxy. It was also taken in the brightest and most variable sky conditions (see Fig. 4.5). However, since there were fewer regions of high Ha-f- [N II] emission, we were able to extract enough S/N to get data out to ~22" (which is only 0.6rd for this galaxy). Only NaD and Fe5015 were excluded from all fits, but the TiO indices are poorly fit and perhaps should have been excluded as well. Many of the index profiles are rising for this galaxy (Fig. 4.28), even the metal-line indices. In the index grid plots (Fig. 4.37), there seem to be discrepancies among „ Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 231 different indicators, especially at the largest radii. This likely is a result of residual sky "features" due to the overwhelming sky for these observations. However, the fits for all 4 cases are quite consistent with each other (Fig. 4.46) and are consistent with a decreasing age, from ~ 3 - 7 Gyr in the center to ~ 1 Gyr at ~ 2 re (or a roughly constant age if the first data point is ignored), and a decreasing metallicity from (ignoring the first point) Z~0.03 to 0.004 at ~ 2 r e (or ~-1.5 dex). 4.6.1 Summary To summarize, we present in Table 4.12 our best fit light-weighted SSP ages and metal-licities for all galaxies at three radii: the centre, re, and r^. The general observation is that these galaxies appear to be very young throughout with very little age gradient (but it is negative, i.e. younger outer regions, when detected) and with mostly high central metallicities with significant negative gradients. These trends are certainly not Name ;e (Gyr) AA/(r/rd) Z Alog 1 0 Z/(r / r d ) r ~ 0 re (Gyr) r ~ 0 re rd (dex) N0173 3.0 20 2.5 -0.5 0.03 0.006 0.003 -1 N0628 2.5 2.0 1.2 -1.3 0.02 0.008 0.004 -0.7 U2124 6.0 6.0 6.0 0 0.03 0.01 0.004 -0.9 N7490 4.0 4.0 4.0 0 0.04 0.02 0.002 -1.3 N7495 1.5 1.5 N / A N / A 0.004 0.002 N / A N / A . N7610 0.8 1.5 N / A N / A 0.006 0.003 N / A . N / A N7741 1.0 1.0 1.0 0 0.004 0.006 0.008 +0.3 10239 2.0 1.5 N / A N / A 0.02 0.01 N / A N / A Table 4.12: Galaxy Ages and Metallicities from Fits consistent with formation scenarios that predict old bulges (monolithic) or that bulges are much older than their disks (hierarchical without secular evolution). The constant age profile is very indicative of secular processes, although disk contamination could be partially responsible, but the large metallicity gradients are difficult to interpret in this scenario unless the star formation efficiency increases towards the central regions. It should again be emphasized that the measurements here are light-weighted and thus Chapter 4. Spectroscopic Age and Metallicity Gradients in Spiral Galaxies 232 sensitive to the most recent episode of star formation. Thus the young ages implied by the indices do not imply that no old age stellar populations exist. An underlying old age component could be present, but is masked by the significant amount of recent star formation that has occured in these galaxies. The anomalous ages determined from the colour gradients for a few of our galaxies (U2124, N7490, & N7610 in particular) should not be taken as a major problem as, in hindsight, we discovered that by bad luck our galaxies were highlighted as having problematic photometric calibration in the BdJOO sample (E. Bell 2005, private com-munication). The source has not yet been tracked down, but it seems most likely that it is the optical calibration that is amiss given the good agreement in the NIR-sensitive metallicity measurements between the colours and indices. As such, we reserve any fur-ther discussion on the comparison between spectroscopic and broadband measurements until a properly calibrated and homogeneous set of photometry is obtained for these galaxies. Chapter 5 Conclusions and Out look 233 We have embarked on a long-term study of the large-scale chemical and evolutionary properties of nearby galaxies. The main objectives of this program are to produce a comprehensive study of the relative age and metallicity distributions in bulges and inner disks and their systematic behavior with galaxy type, bar structure, and environment, and to constrain basic theories of bulge formation (see Chapter 1). The ongoing star formation and the mixing of the gas and stellar population by bars and spiral arms make the interpretation of age and metallicity gradients in spiral galaxies much less clear than for ellipticals. An extensive database is needed to map the many parameters that influence the formation history of galaxies, and provide further insight into the origins of bulges as a function of bulge-to-disk ratio and galaxy magnitude. In Chapter 2, we presented an extensive study of spiral galaxy stellar populations using radial colour profiles. We found that bright spiral galaxies exhibit clear trends for the effective age and metallicity in as a function of surface brightness (SB), lumi-nosity, rotational velocity, and size. Based on their integrated light, higher SB regions of galaxies formed their stars earlier than lower SB ones, or at a similar epoch but on shorter timescales. Also, the SFHs at a given SB level, which lead to age gradients, are modulated by the overall potential of the galaxy, such that brighter/higher rotational velocity galaxies formed earlier. These trends are of course at odds with hierarchical galaxy formation which predicts that more massive galaxies form late. While feed-back processes are often invoked as a possible mechanism to prevent the gas in lower mass galaxies from cooling and forming stars at early times, no viable prescription for feedback has been found thus far that agrees with all observational constraints. An im-portant limitation of our broadband colour gradients analysis is its sensitivity to dust extinction which could mimic (or hide) gradients in age and, particularly, metallicity. Chapter 5. Conclusions and Outlook 234 In order to alleviate concerns that our colour gradients could be affected by dust reddening, we designed a similar spectroscopic investigation. To further reassure our-selves of the relative dust insensitivity of line-indices, we carried out a qualitative study in Chapter 3 which highlighted a few potential dangers in this assumption (especially concerning the 4000 A break and very young composite stellar populations). However, for the most part, this study demonstrated that dust will not play a major role in the determination of ages and metallicities using absorption-line indices. Hence, to address the many astrophysical goals outlined in the introduction (Chap-ter 1), we obtained longslit spectroscopy for a pilot sample of 8 nearby late-type field spirals from which we measure a suite of absorption line-indices in their bulges and inner disks. Our pilot Gemini/GMOS study of nearby face-on spirals reveals that the bulges and inner disks (out to ~ 1 rd) of late-type spiral galaxies are generally consistent with having young light-weighted SSP ages (< 1 to 6 Gyr), with the latest-types having the youngest ages. For the majority of our galaxies, the age profiles are relatively flat, or slightly decreasing to younger ages at larger radii. The metallicities are solar or above in the centres of five of the galaxies, while the other three have central Z ~ 0.004. Most galaxies show large metallicity gradients on the order of —0.3 to —0.7 dex per rd, in the sense that they are more metal rich in their central regions. Three out of the eight galaxies show mild [a/Fe] abundance enhancements, such as those commonly observed in elliptical galaxies and early-type bulges. This is indicative of short star formation timescales, which are not exclusively confined to the bulge region. Four of the galaxies show stronger CN values in their central regions than can be accommodated by the models. Such CN enhancements have been observed before in globular clusters and el-liptical galaxies, and are likely due to an enhancement in nitrogen (rather than carbon). The above trends are certainly not consistent with formation scenarios that predict old bulges (monolithic) or that bulges are much older than their disks (hierarchical without secular evolution). The flat age profile is indicative of secular processes, although disk contamination could be partially responsible, but the large metallicity gradients are dif-ficult to interpret in this scenario unless the star formation efficiency increases towards Chapter 5. Conclusions and Outlook 235 the central regions. These findings are consistent with the study of Proctor & Sansom (2002) who also rule out primordial collapse models for late-type bulge formation, but contrast the results of Goudfooij et al. (1999) who find uniformly old bulge ages with super-solar abundance ratios. Our study is the first to obtain radially resolved spectra well into the galaxy disks, enabling a comparison of the bulge and inner disk populations. We do not see obvious discontinuities between the bulge and inner disk (except perhaps in the case of N7490, one of our earliest-type galaxies), as would be expected in non-secular bulge formation. To assert the absence of a dichotomy between the bulge and inner disk, however, requires the careful removal of any disk contamination into the bulge. This is indeed a major focus of our continued investigation of this rich and complex data set. While the small size of our sample does not allow for a detailed statistical study of galaxy parameters with inferred age and metallicity, several important observations emerge that can guide future studies. The need for high S/N well into the disk, and especially the inter-arm regions, is absolutely crucial and realistically feasible only with the largest aperture telescopes (otherwise prohibitively long exposures would be re-quired). Accurate sky subtraction, ideally measured from the slit edges, are also of the essence for absorption-line spectroscopy. In future telescope proposals, darker skies will be requested along with longer integrations for dithered sky observations (for at least one-third of the science target) when absolutely needed for large diameter galaxies. The longslit technique used is quite adequate for such a study, but full 2D mapping from an integral field unit with a large enough field of view and sufficient spectral coverage would be ideal. This analysis has emphasized the advantage of having a large number of the Lick indices to consider simultaneously when translating from the observed to the physical age and metallicity plane via SPS models. In particular, the age indicators at the blue end of the spectrum ( H 7 , H<5, & G4300) are crucial, and the existence of abundance anomalies of individual metal-line indices (Fe, Mg, CN) require combinations of indices or for certain indices to be omitted altogether. This holds true for any stellar system, Chapter 5. Conclusions and Outlook 236 but is of particular importance for objects with significant emission and/or non-solar abundance ratios (although the latter can be accounted for with models that take such ratios into account). Future studies of spiral galaxies must consider this when designing the observing strategy. A major impediment of the longslit spectroscopy technique is the significant coverage loss due to the Ff II regions that overwhelm the underlying spectrum. Two-dimensional mapping with an integral field spectrograph with large enough field of view and suffi-cient spectral coverage would be ideal for this sort of investigation, but exposure times remain prohibitive and longslit spectroscopy is still attractive. Tunable filters enable the measurement of selected indices through a collection of narrow-band "images" on and around any chosen feature (e.g. Ryder, Fenner, & Gibson 2005). Having full in-dex mapping across an entire galaxy is a most powerful approach to measuring spatial gradients, but the requirement of at least two to three low-luminosity integrations per index and the need for multiple indices to achieve robust age and metallicity diagnostics make for a challenging and taxing technique. An important (and future) application of our spectroscopic analysis is the calibra-tion of techniques for evolutionary studies of high-redshift galaxies which rely largely on colour information. Our analysis of local, well resolved, galaxies enables us to de-termine the limitations and reliability of the colour-based techniques to map out stellar populations of distant galaxies (e.g. whether dust biases the results or not). This ex-ploration is beyond the scope of the present work and a larger well-calibrated database with extended photometric and spectroscopic coverage will also be needed for improved statistical confidence. In our ongoing investigation, we also plan to use SB cuts (rather than azimuthally averaged profiles) along the same position angle as that of the spectrograph slit, to com-pare photometric and spectroscopic features directly. This comparison will confirm the degree to which line-index variations are associated with physical (e.g. arm/inter-arm) features. Related to this, much of the excised H II regions also contain significant gas-phase information that can be studied using the diagnostics such as the [N II] A6584/Ho: Chapter 5. Conclusions and Outlook 237 metallicity (Storchi-Bergmann, Calzetti,& Kinney 1994) and EW(Hai) age indicators1. A proper calibration of the indices on to the Lick system using our standard star observations will also be performed to enable direct (model-independent) comparison with other studies of ellipticals (e.g. Trager et al. 2000) and spirals (e.g. Proctor & Sansom 2000). Given some of the remaining degeneracies in the absorption-line index grids, stellar population parameters will be derived using a number of different template synthesis techniques, all designed to overcome the effects of emission-line contamination. A first technique involves the linear, non-zero, combination of stellar templates, but masking out emission-line regions, observed with the same instrumental set-up. The galaxy spec-trum is then divided by the best-fit template and, if the template match is good, this will highlight the emission lines (not present in the stellar templates). The equivalent widths of the emission lines can then be measured and used to correct the absorption EWs. A second template synthesis technique uses SPS model spectra as templates. Best-fit SSP and mixed populations can be considered in conjunction with dust red-dening (assuming a dust model such as the one described in Chapter 3), and power-law AGN-type contributions. As an example of work in progress and the power of this technique, we show in Figure 5.1 the best fit SSP model spectrum (red line) based on our index fits compared to the galaxy spectrum (black lines) for the emission free r = 0" bin in NGC 628 [top 3 panels] as well as the severely emission contaminated r = 0" bin of NGC 7495 [bottom 3 panels] (the three panels in each plot divides the wavelength coverage into three intervals). The model-data match for NGC 628 is quite good (and reassuring), recalling that only selected narrow wavelength ranges are included in the fit. The fit for NGC 7495 only included indices that are not affected by the emission lines. In the ~ 5100 -5800 A interval, the match is quite remarkable. However, the galaxy spectrum is clearly bluer than the model in the ~ 4000 -4600 A region, which could easily result from a small contribution of a very young (^0.1 Gyr) stellar pop-ulation. Finally, we can also attempt to put limits on H II region contamination to 1Bergmann, Jorgensen, & H i l l (2003) performed such an analysis for a sample of low SB galaxies. Chapter 5. Conclusions and Outlook 238 V 1.2 1.1 1 0.9 0.8 0.7 1.2 1.1 1 0.9 0.8 0.7 4100 4200 4300 4400 4500 4600 4700 4800 4900 - I — | — i — i — i — [ — i — i — i — | — i — i — i — | — i — i — i — | — i — i — i — | — i — i — i — | — i — i — i — | — i — i — i — | — r _l 1 i i i I i i l I i i i I i i i I i i i I i i i I i i i I i i i I i_ 5000 5100 5200 5300 5400 5500 5600 5700 5800 1 ] I 1 I 1 I I I 1 1 1 1 T- i—i—|—i—i—i—|—i—i—i—I—i—i—i—I—r • Ana — OR r.iir V 4 fL Age = 2.5 Gyr Z = 0.02 j 1 i i i 1 i i i I i i i 1 i i i L _i i i i i i_i i i i i i_ 5900 6000 6100 6200 6300 6400 6500 6600 6700 A(A) N0628 I . i i i . . . i . . . i . . . i . . . i . . . i . . . i . 4100 4200 4300 4400 4500 4600 4700 4800 4900 A 1.2 O 5 1.1 s i 1 o5 0.9 3 V 0.8 J \ 0.7 < — 1.2 1.1 1 0.9 0.8 0.7 ' I 1 1 1 I 1 1 1 I 1 1 1 I 1 1 1 I 1 1 ' I 1 1 1 I 1 1 ' I 1 1 1 I J ^ ^ ^ ^ ^ _J I I I I I ! I 1 I I I ! 1 I I I I I I I I I I l_ 5000 5100 5200 5300 5400 5500 5600 5700 5800 Age = 1 . 5 Gyr  I- Z = 0.004 5900 6000 6100 6200 6300 6400 6500 6600 6700 X(A) N7495 Figure 5.1: Best fit SSP model (red line) based on index fits compared to the data (black line) for the r = 0" bins of NGC 628 [top 3 panels] and NGC 7495 [bottom 3 panels]. Chapter 5. Conclusions and Outlook 239 the underlying stellar population using an empirical H II region template (from the galaxy itself), and subtract off increasing fractions to map its effects. These techniques should permit further analysis of the underlying signal that is currently lost due to the overwhelming presence of H II regions. A major limitation in translating from index measurements to light-weighted ages and metallicities is that the models turn over on themselves (i.e. become degenerate) at low metallicities and older ages, such that a given datum can be consistent with being either very old (20 Gyr) or relatively young (< 8 Gyr). This problem could be alleviated if some other age indicator is available that can put a reasonable upper (> 8 Gyr) or lower (^8 Gyr) limit on the age (the dividing line depends on metallicity and which indices are being considered). A few possible indicators include the 4000 A break, the high-resolution H 7 index of Jones & Wothey (1995), and optical-UV colours. A variety of improvements to the models are underway. For example, Schiavon, Rose, & collaborators are currently working on revisions of the Lick/IDS spectral indices us-ing a new spectral library of ~ 1100 stars, including new high metallicity stars, at ~ 2 A resolution over the range 3500-7500A, and stellar atmospheric parameters. In collab-oration with Schiavon, Rose, & Courteau, we have measured very high S/N spectra (~150/A) of Galactic globular clusters to calibrate high-resolution stellar population synthesis models. The addition of clusters with varying metallicities and horizontal branch morphologies will significantly improve the accuracy and predicting power of current SPS models with variable element abundance ratios (e.g. Schiavon et al. 2003; Thomas, Maraston, & Bender 2003). Many modelers are also extending their model predictions into the UV in search of better age indicators (e.g. Peterson et al. 2005), which would be of great use for future studies of spiral galaxy populations. Our primary goal, namely to demonstrate the applicability of longslit spectroscopy to extract age and metallicity sensitive indices in star-forming spiral galaxies and ul-timately link them with SPS models, has been achieved. Considerable analysis is still required before definitive results can be reached, but our study calls for optimism. Ul-timately, we desire a large enough sample to look for trends with galaxy parameters. Chapter 5. Conclusions and Outlook 240 More data have also been obtained at Gemini South using a similar set-up (though further optimised for sensitivity at blue wavelengths), and will provide much needed coverage of the age-sensitive 4000 A break. These data will be reduced over the course of the coming year and will further our ability to establish empirical trends between the spectral features and global galaxy parameters. The majority of those galaxies, and most likely future targets as well, are preferentially selected out of the SINGS (Spitzer Infrared Nearby Galaxies Survey; see http://sings.stsci.edu/) data base which to pro-vide complete gas-phase and dust information in addition to our own stellar and H II region data. The future of stellar population studies in late-type spiral bulges and inner disks still holds the realistic promise of discriminating between different scenarios for the formation and evolution of galaxies. This will not come, however, without concerted efforts and dedicated telescope resources. 241 Bibl iography [I] Abraham, R. G., Ellis, R. S., Fabian, A. C., Tanvir, N. R., and Glazebrook, K. 1999, Mon. Not. Roy. Astro. Soc. 303, 641 [2] Anders, R, Bissantz, N. , Fritze-v. Alvensleben, U., and de Grijs, R. 2004, Mon. Not. Roy. Astro. Soc. 347, 196 [3] Andersen, D. R., Bershady, M . A., Sparke, L. S., Gallagher, J. S., and Wilcots, E. M. 2001, Astrophys. J. Let. 551, L131 [4] Armandroff, T. E. and Zinn, R. 1988, Astron. J. 96, 92 [5] Balcells, M. , Graham, A. W., Dominguez-Palmero, L., and Peletier, R. F. 2003, Astrophys. J. Let. 582, L79 [6] Balcells, M. , Graham, A. W., and Peletier, R. F. 2005, ApJ submitted, astro-ph/0404379 [7] Balogh, M. L., Morris, S. L., Yee, H. K. C., Carlberg, R. G., and Ellingson, E. 1999, Astrophys. J. 527, 54 [8] Bardeen, J. M. 1980, Phys. Rev. D 22, 1882 [9] Barnes, J. E. and Hernquist, L. 1996, Astrophys. J. 471, 115 [10] Baugh, C. M. , Lacey, C. G., Frenk, C. S., Granato, G. L., Silva, L., Bressan, A., Benson, A. J., and Cole, S. 2005, Mon. Not. Roy. Astro. Soc. 356, 1191 [II] Bell, E. F. and Bower, R. G. 2000, Mon. Not. Roy. Astro. Soc. 319, 235 [12] Bell, E. F. and de Jong, R. S. 2000, Mon. Not. Roy. Astro. Soc. 312, 497, [BdJOO] Bibliography 242 [13] Bell, E. F. and de Jong, R. S. 2001, Astrophys. J. 550, 212 [14] Bell, E. F., Mcintosh, D. H., Katz, N., and Weinberg, M. D. 2003, Astrophys. J. Supp. Ser. 149, 289 [15] Bergmann, M . P., Jorgensen, I., and Hill, G. J. 2003, Astron. J. 125, 116 [16] Bertelli, G., Bressan, A., Chiosi, C., Fagotto, F., and Nasi, E. 1994, Astron. Astrophys. Suppl. Ser. 106, 275 [17] Blakeslee, J. P., Vazdekis, A., and Ajhar, E. A. 2001, Mon. Not. Roy. Astro. Soc. 320, 193 [18] Bottinelli, L., Gouguenheim, L., Fouque, P., and Paturel, G. 1990, Astron. Astro-phys. Suppl. Ser. 82, 391 [19] Bruzual, G. and Chariot, S. 2003, Mon. Not. Roy. Astro. Soc. 344, 1000, [BC03, GALAXEV] [20] Bruzual A., G. 1983, Astrophys. J. 273, 105 [21] Bruzual A., G. and Chariot, S. 1993, Astrophys. J. 405, 538 [22] Bruzual A., G., Magris, G., and Calvet, N. 1988, Astrophys. J. 333, 673 [23] Burstein, D. and Heiles, C. 1984, Astrophys. J. Supp. Ser. 54, 33 [24] Byun, Y. I., Freeman, K. C , and Kylafis, N. D. 1994, Astrophys. J. 432, 114 [25] Caldwell, N., Rose, J. A., and Concannon, K. D. 2003, Astron. J. 125, 2891 [26] Calzetti, D. 2001, Pub. Astro. Soc. Pac. in press (astroph/0109035) [27] Cardiel, N. , Gorgas, J., Cenarro, J., and Gonzalez, J. J. 1998, Astron. Astrophys. Suppl. Ser. 127, 597 [28] Cardiel, N., Gorgas, J., Sanchez-Blazquez, P., Cenarro, A. J., Pedraz, S., Bruzual, G., and Klement, J. 2003, Astron. Astrophys. 409, 511 Bibliography 243 [29] Carlberg, R. G. 1984, Astrophys. J. 286, 403 [30] Carollo, C. M. 1999, Astrophys. J. 523, 566 [31] Cenarro, A. J., Cardiel, N., Gorgas, J., Peletier, R. F., Vazdekis, A., and Prada, F. 2001a, Mon. Not. Roy. Astro. Soc. 326, 959 [32] Cenarro, A. J., Gorgas, J., Cardiel, N., Pedraz, S., Peletier, R. F., and Vazdekis, A. 2001b, Mon. Not. Roy. Astro. Soc. 326, 981 [33] Chabrier, G. 2003, Pub. Astro. Soc. Pac. 115, 763 [34] Chariot, S. and Fall, S. M. 2000, Astrophys. J. 539, 718, [CF00] [35] Chariot, S., Kauffmann, G., Longhetti, M. , Tresse, L., White, S. D. M. , Maddox, S. J., and Fall, S. M. 2002, Mon. Not. Roy. Astro. Soc. 330, 876 [36] Chariot, S., Worthey, G., and Bressan, A. 1996, Astrophys. J. 457, 625 [37] Cline, D. B. 2004, in AIP Conf. Proc. 743: The New Cosmology: Conference on Strings and Cosmology, pp 41-56 [38] Cole, S., Lacey, C. G., Baugh, C. M. , and Frenk, C. S. 2000, Mon. Not. Roy. Astro. Soc. 319, 168 [39] Combes, F. 1998, in Astronomical Society of the Pacific Conference Series, pp 300-+ [40] Courteau, S. 1992, Ph.D. Thesis [41] Courteau, S. 1997, Astron. J. 114, 2402 [42] Courteau, S., de Jong, R. S., and Broeils, A. H. 1996, Astrophys. J. Let. 457, L73+ [43] Dalcanton, J. J. and Bernstein, R. A. 2002, Astron. J. 124, 1328 Bibliography 244 [44] Dalcanton, J. J., Spergel, D. N., and Summers, F. J. 1997, Astrophys. J. 482, 659+ [45] de Blok, W. J. G. and Bosma, A. 2002, Astron. Astrophys. 385, 816 [46] de Blok, W. J. G., McGaugh, S. S., and Rubin, V. C. 2001, Astron. J. 122, 2396 [47] de Jong, R. S. 1996, Astron. Astrophys. Suppl. Ser. 118, 557 [48] de Sitter, W. 1930, Bull. Astron. Inst. Neth 5, 211+ [49] de Zeeuw, P. T., Bureau, M. , Emsellem, E., Bacon, R., Marcella Carollo, C., Copin, Y., Davies, R. L., Kuntschner, H., Miller, B. W., Monnet, G., Peletier, R. F., and Verolme, E. K. 2002, Mon. Not. Roy. Astro. Soc. 329, 513 [50] Delisle, S. and Hardy, E. 1992, Astron. J. 103, 711 [51] Diaz, A. I., Terlevich, E., and Terlevich, R. 1989, Mon. Not. Roy. Astro. Soc. 239, 325, [DTT] [52] Disney, M. , Davies, J., and Phillipps, S. 1989, Mon. Not. Roy. Astro. Soc. 239, 939 [53] Eggen, O. J., Lynden-Bell, D., and Sandage, A. R. 1962, Astrophys. J. 136, 748+ [54] Emsellem, E., Davies, R., McDermid, R., Kuntschner, H., Peletier, R., Bacon, R., Bureau, M. , Cappellari, M. , Copin, Y. , Miller, B., Verolme, E., and de Zeeuw, T. 2002, in ASP Conf. Ser. 282: Galaxies: the Third Dimension, pp 189-+ [55] Evans, R. 1994, Mon. Not. Roy. Astro. Soc. 266, 511 [56] Falcon-Barroso, J., Peletier, R. F., Vazdekis, A., and Balcells, M. 2003, Astrophys. J. Let. 588, L17 [57] Fall, S. M. and Efstathiou, G. 1980, Mon. Not. Roy. Astro. Soc. 193, 189 [58] Ferland, G. J. 2002, Hazy, a Brief Introduction to CLOUDY, University of Ken-tucky, Dept. of Physics & Astronomy, Internal Report, [CLOUDY] Bibliography 245 [59] Fioc, M. and Rocca-Volmerange, B. 1997, Astron. Astrophys. 326, 950, [PEGASE] [60] Fisher, D., Franx, M. , and Illingworth, G. 1996, Astrophys. J. 459, 110 [61] Freeman, K. and Bland-Hawthorn, J. 2002, Annu. Rev. Astron. Astrophys. 40, 487 [62] Freeman, K. C. 1970, Astrophys. J. 160, 811 [63] Friedli, D. and Benz, W. 1993, Astron. Astrophys. 268, 65 [64] Friedli, D., Benz, W., and Kennicutt, R. 1994, Astrophys. J. Let. 430, L105 [65] Garnett, D. R. 2002, Astrophys. J. 581, 1019 [66] Garnett, D. R., Shields, G. A., Skillman, E. D., Sagan, S. P., and Dufour, R. J. 1997, Astrophys. J. 489, 63 [67] Gavazzi, G., Bonfanti, C., Sanvito, G., Boselli, A., and Scodeggio, M . 2002, As-trophys. J. 576, 135 [68] Gonzalez, J. J. 1993, Ph.D. Thesis [69] Gordon, K. D., Calzetti, D., and Witt, A. N. 1997, Astrophys. J. 487, 625 [70] Gordon, K. D., Misselt, K. A., Witt, A. N., and Clayton, G. C. 2001, Astrophys. J. 551, 269 [71] Gorgas, J., Cardiel, N. , Pedraz, S., and Gonzalez, J. J. 1999, Astron. Astrophys. Suppl. Ser. 139, 29 [72] Gorgas, J., Faber, S. M. , Burstein, D., Gonzalez, J. J., Courteau, S., and Prosser, C. 1993, Astrophys. J. Supp. Ser. 86, 153 [73] Goudfrooij, P., Gorgas, J., and Jablonka, P. 1999, Astrophys. and Space Science 269, 109 [74] Greggio, L. and Renzini, A. 1983, Astron. Astrophys. 118, 217 Bibliography 246 [75] Grosb0l, P. J. 1985, Astron. Astrophys. Suppl. Ser. 60, 261 [76] Guth, A. H. 1998, The Inflationary Universe, Addison Wesley Longman [77] Hao, L., Strauss, M . A., Tremonti, C. A., Schlegel, D. J., Heckman, T. M. , Kauff-mann, G., Blanton, M. R., Fan, X., Gunn, J. E., Hall, P. B., ivezic, Z., Knapp, G. R., Krolik, J. H., Lupton, R. H., Richards, G. T., Schneider, D. P., Strateva, I. V., Za-kamska, N. L., Brinkmann, J., Brunner, R. J., and Szokoly, G. P. 2005, Astron. J. 129,1783 [78] Haynes, M. P., Giovanelli, R., Chamaraux, P., da Costa, L. N., Freudling, W., Salzer, J. J., and Wegner, G. 1999, Astron. J. 117, 2039 [79] Henry, R. B. C. and Worthey, G. 1999, Pub. Astro. Soc. Pac. I l l , 919 [80] Hook, I. M. , J0rgensen, I., Allington-Smith, J. R., Davies, R. L., Metcalfe, N., Murowinski, R. G., and Crampton, D. 2004, Pub. Astro. Soc. Pac. 116, 425 [81] Jarrett, T. H., Chester, T., Cutri, R., Schneider, S. E., and Huchra, J. P. 2003, Astron. J. 125, 525 [82] Jones, J. E., Alloin, D. M. , and Jones, B. J. T. 1984, Astrophys. J. 283, 457 [83] Jones, L. A. and Worthey, G. 1995, Astrophys. J. Let. 446, L31+ [84] Kauffmann, G. 1996, Mon. Not. Roy. Astro. Soc. 281, 487 [85] Kauffmann, G., Heckman, T. M. , White, S. D. M. , Chariot, S., Tremonti, C , Brinchmann, J., Bruzual, G., Peng, E. W., Seibert, M. , Bernardi, M. , Blanton, M. , Brinkmann, J., Castander, F., Csabai, I., Fukugita, M. , Ivezic, Z., Munn, J. A., Nichol, R. C , Padmanabhan, N., Thakar, A. R., Weinberg, D. H., and York, D. 2003a, Mon. Not. Roy. Astro. Soc. 341, 33 [86] Kauffmann, G., Heckman, T. M. , White, S. D. M. , Chariot, S., Tremonti, C., Peng, E. W., Seibert, M. , Brinkmann, J., Nichol, R. C , SubbaRao, M. , and York, D. 2003b, Mon. Not. Roy. Astro. Soc. 341, 54 Bibliography 247 [87] Kormendy, J. 1982, Astrophys. J. 257, 75 [88] Kormendy, J. and Kennicutt, R. C. 2004, Annu. Rev. Astron. Astrophys. 42, 603 [89] Kraft, R. P. 1994, Pub. Astro. Soc. Pac. 106, 553 [90] Kroupa, P. 2002, Science 295, 82 [91] Kuchinski, L. E., Terndrup, D. M. , Gordon, K. D., and Witt, A. N. 1998a, Astron. J. 115, 1438 [92] Kuchinski, L. E., Terndrup, D. M. , Gordon, K. D., and Witt, A. N. 1998b, Astron. J. 115, 1438 [93] Larson, R. B. 1974, Mon. Not. Roy. Astro. Soc. 166, 585 [94] Lee, J. C., Salzer, J. J., and Melbourne, J. 2004, Astrophys. J. 616, 752 [95] Lemaitre, G. 1931, Mon. Not. Roy. Astro. Soc. 91, 490+ [96] Leonardi, A. J. and Worthey, G. 2000, Astrophys. J. 534, 650 [97] Lin, D. N. C. and Pringle, J. E. 1987, Astrophys. J. Let. 320, L87 [98] Liu, M . C., Chariot, S., and Graham, J. R. 2000, Astrophys. J. 543, 644 [99] Loveday, J., Peterson, B. A., Maddox, S. J., and Efstathiou, G. 1996, Astrophys. J. Supp. Ser. 107, 201 [100] Mollenhoff, C. and Heidt, J. 2001, Astron. Astrophys. 368, 16 [101] MacArthur, L. A. 2005, Astrophys. J. 623, 795 [102] MacArthur, L. A., Courteau, S., Bell, E., and Holtzman, J. A. 2004, Astrophys. J. Supp. Ser. 152, 175 [103] MacArthur, L. A., Courteau, S., and Holtzman, J. A. 2003, Astrophys. J. 582, 689, [Paper I] Bibliography 248 [104] Maraston, C. 1998, Mon. Not. Roy. Astro. Soc. 300, 872 [105] Maraston, C. 2003, in Extragalactic Globular Cluster Systems, pp 237—h [106] Maraston, C. and Thomas, D. 2000, Astrophys. J. 541, 126 [107] Martin, P. 1995, Astron. J. 109, 2428+ [108] Martin, P. and Roy, J. 1994, Astrophys. J. 424, 599 [109] Martinet, L. 1995, Fundamentals of Cosmic Physics 15, 341 [110] Masters, K. L., Giovanelli, R., and Haynes, M. P. 2003, Astron. J. 126, 158 [111] Matteucci, F. 2002, Chemical Evolution of Galaxies and Intracluster Medium., Proc. XII Canary Islands Winter School of Astrophysics, (astro-ph/0203340) [112] Matthews, L. D. and Wood, K. 2001, Astrophys. J. 548, 150 [113] Mehlert, D., Thomas, D., Saglia, R. P., Bender, R., and Wegner, G. 2003, Astron. Astrophys. 407, 423 [114] Mestel, L. 1963, Mon. Not. Roy. Astro. Soc. 126, 553+ [115] Meurer, G. R. 2000, in Astronomical Society of the Pacific Conference Series, pp 81-+ [116] Mo, H. J., Mao, S., and White, S. D. M. 1998, Mon. Not. Roy. Astro. Soc. 295, 319 [117] Moy, E., Rocca-Volmerange, B., and Fioc, M. 2001, Astron. Astrophys. 365, 347 [118] Nilson, P. 1973, Nova Acta Regiae Soc. Sci. Upsaliensis Ser. V pp 0+ [119] Norman, C. A., Sellwood, J. A., and Hasan, H. 1996, Astrophys. J. 462, 114+ V Bibliography 249 [120] Osterbrock, D. E. 1989, Astrophysics of gaseous nebulae and active galactic nu-clei, Research supported by the University of California, John Simon Guggenheim Memorial Foundation, University of Minnesota, et al. Mill Valley, CA, University Science Books, 1989, 422 p. [121] Panuzzo, P., Bressan, A., Granato, G. L., Silva, L., and Danese, L. 2003, Astron. Astrophys. 409, 99 [122] Peebles, P. J. E. 1969, Astrophys. J. 155, 393+ [123] Peebles, P. J. E. 1980, Cosmology: The Physics of Large Scale Structure, Prince-ton University Press, Princeton [124] Peletier, R., Balcells, M. , Falcon-Barroso, J., and Graham, A. 2004, Baryons in Dark Matter Halos [125] Peletier, R. F. and Balcells, M. 1996, Astron. J. I l l , 2238 [126] Peletier, R. F., Balcells, M. , Davies, R. L., Andredakis, Y. , Vazdekis, A., Burkert, A., and Prada, F. 1999, Mon. Not. Roy. Astro. Soc. 310, 703 [127] Peletier, R. F., Valentijn, E. A., Moorwood, A. F. M. , Freudling, W., Knapen, J. Ff., and Beckman, J. E. 1995, Astron. Astrophys. 300, L1+ [128] Pfenniger, D. 1993, in IAU Symp. 153: Galactic Bulges, Vol. 153, pp 387+ [129] Prantzos, N. and Boissier, S. 2000, Mon. Not. Roy. Astro. Soc. 313, 338, [PB00] [130] Primack, J. R. 2000, in ASP Conf. Ser. 201: Cosmic Flows Workshop, pp 389+ [131] Primack, J. R. 2004a, in Proceedings of 6th UCLA Symposium on Sources and Detection of Dark Matter in the Universe, Marina del Rey, February 2004 [132] Primack, J. R. 2004b, in IAU Symposium, pp 53-+ [133] Proctor, R. N. and Sansom, A. E. 2002, Mon. Not. Roy. Astro. Soc. 333, 517 Bibliography 250 [134] Renzini, A. 1983, Memorie delta Societa Astronomica Italiana 54, 335 [135] Rose, J. A. 1984, Astron. J. 89, 1238 [136] Rose, J. A. 1985, Astron. J. 90, 1927 [137] Rose, J. A. 1994, Astron. J. 107, 206 [138] Ryder, S. D., Fenner, Y., and Gibson, B. K. 2005, Mon. Not. Roy. Astro. Soc. 358, 1337 [139] Saio, H. and Yoshii, Y. 1990, Astrophys. J. 363, 40 [140] Salpeter, E. E. 1955, Astrophys. J. 121, 161 [141] Sandage, A. 1986, Astron. Astrophys. 161, 89 [142] Savage, B. D. and Mathis, J. S. 1979, Annu. Rev. Astron. Astrophys. 17, 73 [143] Schiavon, R. P., Faber, S. M. , Castilho, B. V., and Rose, J. A. 2002a, Astrophys. J. 580, 850 [144] Schiavon, R. P., Faber, S. M. , Rose, J. A., and Castilho, B. V. 2002b, Astrophys. J. 580, 873 [145] Schlegel, D. J., Finkbeiner, D. P., and Davis, M. 1998, Astrophys. J. 500, 525 [146] Searle, L., Sargent, W. L. W., and Bagnuolo, W. G. 1973, Astrophys. J. 179, 427 [147] Seigar, M. , Carollo, C. M. , Stiavelli, M. , de Zeeuw, P. T., and Dejonghe, H. 2002, Astron. J. 123, 184 [148] Sellwood, J. A. and Binney, J. J. 2002, Mon. Not. Roy. Astro. Soc. 336, 785 [149] Sellwood, J. A. and Moore, E. M. 1999, Astrophys. J. 510, 125 [150] Sellwood, J. A. and Wilkinson, A. 1993, Rep. Prog. Phys. 56, 173 Bibliography 251 [151] Sersic, J. L. 1968, Atlas de Galaxias Australes, Observatorio Astronomico, Cor-doba, Argentina [152] Silva, L., Granato, G. L., Bressan, A., and Danese, L. 1998, Astrophys. J. 509, 103, [GRASIL] [153] Smoot, G. F., Bennett, C. L., Kogut, A., Wright, E. L., Aymon, J., Boggess, N. W., Cheng, E. S., de Amici, G., Gulkis, S., Hauser, M. G., Hinshaw, G., Jackson, P. D., Janssen, M. , Kaita, E., Kelsall, T., Keegstra, P., Lineweaver, C , Loewenstein, K., Lubin, P., Mather, J., Meyer, S. S., Moseley, S. H., Murdock, T., Rokke, L., Silverberg, R. F., Tenorio, L., Weiss, R., and Wilkinson, D. T. 1992, Astrophys. J. Let. 396, L l [154] Somerville, R. S. and Primack, J. R. 1999, Mon. Not. Roy. Astro. Soc. 310, 1087 [155] Spergel, D. N., Verde, L., Peiris, H. V., Komatsu, E., Nolta, M . R., Bennett, C. L., Halpern, M. , Hinshaw, G., Jarosik, N., Kogut, A., Limon, M. , Meyer, S. S., Page, L., Tucker, G. S., Weiland, J. L., Wollack, E., and Wright, E. L. 2003, Astrophys. J. Supp. Ser. 148, 175 [156] Storchi-Bergmann, T., Calzetti, D., and Kinney, A. L. 1994, Astrophys. J. 429, 572 [157] Struck-Marcell, C. 1991, Astrophys. J. 368, 348 [158] Tantalo, R. and Chiosi, C. 2004, Mon. Not. Roy. Astro. Soc. 353, 917 [159] Theureau, G., Bottinelli, L., Coudreau-Durand, N., Gouguenheim, L., Hallet, N., Loulergue, M. , Paturel, G., and Teerikorpi, P. 1998, Astron. Astrophys. Suppl. Ser. 130,333 [160] Thomas, D. and Kauffmann, G. 1999, in ASP Conf. Ser. 192: Spectrophotometry Dating of Stars and Galaxies, pp 261—h Bibliography 252 [161] Thomas, D., Maraston, C , and Bender, R. 2003, Mon. Not. Roy. Astro. Soc. 3 3 9 , 897, [TMB03] [162] Tinsley, B. M . 1972, Astrophys. J. 1 7 8 , 319 [163] Tinsley, B. M . and Gunn, J . E . 1976, Astrophys. J. 2 0 3 , 52 [164] Tody, D. 1993, in ASP Conf. Ser. 52: Astronomical Data Analysis Software and Systems II, pp 173—h [165] Trager, S. C., Dalcanton, J . J. , and Weiner, B. J. 1999, in The Formation of Galactic Bulges, pp 42—h [166] Trager, S. C., Faber, S. M . , Worthey, G., and Gonzalez, J. J. 2000, Astron. J. 119,1645 [167] van den Bosch, F. C. 2000, Astrophys. J. 5 3 0 , 177 [168] van Zee, L. , Salzer, J. J. , Haynes, M . P., O'Donoghue, A . A. , and Balonek, T. J . 1998, Astron. J. 1 1 6 , 2805 [169] Vazdekis, A . 1999, Astrophys. J. 5 1 3 , 224 [170] Vazdekis, A . and Arimoto, N . 1999, Astrophys. J. 5 2 5 , 144 [171] Vazdekis, A. , Cenarro, A . J., Gorgas, J., Cardiel, N . , and Peletier, R. F. 2003, Mon. Not. Roy. Astro. Soc. 3 4 0 , 1317 [172] Vazdekis, A. , Salaris, M . , Arimoto, N . , and Rose, J . A . 2001, Astrophys. J. 5 4 9 , 274 [173] Walker, I. R., Mihos, J . C , and Hernquist, L. 1996, Astrophys. J. 4 6 0 , 121+ [174] Witt, A . N . , Thronson, H. A. , and Capuano, J. M . 1992, Astrophys. J. 3 9 3 , 611 [175] Worthey, G. 1993, Astrophys. J. 4 0 9 , 530 [176] Worthey, G. 1994, Astrophys. J. Supp. Ser. 9 5 , 107 Bibliography 253 [177] Worthey, G., Faber, S. M. , Gonzalez, J. J., and Burstein, D. 1994, Astrophys. J. Supp. Ser. 94, 687 [178] Worthey, G. and Ottaviani, D. L. 1997, Astrophys. J. Supp. Ser. I l l , 377, [W097] [179] Wyse, R. F. G. 1999, in The Formation of Galactic Bulges, pp 195-+ [180] Xilouris, E. M. , Byun, Y. I., Kylafis, N. D., Paleologou, E. V., and Papamas-torakis, J. 1999, Astron. Astrophys. 344, 868 [181] Zaritsky, D., Rieke, M. , and Rix, H.-W. 1993, in IAUSymp. 153: Galactic Bulges, Vol. 153, pp 441+ [182] Zhang, B. and Wyse, R. F. G. 2000, Mon. Not. Roy. Astro. Soc. 313, 310 

Cite

Citation Scheme:

        

Citations by CSL (citeproc-js)

Usage Statistics

Share

Embed

Customize your widget with the following options, then copy and paste the code below into the HTML of your page to embed this item in your website.
                        
                            <div id="ubcOpenCollectionsWidgetDisplay">
                            <script id="ubcOpenCollectionsWidget"
                            src="{[{embed.src}]}"
                            data-item="{[{embed.item}]}"
                            data-collection="{[{embed.collection}]}"
                            data-metadata="{[{embed.showMetadata}]}"
                            data-width="{[{embed.width}]}"
                            async >
                            </script>
                            </div>
                        
                    
IIIF logo Our image viewer uses the IIIF 2.0 standard. To load this item in other compatible viewers, use this url:
http://iiif.library.ubc.ca/presentation/dsp.831.1-0085579/manifest

Comment

Related Items