The CALIFA Survey Across the Hubble Sequence Spatially Resolved Stellar Population Properties in Galaxies?

A&A 581, A103 (2015) Astronomy DOI: 10.1051/0004-6361/201525938 & c ESO 2015 Astrophysics

The CALIFA survey across the Hubble sequence Spatially resolved stellar population properties in galaxies?

R. M. González Delgado1, R. García-Benito1, E. Pérez1, R. Cid Fernandes2, A. L. de Amorim2, C. Cortijo-Ferrero1, E. A. D. Lacerda1,2, R. López Fernández1, N. Vale-Asari2, S. F. Sánchez3, M. Mollá4, T. Ruiz-Lara5, P. Sánchez-Blázquez6, C. J. Walcher7, J. Alves8, J. A. L. Aguerri9, S. Bekeraité7, J. Bland-Hawthorn10, L. Galbany11,12, A. Gallazzi13, B. Husemann14, J. Iglesias-Páramo1,15, V. Kalinova16, A. R. López-Sánchez17, R. A. Marino18, I. Márquez1, J. Masegosa1, D. Mast19, J. Méndez-Abreu20, A. Mendoza1, A. del Olmo1, I. Pérez5, A. Quirrenbach21, S. Zibetti12, and CALIFA collaboration?? (Aﬃliations can be found after the references)

Received 21 February 2015 / Accepted 11 June 2015

ABSTRACT

Various different physical processes contribute to the star formation and stellar mass assembly histories of galaxies. One important approach to understanding the significance of these different processes on galaxy evolution is the study of the stellar population content of today’s galaxies in a spatially resolved manner. The aim of this paper is to characterize in detail the radial structure of stellar population properties of galaxies in the nearby universe, based on a uniquely large galaxy sample, considering the quality and coverage of the data. The sample under study was drawn from the CALIFA survey and contains 300 galaxies observed with integral field spectroscopy. These cover a wide range of Hubble types, 9 11 from spheroids to spiral galaxies, while stellar masses range from M? ∼ 10 to 7 × 10 M . We apply the fossil record method based on spectral synthesis techniques to recover the following physical properties for each spatial resolution element in our target galaxies: the stellar mass surface density (µ?), stellar extinction (AV ), light-weighted and mass-weighted ages (hlog ageiL, hlog ageiM), and mass-weighted metallicity (hlog Z?iM). To study mean trends with overall galaxy properties, the individual radial profiles are stacked in seven bins of galaxy morphology (E, S0, Sa, Sb, Sbc, Sc, and Sd). We confirm that more massive galaxies are more compact, older, more metal rich, and less reddened by dust. Additionally, we find that these trends are preserved spatially with the radial distance to the nucleus. Deviations from these relations appear correlated with Hubble type: earlier types are more compact, older, and more metal rich for a given M?, which is evidence that quenching is related to morphology, but not driven by mass. Negative gradients of hlog ageiL are consistent with an inside-out growth of galaxies, with the largest hlog ageiL gradients in Sb–Sbc galaxies. Further, the mean stellar ages of disks and bulges are correlated and with disks covering a wider range of ages, and late-type spirals hosting younger disks. However, age gradients are only mildly negative or flat beyond R ∼ 2 HLR (half light radius), indicating that star formation is more uniformly distributed or that stellar migration is important at these distances. The gradients in stellar mass surface density depend mostly on stellar mass, in the sense that more massive galaxies are more centrally concentrated. Whatever sets the concentration indices of galaxies obviously depends less on quenching/morphology than on the depth of the potential well. There is a secondary correlation in the sense that at the same M? early-type galaxies have steeper gradients. The µ? gradients outside 1 HLR show no dependence on Hubble type. We find mildly negative hlog Z?iM gradients, which are shallower than predicted from models of galaxy evolution in isolation. In general, metallicity gradients depend on stellar mass, and less on morphology, hinting that metallicity is affected by both – the depth of the potential well and morphology/quenching. Thus, the largest hlog Z?iM gradients occur in Milky Way-like Sb–Sbc galaxies, and are similar to those measured above the Galactic disk. Sc spirals show flatter hlog Z?iM gradients, possibly indicating a larger contribution from secular evolution in disks. The galaxies from the sample have decreasing-outward stellar extinction; all spirals show similar radial profiles, independent from the stellar mass, but redder than E and S0. Overall, we conclude that quenching processes act in manners that are independent of mass, while metallicity and galaxy structure are influenced by mass-dependent processes. Key words. techniques: spectroscopic – Galaxy: evolution – Galaxy: stellar content – galaxies: structure – Galaxy: fundamental parameters – galaxies: spiral

1. Introduction Theoretically, the formation of large-scale structures arise through the evolution of cold dark matter. In this picture, Galaxies are a complex mix of stars, gas, dust, and dark mat- small-scale density perturbations in the dark matter collapse ter distributed in different components (bulge, disk, and halo) and form the first generation of dark matter halos, which sub- whose present day structure and dynamics are intimately linked sequently merge to form larger structures such as clusters and to their assembly and evolution over the history of the Universe. superclusters (Springel et al. 2005; De Lucia et al. 2006). This Different observational and theoretical approaches can be fol- basic hierarchical picture explains the global evolution of the lowed to learn how galaxies form and evolve. star formation rate density of the universe, with galaxy peak formation epoch at redshift 2−3 (e.g. Madau & Dickinson 2014, and references therein). The stellar components formed at earlier ? Appendices are available in electronic form at epochs likely evolved into elliptical galaxies and bulges through http://www.aanda.org mergers of the primordial star-forming disks (Elmegreen et al. ?? http://califa.caha.es 2007; Bournaud et al. 2007). However, this framework fails to

Article published by EDP Sciences A103, page 1 of 44 A&A 581, A103 (2015) explain how the galaxy population emerges at z ∼ 1, and how extinction and kinematics, as well as a suit of nebular proper- the present day Hubble sequence of galaxies was assembled. ties, such as gas kinematics, metallicity, excitation, etc. Until a The growth of galaxies is not related in a simple way to few years ago, IFS was used to target small samples of galax- the build up of dark matter; the interplay of energy and mat- ies. Detailed programs, such as SAURON (Bacon et al. 2001), ter exchange (between the process of gas accretion and cool- VENGA (Blanc et al. 2013), (U)LIRs at z ≤ 0.26 (Arribas et al. ing and star formation) is essential for the growth the gaseous 2010), PINGS (Rosales-Ortega et al. 2010), or DiskMass Survey and stellar components in galaxies. Feedback mechanisms re- (Bershady et al. 2010), have been limited to less than a hundred sulting from stellar winds, supernova explosions, and AGN are galaxies, but it is more than fair to recognize that to get these relevant to stop the gas collapse and cooling, quenching the star amounts of IFU data was a challenge at the time. ATLAS3D formation and hence galaxy growth (Silk & Rees 1998; Hopkins (Cappellari et al. 2011) represented a step forward, with the ob- et al. 2011). Although these processes are difficult to implement servation of a volume-limited sample of 260 galaxies, but with in theoretical models, they are essential to explain the masses, three important limitations: the sample only includes early-type structures, morphologies, stellar populations, and chemical com- galaxies, the field of view is limited to 1 effective radius, and the positions of galaxies, and the evolution of these properties with spectral range is restricted from Hβ to [NI]λ5200. cosmic time. CALIFA (Calar Alto Legacy Integral Field Area) is our Recently, a new set of cosmological hydrodynamic simula- ongoing survey of 600 nearby galaxies at the 3.5 m at Calar tions have started to predict how the spatially resolved informa- Alto (Sánchez et al. 2012)1. The data set provided by the tion of the properties of stellar populations in galaxies can con- survey (see Husemann et al. 2013, for DR1; García-Benito strain the complex interplay between gas infall, outflows, stellar et al. 2015 for DR2) is unique to advance in these issues migration, radial gas flows, and star formation efficiency, in driv- not only because of its ability to provide spectral and spa- ing the inside-out growth of galactic disks (Brook et al. 2012; tial information, but also because of the following: a) It in- Gibson et al. 2013; Few et al. 2012; Pilkington et al. 2012a; cludes a large homogeneous sample of galaxies across the color- Minchev et al. 2014). Radiative cooling, star formation, feed- magnitude diagram, covering a large range of masses (109 to 12 back from supernovae, and chemical enrichment are also in- 10 M , González Delgado et al. 2014c), and morphologies cluded in simulations to predict radial metallicity gradients as from Ellipticals (E0-E7), Lenticulars (S0-S0a), to spirals (Sa to a function of merging history. Shallow metallicity gradients are Sm; see Walcher et al. 2014 for a general description of the sam- expected if elliptical galaxies result from major mergers (e.g. ple). b) It has a large field of view (7400 ×6500) with a final spatial Kobayashi 2004), but a minor merger picture for the formation sampling of 1 arcsec, and a resolution of ∼2.5 arcsec, allowing of ellipticals can successfully explain the strong size evolution of it to spatially resolve the stellar population properties well, and massive galaxies (Oser et al. 2012). This late-time accretion of to obtain the total integrated properties, such as galaxy stellar low-mass and metal poor galaxies (dry mergers) into the already mass, and stellar metallicity. c) It covers the whole rest-frame formed massive galaxy can produce a variation of the age and optical wavelength at intermediate spectral resolution, including metallicity radial structure of the galaxy as it increases in size. the most relevant absorption diagnostics for deriving the stellar Galactic stellar winds and metal cooling also have an important population properties. effect on these ex situ star formation models, predicting different Previous papers in this series used the first ∼100 datacubes of behavior of the mass and metallicity assembly in massive early- the survey to derive spatially resolved stellar population proper- type galaxies, and in the radial gradient of present stellar popu- ties by means of full spectral fitting techniques. We obtained the lations properties of galaxies (Hirschmann et al. 2013, 2015). following: 1) massive galaxies grow their stellar mass inside-out. In summary, these theoretical works show that observational The signal of downsizing is shown to be spatially preserved, with data with spatial information of the mass and metallicity assem- both inner and outer regions growing faster for more massive bly and their cosmic evolution, and the present radial structure galaxies. The relative growth rate of the spheroidal component of stellar population properties (stellar mass surface density, age, (nucleus and inner galaxy), which peaked 5–7 Gyr ago, shows 10 metallicity) contain relevant information to constrain the forma- a maximum at a critical stellar mass M? ∼ 7 × 10 M (Pérez tion history of galaxies and the physics of feedback mechanisms et al. 2013). 2) The inside-out scenario is also supported by the involved. negative radial gradients of the stellar population ages (González Observationally, a first step is to study what kinds of galax- Delgado et al. 2014c). 3) Global and local relations between stel- ies are there in the Universe, and which are their physical prop- lar mass, stellar mass surface density, and stellar metallicity re- erties. Attending to their form and structure, galaxies can be lation were investigated, along with their evolution (as derived grouped into a few categories. Results show that most of the from our fossil record analysis). In disks, the stellar mass sur- massive nearby galaxies are ellipticals, S0’s, or spirals (Blanton face density regulates the ages and the metallicity. In spheroids, & Moustakas 2009) following the Hubble tuning-fork diagram. the galaxy stellar mass dominates the physics of star formation In this scheme, S0’s are a transition between spirals and ellipti- and chemical enrichment (González Delgado et al. 2014c,a). 4) cals (Cappellari et al. 2013), and the bulge/disk ratio increases In terms of integrated versus spatially resolved properties, the from late to early-type spirals. At the same time, galaxy prop- stellar population properties are well represented by their values erties such as color, mass, surface brightness, luminosity, and at 1 HLR (González Delgado et al. 2014c,a). The CALIFA col- gas fraction are correlated with Hubble type (Roberts & Haynes laboration has also compared the age and metallicity gradients in 1994), suggesting that the Hubble sequence somehow reflects a subsample of 62 face-on spirals and it was found that there is possible paths for galaxy formation and evolution. However, the no difference between the stellar population properties in barred processes structuring galaxies along the Hubble sequence are and unbarred galaxies (Sánchez-Blázquez et al. 2014). still poorly understood. In this paper, we extend our study of the spatially resolved Integral field spectroscopy (IFS) enables a leap forward, pro- star formation history (SFH) of CALIFA galaxies to derive the viding 3D information (2D spatial + 1D spectral) on galax- radial structure of the stellar population properties as a function ies. These datacubes allow one to recover two-dimensional maps of stellar mass surface density, stellar ages, metallicities, 1 http://califa.caha.es

A103, page 2 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence

Fig. 1. Left: comparison of the distribution of Hubble types in the CALIFA mother sample (empty bars) and the galaxies analyzed here (ﬁlled bars). The number of galaxies in our sample are labeled in colors. The histograms are normalized to form a probability density, i.e., each bar scales with the ratio of the number of galaxies in each bin and the total number of galaxies, such that the two distributions are directly comparable. Right: color-magnitude diagram. Mother sample galaxies are plotted in grey, while the 300 galaxies analyzed in this work are marked as colored points.

of Hubble type, and galaxy stellar mass, M?. The goals are: 1) to The sample for this paper comprises the 312 CALIFA galax- characterize in detail the radial structure of stellar population ies observed in both V1200 and V500 setups as of January 2014. properties of galaxies in the local universe; 2) to find out how The 12 galaxies showing “merger or interacting features” are these properties are correlated with Hubble type, and if the not discussed here, leaving a main sample of 300 objects with a Hubble sequence is a scheme to organize galaxies by mass and well defined morphology. For this work we have grouped galax- age, and/or mass and metallicity; 3) to establish observational ies into seven morphology bins: E, S0 (including S0 and S0a), constraints to galaxy formation models via the radial distribu- Sa (Sa and Sab), Sb, Sbc, Sc (Sc and Scd), and Sd (13 Sd, 1 Sm tions and gradients of stellar populations for disk and bulge dom- and 1 Irr). inated galaxies. Figure1 shows that these 300 galaxies provide a fair rep- This paper is organized as follows: Sect. 2 describes the resentation of the CALIFA survey as a whole. The left panel observations and summarizes the properties of the CALIFA shows scaled histograms of the Hubble type in the mother sam- galaxies analyzed here. In Sect. 3 we summarize our method ple (empty bars) and in our sample (filled bars). The number for extracting the SFH, based on the fossil record method, and of objects in each morphology bin for our sample is indicated at we explain the main differences between the analysis presented the top, with a brown to blue color palette that represents Hubble here and that in previous works. Section 4 presents results on the types from ellipticals to late spirals. This same color scheme is galaxy stellar mass, half-light, and half-mass radii (HLR, HMR, used throughout this paper. The similarity of the distributions respectively), and galaxy averaged stellar metallicity. Section 5 reflects the random sampling strategy of CALIFA, with targets deals with the spatially resolved properties of the stellar popula- picked from the mother sample on the basis of visibility crite- tion: stellar mass surface density, µ?; luminosity weighted mean ria alone. The right panel in Fig.1 shows the u − r versus Mr age, hlog ageiL; mass weighted mean metallicity, hlog Z?iM; and CMD, with grey points representing the mother sample and col- stellar extinction, AV . We discuss the results in Sect. 6; and ored points the 300 galaxies. As for the Hubble type distribution, Sect. 7 presents the conclusions. a simple visual inspection shows that our subsample is representative of the full CALIFA sample in terms of CMD coverage. 2. Sample, observations, and data reduction 2.2. Observations and data reduction 2.1. Sample and morphological classification The observations were carried out with the Potsdam Multi- The CALIFA mother sample consists of 939 galaxies selected Aperture Spectrometer (Roth et al. 2005, PMAS) in the PPaK from SDSS survey in the redshift range z = 0.005–0.03, and mode (Verheijen et al. 2004) at the 3.5 m telescope of Calar Alto with r-band angular isophotal diameter of 45−8000. These crite- observatory. PPaK contains 382 fibers of 2.700 diameter each, ria guarantee that the objects fill the 7400 × 6400 FoV well. The and a 7400 × 6400 FoV (Kelz et al. 2006). Each galaxy is ob- sample includes a significant number of galaxies in different bins served with two spectral settings, V500 and V1200, with spectral in the color-magnitude diagram (CMD), ensuring that CALIFA spans a wide and representative range of galaxy types. resolutions ∼6 Å (FWHM) and 2.3 Å, respectively. The V500 The galaxies were morphologically classified as grating covers from 3745 to 7300 Å, while the V1200 covers Ellipticals (E0–7), Spirals (S0, S0a, Sab, Sb, Sbc, Sc, Scd, 3650−4840 Å. Detailed descriptions of the observational strat- Sd, Sm), and Irregulars (I). The classification was carried egy and of the data can be found in Sánchez et al.(2012), and out through visual inspection of the r-band images averaging Husemann et al.(2013). the results (after clipping outliers) from five members of the The datacubes analyzed here have been calibrated with ver- collaboration. Galaxies are also classified as B for barred, sion 1.5 of the reduction pipeline. The main issues addressed otherwise A, or AB if it is unsure, and as M if it shows “merger” by this new version are: (i) correction of the sensitivity curve or “interaction features” (Walcher et al. 2014). for the V500 grating; (ii) new registering method to determine,

A103, page 3 of 44 A&A 581, A103 (2015) for each galaxy, the relative positioning of the three pointings of SSP models, labeled as GM, CB, and BC in Cid Fernandes et al. the dithering pattern, and absolute WCS registration; (iii) a new (2014). The first two are used again here, but extended to a wider cube interpolation method. CALIFA pipeline v1.5 improves the range of metallicities, producing what we denote as bases GMe flux calibration to an accuracy of 2−3% and is the current official and CBe. data release. A detailed account of this new pipeline is presented Base GMe is a combination of the SSP spectra provided by in the Data Release 2 article (García-Benito et al. 2015). Vazdekis et al.(2010) for populations older than t = 63 Myr and To reduce the effects of vignetting on the data, we combine the González Delgado et al.(2005) models for younger ages. The the observations in the V1200 and V500 setups. The combined evolutionary tracks are those of Girardi et al.(2000), except for datacubes were processed as described in Cid Fernandes et al. the youngest ages (1 and 3 Myr), which are based on the Geneva (2013). Our analysis requires that spectra have signal-to-noise tracks (Schaller et al. 1992; Schaerer et al. 1993; Charbonnel ratio S/N ≥ 20 in a 90 Å window centered at 5635 Å (rest- et al. 1993). The Initial Mass Function (IMF) is Salpeter. In our frame). When individual spaxels do not meet this S/N thresh- previous studies of the first 100 CALIFA galaxies, we defined old, they are coadded into Voronoi zones (Cappellari & Copin base GM as a regular (t, Z) grid of these models, with 39 ages 2003). Further pre processing steps include spatial masking of spanning t = 0.001−14 Gyr and four metallicities from 0.2 to foreground and background sources, rest-framing, and spec- 1.5 Z . We now extend the Z range to use of all seven metallicites tral resampling. The resulting 253 418 spectra were then pro- provided by Vazdekis et al.(2010) models: log Z/Z = −2.3, cessed through  and y (the Python CALIFA −1.7, −1.3, −0.7, −0.4, 0, and +0.22. Because these models lack  Synthesis Organizer), producing the stellar popula- ages below 63 Myr, these young ages are only covered by the tion properties discussed here as described in detail in the next four largest metallicities, such that our extended GM base is no section. longer regular in t and Z. Base GMe contains N? = 235 elements. Base CBe is built from an update of the Bruzual & Charlot 3. Stellar population analysis: differences (2003) models (Charlot & Bruzual 2007, private communi- with respect to previous work cation), replacing STELIB (Le Borgne et al. 2003) with a combination of the MILES (Sánchez-Blázquez et al. 2006; Our method to extract stellar population properties from dat- Falcón-Barroso et al. 2011) and  (Martins et al. 2005) acubes has been explained and applied to CALIFA in Pérez spectral libraries (the same ones used in base GMe). The evo- et al.(2013), Cid Fernandes et al.(2013, 2014), and González lutionary tracks are those collectively known as Padova 1994 Delgado et al.(2014c,a). In short, we analyze the data with the (Alongi et al. 1993; Bressan et al. 1993; Fagotto et al. 1994a,b;  code (Cid Fernandes et al. 2005), which fits an ob- Girardi et al. 1996). The IMF is that of Chabrier(2003). Whereas served spectrum (O ) in terms of a model (M ) built by a non- λ λ in previous works we limited the Z range to ≥ 0.2 solar, we parametric linear combination of N simple stellar populations ? now extend this base to six metalicities: log Z/Z = −2.3, −1.7, (SSP) from a base spanning different ages (t) and metallici- −0.7, −0.4, 0, and +0.4. Base CBe contains N = 246 elements ties (Z). Dust effects are modeled as a foreground screen with ? (41 ages from 0.001 to 14 Gyr and the six metallicities above). a Cardelli et al.(1989) reddening law with R = 3.1. Windows V The main similarities and differences between bases GMe around the main optical emission lines and the NaI D absorption and CBe are the same as between the original GM and CB bases, doublet (because of its interstellar component) are masked in all which are thoroughly discussed in Cid Fernandes et al.(2014). spectral fits2. Bad pixels (identified by the reduction pipeline) Throughout the main body of this paper we focus on results ob- are also masked. Results for each spectrum are then packed and tained with base GMe, but we anticipate that our overall quali- organized with the y pipeline. tative findings remain valid for base CBe. The role of base CBe This working scheme is preserved here, but with three new is to enable a rough assessment of the uncertainties associated developments: with the model choice. 1. The datacubes come from the version 1.5 of the reduction A minor technical difference with respect to our previous pipeline (García-Benito et al. 2015). analysis is that we now smooth the spectral bases to 6 Å FWHM 2. Larger and more complete SSP bases are employed. effective resolution prior to the fits. This is because the kine- 3. A somewhat different definition of mean stellar metallicity is matical filter implemented in  operates in velocity- adopted (see González Delgado et al. 2014a). space, whereas both CALIFA and the SSP model spectra have a constant spectral resolution in λ-space, so that effects of the This section describes the novelties related to the stellar popula- instrumental broadening can only be mimicked approximately tion synthesis. Improvements resulting from the new reduction by . We have verified that this modification does not pipeline are described in Appendix A. affect the stellar population properties used in this paper. AppendixB presents some comparisons of the results ob- 3.1. SSP spectral bases tained with these two bases. Experiments were also performed with bases which extend the age range to 18 Gyr and configure SSP models are a central ingredient in our analysis, linking the  to allow negative values of AV . These tests are also results of the spectral decomposition to physical properties of discussed in AppendixB, which adds to the collection of “sanity the stellar populations. Our previous applications of  checks” on the results of our analysis. to CALIFA explored spectral bases built from three sets of

2 To test the effect of this process in age estimation, we have compared 4. Galaxy mass, metric, and stellar metallicity with the Sánchez-Blázquez et al.(2014) results for the 60 galaxies in common. This work uses Steckmap (Ocvirk et al. 2006) and the Hβ line This section addresses three relatively unrelated aspects, which (previously corrected for the emission contribution). Statistically, we are all important to better understand the results presented in find that there is no difference in age (mean = –0.04, std = 0.15 dex) if the next section, where we examine how the spatial distribu- the same SSP models are used in the two methods. tion of stellar population properties relates to a galaxy’s stellar

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CBe-based masses (Chabrier IMF) are on average smaller by a factor 1.84. As for the general galaxy population, mass is well correlated with Hubble type, decreasing from early to late types. High bulge-to-disk ratios (E, S0, Sa) are the most mas- 11 sive ones (≥10 M ), while galaxies with small bulges (Sc–Sd) 10 have M? ≤ 10 M . The average log M?(M ) is 11.4, 11.1, 11.0, 10.9, 10.7, 10.1, and 9.5 for E, S0, Sa, Sb, Sbc, Sc, and Sd, respectively. The dispersion is typically 0.3 dex, except for Sc galaxies, which have a dispersion of ∼0.5 dex.

Because CALIFA is not complete for Mr ≥ −19.5, this distribution in mass is not completely representative of the local Universe. In particular, it is important to remember that dwarf ellipticals are not included, so M? or any other property discussed here for E’s are restricted to massive ellipticals. Fig. 2. Distribution of the stellar masses obtained from the spatially resolved spectral fits of each galaxy for each Hubble type (grey small 4.2. The HMR/HLR points). The colored dots (stars) are the mean galaxy stellar mass in each Hubble type obtained with the GMe (CBe) SSP models. The bars As explained in Cid Fernandes et al.(2013), we define the HLR show the dispersion in mass. as the semi-major axis length of the elliptical aperture which contains half of the total light of the galaxy at the rest-frame Table 1. Number of galaxies for each Hubble type and M? interval wavelength 5635 Å. Similarly, the HMR is derived from the (GMe). 2D distribution of the stellar mass, as the elliptical aperture at which the mass curve of growth reaches 50% of its asymptote. log M? (M ) bin E S0 Sa Sb Sbc Sc Sd M L The ratio between the HMR and the HLR (a50/a50) reflects the ≤9.1 – – – – – 2 2 spatial variation of the SFH in a galaxy. This ratio is lower than 9.1–9.6 – – – – – 9 8 1 in almost all cases (González Delgado et al. 2014c), a signpost 9.6–10.1 – – – – 2 10 5 of the inside-out growth found by Pérez et al.(2013). 10.1–10.6 – – 7 11 16 21 – M L 10.6–10.9 3 8 9 14 21 4 – Figure3 shows the relation between a50/a50 and Hubble type 10.9–11.2 8 14 22 17 16 3 – (left panel), and galaxy stellar mass (right panel). These plots 11.2–11.5 17 8 13 10 3 1 – confirm our earlier finding that galaxies are more compact in 11.5–11.8 12 2 – 1 – – mass than in light. If the gradient in stellar extinction is taken ≥11.8 1 – – – – – M Lintrin ± into account, the average a50/a50 = 0.82 (0.80) 0.10 (0.13) for total 40 32 51 53 58 50 15 base GMe (CBe). Figure3 shows that the ratio decreases from late to early-type spirals; while lenticulars and ellipticals have M L similar a50/a50. mass and morphology. First, Sect. 4.1 reviews the relation be- These results are also in agreement with our previous result, tween stellar mass and morphological type for our sample. This M L which a50/a50 shows a dual dependence with galaxy stellar mass: strong relation is imprinted on virtually all results discussed it decreases with increasing mass for disk galaxies, but is almost in Sect.5. Second, Sect. 4.2 compares our measurements of constant in spheroidal galaxies, as confirmed in the right panel the HLR and HMR. As discussed by González Delgado et al. of Fig.3. Sb–Sbc galaxies are those with the lowest aM /aL . (2014c), these two natural metrics for distances are not identical 50 50 because of the inside-out growth of galaxies. Here we inspect how the HMR/HLR ratio varies as a function of Hubble type 4.3. Stellar metallicity and stellar mass in our sample. Finally, Sect. 4.3 presents our definition of mean stellar metallicity. González Delgado et al. Metallicity is one of the most difficult stellar population proper- (2014c) showed that stellar mass surface densities, mean ages, ties to estimate. Reasons for this include: (i) the coarse metal- and extinction values defined from the integrated spectrum, from licity grid of the SSP bases; (ii) the limitation of the stellar galaxy-wide spatial averages, and measured at R = 1 HLR all libraries to the solar neighborhood; and (iii) inherent degen- agree very well with each other. Here we extend this test to stel- eracies like the dependence of the continuum shape on ex- lar metallicities. Throughout this section, results for the two SSP tinction, age, and metallicity, whose effects are hard to disen- models discussed in Sect. 3.1 are presented. tangle. Notwithstanding these difficulties, meaningful estimates of Z? can be extracted from observed spectra, particularly by means of full spectral synthesis methods (Sánchez-Blázquez 4.1. Stellar masses et al. 2011).

To obtain the total stellar mass of a galaxy, we add the mass in -based estimates of Z? for the same CALIFA each zone, thus taking spatial variations of the stellar extinction sample used in this paper were previously used by González and M/L ratio into account. Masked spaxels (e.g., foreground Delgado et al.(2014a) to study global and local relations of Z? stars) are accounted for using the µ? radial profile as explained with the stellar mass and stellar mass surface density. We have in González Delgado et al.(2014c). shown there that: (i) our sample follows a well defined stellar Figure2 shows the distribution of M? as a function of mass-metallicity relation (MZR); (ii) this relation is steeper that Hubble type. Table1 shows the distribution of galaxies by one obtained from O/H measurements in HII regions, but that Hubble type in several bins of M?. The masses range from considering only young stellar populations the two MZR’s are 8 11 7 × 10 to 7 × 10 M for fits with GMe (Salpeter IMF). similar; and (iii) Z? is strongly related to µ? in galaxy disks

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M L M L Fig. 3. Left: the ratio between half mass and half light radius (a50/ a50) with the Hubble type (left). Big colored dots represent the averaged a50/ a50 in each Hubble type bin, and the lines the dispersion. Stars and big circles show the results obtained with the GMe and CBe bases, respectively. M L Right: a50/ a50 as a function of the galaxy stellar mass. The black crosses show the averaged correlation independent of the morphological type. Large circles represent the averaged relation in mass intervals of 0.25 dex for each color-coded morphological type.

and to M? in spheroids. All these results lend confidence to our use of the logarithm). This is the same reasoning behind the use Z? estimates. of hlog agei instead of loghagei. Here we review our definition of the mean stellar metallicity, To some degree, the definition of mean Z is largely a matter and test whether its value at 1 HLR matches the galaxy wide of taste (albeit one with mathematical consequences because of average value as well as at obtained from the spatially collapsed the inequality of the arithmetic and geometric means, hlog Zi ≤ data cube. loghZi), so much so that one finds both types of averaging in the literature. For instance, in Gallazzi et al.(2005) metallicities are averaged logarithmically, whereas Asari et al.(2007) work with 4.3.1. Mean stellar metallicity arithmetic averages. The main properties analyzed in this paper are the stellar As shown in González Delgado et al.(2014a; see also Fig.4), mass surface density (µ∗), stellar extinction (AV ), mean age our metallicities span about 1 dex for galaxy masses ranging 9 12 (hlog ageiL), and metallicity of the stellar population, whose spa- from 10 to 10 M , with an MZR which matches well the stel- tial distributions are studied as a function of Hubble type and to- lar metallicities of both Milky Way and LMC-like galaxies. tal stellar mass (M?). These properties were defined in previous articles in this series. For instance, the mean light weighted log stellar age is defined as 4.3.2. Galaxy averaged stellar metallicity X González Delgado et al.(2014c) obtained the important result hlog ageiL = xtZ × log t (1) that galaxy-averaged stellar ages, mass surface density, and ex- t,Z tinction are well matched by the corresponding values of these properties at R = 1 HLR and also with the values obtained from (Eq. (9) of Cid Fernandes et al. 2013), where x is the fraction of tZ the analysis of the integrated spectrum (i.e., the one obtained by flux at the normalization wavelength (5635 Å) attributed to the collapsing the datacube to a single spectrum). The general pat- base element with age t and metallicity Z. The mass-weighted tern, therefore, is that galaxy averaged properties match both the version of this index, hlog ageiM, is obtained by replacing xtZ by values at 1 HLR and those obtained from integrated spectra. Do its corresponding mass fraction mtZ. our stellar metallicities comply with this rule? While Cid Fernandes et al.(2013) average the base metal- To answer this question, we first define the galaxy-wide av- licities linearly (their Eq. (10)), as in González Delgado et al. erage stellar metallicity following Eq. (2) in González Delgado (2014a), we employ a logarithmic average: et al.(2014a), which gives the mass-weighted mean value of X hlog Z?,xyiM as hlog Z?iM = mtZ × log Z (2) t,Z P galaxy xy M?,xyhlog Z?iM,xy hlog Z?iM = P , (4) for the mass-weighted mean log Z? and xy M?,xy X hlog Z?iL = xtZ × log Z (3) where M?,xy is the stellar mass in spaxel xy. t,Z galaxy Figure4 compares our results for hlog Z?iM with the for the luminosity-weighted mean log Z . The motivation to use mass-weighted mean hlog Z?iM values obtained at R = 1 HLR ? HLR this definition is that the extended SSP bases used in this study (hlog Z?iM , bottom panels) and those derived from the inte- integrated span a much wider dynamical range in Z? (nearly three orders grated spectrum (hlog Z?iM , top panels), analyzed in the of magnitude, compared to barely one in our previous papers), exact same way as the individual zone spectra. Results are shown which is better handled with a geometric mean (implicit in the for both base GMe (left panels) and CBe (right).

A103, page 6 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence

evaluating and comparing the roles of morphology and total stellar mass in shaping the observed behavior of these four properties. Before discussing the results, we brieﬂy explain how these quantities are obtained and how they are presented.

2D maps in the CMD: using y we obtain, for each galaxy, 2D maps of each of the four properties. The results for all the galaxies are presented in the framework of the color-magnitude diagram, where each map is placed at the galaxy’s coordinates in the u − r vs. Mr CMD. Because absolute magnitude is related to M? and redder galaxies are (usually) older and more metal rich, these plots show the correlations M?-µ?, M?-age, and M?-metallicity in a 2D fashion. Because in our sample the galaxy Hubble type is correlated with color and luminosity, these plots not only show how the galaxy averaged properties and their radial structure change with the galaxy stellar mass, but also with the morphological type. These maps are shown in the Fig. 4. Upper panels: comparison of the galaxy-wide average stel- AppendixC (Figs. C.1 −C.4). lar metallicity (weighted in mass) derived from the spatially resolved galaxy Radial profiles: each 2D map is azimuthally averaged to study spectral analysis (hlog Z?iM ) and the integrated metallicity derived integrated the radial variations of each of the four stellar population prop- from fitting the total (integrated) galaxy spectrum (hlog Z?iM ). galaxy erties. Elliptical apertures 0.1 HLR in width are used to extract Lower panel: comparison of hlog Z?iM with the value measured at HLR the radial profiles, with ellipticity and position angle obtained R = 1 HLR (hlog Z?iM ). Left and right panels show results obtained with base GMe and CBe SSPs, respectively. All panels include from the moments of the 5635 Å flux image. Expressing radial 300 galaxies. The difference between the y-axis and x-axis is labeled in distances in units of HLR allows the profiles of individual galax- each panel as ∆, and the dispersion as σ. ies to be compared on a common metric, as well as averaging (“stacking”) them as a function of Hubble type or stellar mass. Radial profiles expressed in units of the HMR were also ana- The agreement is remarkable. The galaxy averaged metal- lyzed and lead to similar shapes, so all profiles presented below licity and at 1 HLR are the same to within a dispersion of use HLR as the unit for radius. 0.1 dex. The integrated metallicity also matches the galaxy averaged value with only slightly larger dispersions. The largest de- Radial gradients: inner and outer gradients are defined as differ- viations occur for low-metallicity systems. Similar conclusions ences between the values at R = 1 and 0 (5 ), and R = 2 and are reached if the comparison in Fig.4 is done using the light in 1 (5 ), respectively. For instance, weighted version of Eq. (4), out

P 5in log µ? = log µ?(1 HLR) − log µ?(0), (6) galaxy xy L?,xyhlog Z?iL,xy hlog Z?iL = P , (5) 5out log µ? = log µ?(2 HLR) − log µ?(1 HLR), (7) xy L?,xy h i h i where L is the luminosity (corrected by stellar extinction) in for log µ?, and similarly for log age L, log Z? M and AV . ?,xy Defined in this way, the gradients have units of dex/HLR each spaxel evaluated at a reference wavelength (5635 Å in our (mag/HLR for 5AV ). Since the stellar population properties of case). galaxies at 1 HLR represent very well the galaxy-wide average, The stellar metallicities behave as expected, in the sense 5in (5out) effectively measures how the bulge (disk) properties that like other properties, their galaxy-wide averages match the change with respect to those of the galaxy as a whole3. values at R = 1 HLR, and also the values derived from in- Unless otherwise noted, all results reported below are for tegrated spatially unresolved spectroscopy (González Delgado the GMe base, although the whole analysis was carried out with et al. 2014c). We thus conclude that galaxy-wide spatially av- properties derived with both sets of SSP models discussed in 3.1. eraged stellar population properties (stellar mass, mass surface Differences between GMe and CBe SSPs are as: (a) The stel- density, age, metallicity, and extinction) match those obtained lar mass surface density is lower with CBe than with GMe by from the integrated spectrum, and that these spatially averaged 0.27 dex on average, mostly due to the different IMFs (Salpeter properties match those at R = 1 HLR, proving that effective radii in GMe versus Chabrier in CBe). (b) Variations in stellar ex- are really effective (González Delgado et al. 2014b). tinction are negligible. (c) CBe yields somewhat younger ages and higher metallicities than GMe, by an average of 0.14 dex 5. Spatially resolved stellar population properties in hlog ageiL and 0.12 dex in hlog Z?iM. These shifts reflect the as a function of morphology and mass age-metallicity degeneracy, and are mainly a consequence of the different sets of metallicities available in these bases. However,

This section presents a series of results derived from our spa- 3 tially resolved spectral synthesis analysis of CALIFA galax- Based on the exponential fit analysis developed by Sánchez et al. ies. We focus on the following four stellar populations proper- (2013) and Sánchez-Blázquez et al.(2014), we conclude that the regions between 1 and 2 HLR are dominated by the disk component; ties: mass surface density (µ∗, Sect. 5.1), mean ages (hlog ageiL, thus, 5 measures the disk gradient. However, 5 is not measuring Sect. 5.2), metallicities (hlog Z i , Sect. 5.3), and extinction out in ? M the bulge gradient. The reason is that the effective radius (Re) of the (AV , Sect. 5.4). Each of these properties is studied by means of spheroidal component can be smaller than 1 HLR, and it shows a de- (i) 2D maps of the individual galaxies; (ii) radial profiles; and pendence with the morphological type. Thus, Re ∼ 1 HLR for E, but is (iii) radial gradients. Throughout the section, the emphasis is on significantly smaller in late-type spirals.

A103, page 7 of 44 A&A 581, A103 (2015) radial gradients are not aﬀected by this degeneracy. A detailed in our sample are populated mainly by E and S0, and some comparison of properties derived with the two bases is given in Sa. On the other hand, these early types are all massive, with 11 AppendixB. M? ≥ 10 M .

5.1. Stellar mass surface density 5.1.3. Radial gradients

2D maps of the stellar mass surface density for the 300 indi- Inner (0–1 HLR) and outer (1–2 HLR) gradients in log µ?, as de- vidual galaxies of our sample are presented in the AppendixC fined by Eqs. (6) and (7), are plotted as a function of morphology (Fig. C.1). Here we discuss the radial structure of log µ? as a and stellar mass in Fig.7. 5in log µ? values (corresponding to function of Hubble type and M?. the core region) are plotted in grey-red, while 5out log µ? (which trace the disks of spirals and S0 and the extended envelope of ellipticals) are plotted in grey-blue. Circles and stars show re- 5.1.1. µ?-morphology and µ?-mass relations sults for bases GMe and CBe, respectively, illustrating that the resulting gradients are nearly identical even though these bases Figure5 shows how µ? measured at 1 HLR changes with Hubble type (left panel), and with the galaxy stellar mass (right). Recall yield different absolute values of µ?. from González Delgado et al.(2014c) that properties measured A clear correlation exists between 5in log µ? and Hubble at 1 HLR match the corresponding galaxy-wide average value type. The gradient in the inner HLR increases (in absolute val- very well, so these plots ultimately show how the global µ? de- ues) significantly from late to early spirals, converging to a con- pends on the morphology and on M?. stant value for E and S0. This relation reflects the variation of HLR the bulge to disk ratio in spirals, and the dominance of the bulge The plot shows hlog µ? i increasing from late spirals to spheroids, with average and dispersion values of 3.1 ± 0.2, component in spheroids (S0 and E). The outer gradient is weaker 3.10 ± 0.18, 3.05 ± 0.25, 2.70 ± 0.17, 2.65 ± 0.24, 2.40 ± 0.28, (smaller in absolute value) than the inner gradient, as expected 2.04±0.27, for E, S0, Sa, Sb, Sbc, Sc, and Sd, respectively. Note if a disk component dominates the mass outward of 1 HLR. that E and S0 are remarkably similar. The right panel of Fig.7 shows the relation between the inner gradient and stellar mass. There is a clear increase (in absolute Surface densities also increase with M?, as seen in the right HLR values) of 5in log µ? with M?, with the more massive galaxies panel of Fig.5. The overall µ? −M? relation is relatively smooth, with no evidence of an abrupt change of behavior as having a steeper increase of the central density. The dispersion that discussed by Kauffmann et al.(2003) for SDSS galaxies. with respect to the average values (black crosses) within M?- Figure5, however, reveals that morphology is also behind the bins is significant. To check the effect of morphology on this dispersion we have averaged 5 log µ? in mass intervals for each dispersion in the µ?−M? relation. The black line shows the rela- in Hubble type and plotted the resulting averages (large colored tion for the full sample, obtained by averaging log µ? in 0.4 dex- wide bins in mass, while the big circles break this general re- circles). The general trend that emerges is that, for galaxies of the lation into different (color-coded) morphological types for the same mass, early type galaxies tend to be overall centrally denser same mass bins. Despite the reduced statistics, it is evident that: than later types, in agreement with Fig.5; although, there are a (a) for the same stellar mass, early-type galaxies are denser than few intervals of stellar mass (e.g. log M? = 11.4 M ), in which late-type galaxies; and (b) Sa and earlier type galaxies exhibit the variations in 5in log µ? with Hubble type are not significant. It is also worth mentioning that 5 log µ in Sa and Sb is a much flatter µ?−M? relation than later types. The overall im- in ? pression from these results is that morphology, and not only stel- very close to that in S0 and E, and in this sense it would be easy lar mass, plays a fundamental role in defining stellar surface den- to fade early-type spirals into S0’s. sities, and it is responsible for the change of slope in the SDSS µ?−M? relation. 5.2. Ages of the stellar populations 2D maps of the luminosity-weighted mean log stellar ages 5.1.2. Radial profiles (Eq. (1)) for the 300 galaxies are presented in Fig. C.2. Here we discuss the radial structure of hlog agei and its relation to Azimuthally averaged radial profiles of log µ? are shown in L Fig.6. Results are stacked by Hubble type (left panel) and mass Hubble type and M?. The presentation follows the same script (right). In the left panel galaxies are grouped in our seven mor- used in the presentation of µ?-related results in Sect. 5.1. phological classes. The typical dispersion within these bins is illustrated by the error bar, which shows the standard deviation 5.2.1. Age-morphology and age-mass relations in log µ? (R = 1 HLR) for galaxies of the Sa class. A clear trend with Hubble type is seen: the µ?(R) profiles Figure8 shows how the mean age of the stellar populations at scale with Hubble type from late to early spirals, and this mod- 1 HLR changes along the Hubble sequence (left panel), and with HLR HLR ulation with morphology is preserved at any given distance. E the galaxy stellar mass (right). Similar to log µ? , hlog ageiL and S0 have remarkably similar profiles, with core and extended represents well the galaxy-wide averaged stellar population age galaxy envelope equally dense at any given distance, suggesting that the (hlog ageiL , González Delgado et al. 2014c). HLR disk of S0 galaxies and the extended envelope of ellipticals have Clearly, hlog ageiL scales with Hubble type, increasing grown their mass by similar processes. steadily from Sd to Sa. S0 and ellipticals have stellar populations The right panel of Fig.6 shows the radial profiles grouped of similar mean age, and they are older than spirals. The aver- HLR in seven bins of stellar mass spanning the log M?(M ) = age and dispersion values of hlog ageiL (yr) are 8.62 ± 0.22, 9.1−11.8 range. These also show that the average of log µ?(R) is 8.89 ± 0.22, 9.07 ± 0.19, 9.33 ± 0.21, 9.55 ± 0.19, 9.71 ± 0.11, modulated by M?. However, this µ?(R)–M? modulation breaks and 9.74 ± 0.11, for Sd, Sc, Sbc, Sb, Sa, S0, and E, respectively. for early-type galaxies (concentration index C (r90/r50) ≥ 2.8; Mean ages also increase with the galaxy mass (right panel of see also Fig. 13 in González Delgado et al. 2014c), which Fig.8), a “downsizing” behavior that has been confirmed with

A103, page 8 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence

Fig. 5. Left panel: stellar mass surface density measured at 1 HLR as a function of Hubble type. Small dots represent log µ? for each galaxy; the colored circles are the average log µ? for each Hubble type, and the error bars are the dispersion in log µ? for each morphological type. Right panel: log µ?− log M? relation. Individual galaxies are represented by small dots colored by their morphological type. The black line is the average log µ? in galaxy stellar mass bins of 0.4 dex. Large colored circles are the average log µ? in each bin of mass for each Hubble type.

Fig. 6. Left: radial proﬁles (in units of HLR) of the stellar mass surface density obtained with base GMe. The results are stacked in seven morphology bins. The error bar in the panel indicates the dispersion at one HLR distance in the galaxies of the Sa bin. It is similar for other Hubble types and radial distances. Right: radial proﬁles stacked in seven bins of galaxy stellar mass, log M?(M ): 9.1−9.6, 9.6−10.1, 10.1−10.6, 10.6−10.9, 10.9−11.2, 11.2−11.5, 11.5−11.8.

HLR widely different samples and methods. For instance, our age- circles represent the mass-binned average hlog ageiL for each mass relation is similar to that derived for SDSS galaxies by Hubble type. As with the µ?−M? relation, breaking the age- Gallazzi et al.(2005, their Fig. 8;). They found that there is a mass relation into morphological types reveals clean trends. In 10 4 transition at M? ∼ 3 × 10 M , below which galaxies are typi- this case, we see that, for a fixed M?, earlier type galaxies are cally young and above which they are old. This is the same mass older than later types. The corollary is that mass is not the sole at which Kauffmann et al.(2003) find the µ?−M? relation to property controlling the SFH of a galaxy. In fact, given the ∼flat flatten. age-mass relations for Sa, S0, and E, morphology seems to be a Unlike in these SDSS-based works, we do not see sharp tran- more relevant factor, at least in these cases. sitions as a function of M? in neither µ? nor hlog ageiL, although differences in sample selection and statistics prevent a proper 5.2.2. Radial profiles comparison. We do, however, note a common behavior in the Figure9 shows the age radial profiles obtained by stack- right panels of Figs.5 and8, in the sense that the dispersion ing galaxies as a function of Hubble type and mass. The ∼ 10 M above 10 is strongly related to morphology. hlog ageiL(R) profiles scale with Hubble type, but by different Like in Fig.5 (right panel), the black line in the right amounts at the center than at 1 HLR. At any radial distance, panel of Fig.8 shows the age-mass relation for the whole sam- however, the early-type galaxies are older than later-ype galax- HLR ple, obtained by averaging hlog ageiL values in M? bins. ies. E and S0 are again very similar at all radii. This suggests Small dots show individual galaxies, while the large colored that E and S0 have similar histories not only on average, but also spatially resolved, at least in the inner 2 HLR. Negative age gra- 4 10 Equivalent to ∼5.5 × 10 M for our IMF. dients are detected in all galaxies (except perhaps in Sd, whose

A103, page 9 of 44 A&A 581, A103 (2015)

Fig. 7. Left panel: correlation between the inner (grey–red) and outer (grey–blue) gradient of log µ? and the morphological type. The results are shown for the GMe (stars) and CBe (circles) SSP models. The inner gradient is calculated between the galaxy nucleus and 1 HLR, and the outer gradient between 1 HLR and 2 HLR. Right panel: correlation between the inner gradient of log µ? and the galaxy stellar mass. Small dots represent the results for each galaxy, and black crosses the average for each 0.3 dex mass bin. Large circles represent the averaged inner gradient in mass intervals of 0.3 dex for each color-coded morphological type. Black crosses show the average correlation between the inner gradient of log µ? and galaxy mass independent of the morphological type.

Fig. 8. hlog ageiL measured at 1 HLR as a function of Hubble type (left) or galaxy stellar mass (right). Symbols and colors are the same as in Fig.5. The black line is the average hlog ageiL in galaxy stellar mass bins of 0.4 dex.

5 ages profiles are flatter than in the other spirals ). These negative profiles between galaxies of the same M? and different Hubble gradients reflect the inside-out growth of galaxies. Furthermore, type is significant, and larger than between the hlog ageiL(R) pro- the decrease of hlog ageiL with R indicates that quenching hap- files of galaxies of different M? but the same Hubble type. These pens earlier at the galaxy center; and also earlier in early-type results, in agreement with Fig.8, indicate that the age profiles are galaxies (spheroids and Sa) than in later-type spirals (Sbc–Sc). more related to morphology than to M?. Since hlog ageiL(R) is The radial profiles also show a clear trend with M? (Fig.9, essentially a first moment of the spatially resolved SFH, we can right), with the more massive galaxies being older everywhere, conclude that the SFH and its radial variation are modulated pri- hence preserving the downsizing pattern at all radial distances. marily by the galaxy morphology, with mass playing a secondary Comparing the left and right panels in Fig.9, one sees that role. grouping galaxies by their stellar mass leads to a reduced ver- 5.2.3. Radial gradients tical stretch in their hlog ageiL(R) profiles than when the averaging is done by morphology. But the profiles expand to a similar Gradients in hlog ageiL, computed as indicated in Eqs. (6) vertical scale if galaxies earlier than Sd and more massive than and (7), are plotted in Fig. 10 against Hubble type (left panel) 9.6 10 M are considered indicating that the effect of morphology and stellar mass (right). The figure layout is exactly as in Fig.7. and stellar mass are not easily disentangled here. However, in Whilst in that plot results obtained with bases GMe and CBe Sect. 6.3, Fig. 20 shows that the dispersion in the hlog ageiL(R) (circles and stars in the left panel, respectively) could hardly be distinguished, here the results for these two sets of SSPs do not 5 overlap so precisely, although the differences in 5hlog ageiL are The small drop of hlog ageiL toward the center of Sd galaxies is caused by a couple of galaxies with young nuclear regions. Given that clearly very small (see Sect. B.2). this group is the least populated in our analysis (only 15 galaxies), better A clear relation exists between 5inhlog ageiL and morphol- statistics is needed to evaluate the reality of this feature. ogy: the inner age gradient increases from early-type galaxies to

A103, page 10 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence

Fig. 9. Radial proﬁles of hlog ageiL as a function of Hubble type (left) and in seven bins of galaxy stellar mass (right). These bins are log M?(M ) = 9.1−9.6, 9.6−10.1, 10.1−10.6, 10.6−10.9, 10.9−11.2, 11.2−11.5, 11.5−11.8. Symbols and colors are the same as in Fig.6. These results are obtained with base GMe.

Fig. 10. Left panel: as Fig.7 except for hlog ageiL. The inner gradient shows a clear dependence with Hubble type, that seems to be stronger than with the galaxy mass. Sb–Sbc–Sc galaxies have larger inner gradients with CBe than with GMe, but both sets of models show a similar dependence with Hubble type. Right panel: the inner gradient of hlog ageiL as a function of galaxy mass. Colors and symbols are the same as in Fig.7.

Sb–Sbc spirals, which are the galaxies with the largest variation having the largest variations in hlog ageiL between the central between the age of the stellar population at the bulge and the core and the disk. This bimodal behavior seen in Fig. 10 suggests disk. Spirals of later type (Sc and Sd) have flatter radial profiles that other physical properties are also important in establishing than Sb–Sbc. The outer (between 1 and 2 HLR) age gradient the spatial variation of the SFH, which on the other hand is re- shows a similar bimodal behavior as 5in hlog ageiL, but with a flecting the different bulge formation processes along the Hubble smaller amplitude. sequence. The right panel of Fig. 10 shows the behavior of 5 hlog agei with M . The gradient tends to increase (become in L ? 5.3. Stellar metallicity more negative) from low-mass galaxies (which have roughly flat 11 profiles) up to about 10 M , at which point the trend reverses Figure C.3 presents the images of the mass-weighted mean (log- and 5inhlog ageiL decreases with increasing M?. This is best arithmic) stellar metallicity (cf. Eq. (2)). Here we discuss the ra- seen following the black crosses, which trace the mass-binned dial structure of hlog Z i as a function of Hubble type and M . mean relation. The dispersion with respect to this relation is sig- ? M ? nificant and is related to the morphology, as seen through the large colored circles. The tendency is that, at a given mass, S0 5.3.1. Metallicity-morphology and mass-metallicity relations and early-type spirals have weaker 5inhlog ageiL than Sb–Sbc. This dependence of age gradients with the Hubble type at a fixed Figure 11 shows how the stellar metallicity measured at 1 HLR M? indicates again that the spatial variation of the SFH is mainly changes with the Hubble type (left panel) and with the galaxy driven by morphology and not by stellar mass. stellar mass (right). However, the morphology (understood as the B/D ratio Stellar metallicities grow systematically from late to early- Graham & Worley 2008) cannot be the only driver of the spatial type galaxies. The statistics within each Hubble type are HLR variation of the SFH along all the Hubble sequence. Figure 10 hlog Z?iM (Z ) = −0.05 ± 0.13, −0.05 ± 0.33, −0.21 ± 0.16, shows that there is not a monotonic relation between the B/D −0.10 ± 0.18, −0.05 ± 0.15, +0.06 ± 0.08, and +0.10 ± 0.08 for ratio and 5inhlog ageiL, with galaxies with the smaller B/D ratio Sd, Sc, Sbc, Sb, Sa, S0, and E, respectively.

A103, page 11 of 44 A&A 581, A103 (2015)

pivotal position and late-type spirals with the flattest gradients, at least in a statistical sense. The right panel of Fig. 13 shows 5inhlog Z?iM as a function of M?. The dispersion is significant, but on average there is a tendency to turn flat profiles into negative gradient profiles as M? 9 10 increases from 10 to 10 M . The largest gradients are found 10 11 between 10 and 10 M . More massive galaxies tend to have weaker stellar metallicity gradients. The dispersion is significant throughout this relation. A trend with morphology is seen in the sense that, for a given mass, early types are those with weaker gradients.

Fig. 11. hlog Z?iM measured at 1HLR as function of Hubble type (left) and galaxy stellar mass (right). Symbols and colors are the same as in Fig.5. The black line is the average hlog Z?iM obtained in 0.4 dex bins 5.4. Stellar extinction of log M . ?  models the stellar extinction as a foreground screen, parametrized by AV and following the Galactic reddening law. A Not surprisingly, metallicities also grow with M?, as shown Images showing the spatial distribution of V for our 300 galax- in the right panel of Fig. 11. Since we have shown in ies are presented in Fig. C.4. Here we present AV related results Sect. 4.3.2 that the galaxy-wide average stellar metallicity is as a function of Hubble type and M?, following the same script well represented by the metallicity at 1 HLR, this plot is in fact adopted in the discussion of µ?, hlog ageiL, and hlog Z?iM in the equivalent to the global mass-stellar metallicity relation (MZR). previous subsections, thus completing the list of stellar popula- We have previously found that this relation is steeper than that tion properties studied in this work. Unlike masses, ages, and derived from HII regions, which is similar to the flatter stellar metallicities, extinction is more easily affected by inclination ef- MZR obtained when we consider only young stars (González fects, so the results reported below should be interpreted with Delgado et al. 2014a). As in Fig.5, the smoothed black curve is caution. Section 5.5 explores this issue in depth. HLR obtained by averaging hlog Z?iM in 0.4 dex bins of log M?. The dispersion in the MZR is significant, and larger than the 5.4.1. Extinction-morphology and extinction-mass relations dispersion produced by the galaxy morphology as shown by the distribution of large colored circles. These circles are the aver- Figure 14 shows how the stellar extinction at 1 HLR changes HLR age hlog Z?iM in each mass bin for each Hubble type, and with Hubble type (left panel) and with stellar mass (right panel). show the tendency of earlier type galaxies to be more metal rich HLR As with other properties, AV represents well the mean ex- than late-type galaxies of the same stellar mass. 6 tinction of the galaxy as well as the AV value derived from spectral fits of the integrated spectra7. The left panel in Fig. 14 HLR 5.3.2. Radial profiles shows AV as a function of morphology. Ellipticals and S0s HLR have almost no extinction, with mean AV = 0.01 ± 0.01, and Figure 12 shows the results of stacking the radial profiles of 0.06±0.07 mag, respectively. Sa, Sb, and Sc galaxies have AHLR h i V log Z? M as a function of Hubble type and M?. Outwards de- around 0.25 mag, and somewhat smaller (0.19 ± 0.08 mag) in h i creasing log Z? M is detected for most morphological classes, Sd’s. but flat profiles are found for Sc–Sd galaxies. Intermediate type- spirals (Sb–Sbc) stand out as those with the largest variations in There is no clear behavior of stellar extinction with galaxy ≤ 11 stellar metallicity. stellar mass. In general, galaxies with M? 10 M have AHLR = 0.2–0.3 mag. More massive galaxies are less The behavior of the radial variation of the stellar metallic- V ity with M (right panel in Fig. 12) is similar to the behavior extinguished, and for fixed mass early types tend to have ? smaller AHLR, but the dispersion is large. with morphology. Most galaxies have hlog Z?iM that decreases V with R, except for the two lowest mass bins, which show flat profiles. The largest spatial variations are also found in galaxies in 5.4.2. Radial profiles the intermediate mass bins (10 ≤ log M?(M ) ≤ 11). These negative radial gradients of the metallicity are also an Figure 15 shows AV (R) profiles stacked by Hubble type (left indicator of the inside-out formation processes in galaxies. The panel), and mass (right). Spirals have AV ∼ 0.2 mag in the disk inversion of the gradient in late-type spirals and in low-mass spi- and up to 0.6 mag in the center. Their AV profiles are similar rals may be an indicator of the secular processes or the outside-in for all Hubble types except for Sd’s, where AV does not change formation scenario in these galaxies (Pérez et al. 2013). much from disk to center. Ellipticals and S0’s also show negative AV gradients, although at distances larger than 1 HLR they are almost dust-free. The radial profiles in different bins of M? (right 5.3.3. Radial gradients panel) show a similar behavior to that with morphology. Except

Figure 13 clones Figs.7 and 10 for hlog Z?iM gradients. On the for the most massive bins, shifted to lower extinction values, all left panel, one sees that, as for stellar densities and ages, results other mass-binned AV (R) profiles are similar. for bases GMe and CBe are very similar. On average, galaxies 6 have hlog Z?iM gradients ∼−0.1 dex/HLR similar to the value The galaxy average extinction for each galaxy is calculated as the obtained from nebular oxygen abundances (Sánchez et al. 2013). mean of all the 20 radial values obtained for each galaxy between the Outer and inner gradients are not significantly different. Despite center and 2 HLR. 7 HLR galaxy the large scatter, there is a hint of a bimodal distribution as that The difference between AV and hAV i is −0.03 ± 0.06, while HLR integrated found for stellar ages, also with intermediate-type spirals in a between AV and AV is −0.0 ± 0.1. A103, page 12 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence

Fig. 12. Radial proﬁles of hlog Z?iM as a function of Hubble type (left) and of galaxy stellar mass (right). Mass bins are log M?(M ) = 9.1−9.6, 9.6−10.1, 10.1−10.6, 10.6−10.9, 10.9−11.2, 11.2−11.5, 11.5−11.8. Symbols and colors are the same as Fig.6. These results are obtained with base GMe.

Fig. 13. Left: as Fig.7 except for hlog Z?iM. Right: the inner gradient as a function of the galaxy stellar mass. Symbols and colors are the same as in Fig.7.

region. With the exception of Sd galaxies, spirals have 5in AV ∼ −0.25 mag/HLR.

On average, 5inAV gets stronger with increasing M? up 11 to 10 M (Fig. 16, right) and weakens toward higher mass, spheroid dominated systems. The dispersion with respect to the mass-binned relation (traced by the black crosses) is large, and not clearly related to morphology (coded by the colored circles). As a whole, and despite the general trends summarized above, of the four properties analyzed in this section, AV is the one for which tendencies with Hubble type and stellar mass are less clear. A major reason for this is that, unlike for µ?, Fig. 14. AV measured at 1HLR as function of Hubble type (left) and hlog ageiL, and hlog Z?iM, AV estimates are sensitive to inclina- galaxy stellar mass (right). Symbols and colors are the same as in Fig.5. tion eﬀects, as explained next. The black line is the average hlog Z?iM obtained in 0.4 dex bins of log M?. 5.5. Effect of inclination on the radial proﬁles

5.4.3. Radial gradients An implicit hypothesis throughout the analysis presented so far is that galaxy inclination does not affect our estimates of the stel- AV gradients are shown in Fig. 16, which is formatted as lar population properties and their radial distributions. One ex- Figs.7, 10 and 13. As for the previous properties, results pects this assumption to break down in the case of AV , which for bases GMe and CBe are nearly indistinguishable, as il- should increase from face on to edge on galaxies, although it is lustrated by the overlap of circles and stars in the left panel. not unreasonable to conceive that inclination effects propagate 5inAV and 5outAV show similar behavior with morphology, al- to the spectral synthesis-based estimates of stellar mass surface though the inner gradient is always higher than the outer gradi- densities, mean ages, and metallicities. It is therefore relevant to ent. In Ellipticals the gradient of AV only exists in the central evaluate if and how inclination affects our results.

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Fig. 15. Radial proﬁles of AV as a function of Hubble type (left), and in seven bins of galaxy stellar mass (right). These bins are log M?(M ) = 9.1−9.6, 9.6−10.1, 10.1−10.6, 10.6−10.9, 10.9−11.2, 11.2−11.5, 11.5−11.8. Symbols and colors are the same as Fig.6. These results are obtained with base GMe.

Fig. 16. Left: as Fig.7 except for AV . Right: the inner gradient as a function of galaxy stellar mass. Symbols and colors are the same as in Fig.7.

In order to do so, we have divided the 300 galaxies in three happens, however, that the sub-group of b/a ≤ 0.39 Sc’s has a subsamples on the basis of the b/a ratio (minor to major isopho- mean stellar mass 0.4 dex lower than other Sc’s, which could tal axes), as measured in SDSS R-band images. The three sub- explain their lower metallicities without implying inclination samples, each containing 100 galaxies, cover (i) b/a ≤ 0.39, effects. edge on; (ii) 0.39 < b/a ≤ 0.63; and (iii) b/a > 0.63, face on. The one property that does vary systematically with b/a is Galaxies in each subsample were grouped by Hubble type, and AV , and it does so in the expected sense: spirals with lower b/a their radial profiles of log µ?, hlog ageiL, hlog Z?iM, and AV av- have larger extinction. This is particularly evident in Sb’s. This eraged as previously done for the whole sample in the left panels dependence hinders the interpretation of the stacking results pre- of Figs.6,9, 12, and 15. sented in Sect. 5.4, and explains why no clean tendencies of the Figure 17 shows the resulting stacked profiles of log µ?, AV values and profiles with morphology and stellar mass were hlog ageiL, hlog Z?iM, and AV . Solid, dashed and dotted lines identified. show profiles for the “face-on” (b/a > 0.63), intermediate in- . < b/a ≤ . b/a ≤ . clination (0 39 0 63), and “edge-on” ( 0 39) 6. Discussion samples, respectively, and each column is for one of the seven Hubble types used throughout the paper. Average profiles were This section is divided into four main parts. First we summarize only computed for morphology-inclination bins containing at our results in the context of related studies. We then discuss our least four galaxies. findings in the context of the growth of galaxies – theoretical Stellar mass surface density, age, and metallicity profiles expectations, high-redshift observations, and previous results of show a negligible variation among the b/a-based subsamples. inside-out growth for CALIFA galaxies. In the third part, we ex- This result indicates that inclination does not affect the estimates plore what the results tell us about the quenching of star forma- of these properties in any systematic way. Any difference is at tion in galaxies. Finally, we discuss the theoretical predictions a level not significant insofar as this paper is concerned. The for the radial variations of age and metallicity in early types and exception is the hlog Z?iM profiles for “edge-on” Sc’s, which in spirals from different simulations of galaxy formation. We differ substantially from the profiles of less inclined Sc’s. It so compare our results for variations of the radial structure in the

A103, page 14 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence

Fig. 17. Radial profiles of log µ?, hlog ageiL, hlog Z?iM, and AV (from the upper to the bottom panels) for the different Hubble types (from E to Sd from left to right panel) and three different bins of the ratio of minor to major photometric radius of the galaxy: solid line (0.63 < b/a, face on), dashed line (0.39 ≤ b/a ≤ 0.63), and dotted line (b/a < 0.39, edge on). inner (R ≤ 1 HLR) and outer (1 ≤ R ≤ 3 HLR) parts with other and (ii) in the outer part of galaxies, values averaged in a ring observational results in the literature. at 1.5 ± 0.1 HLR, as representative of disks. Figure 18 plots the individual results as small dots; large dots represent the average values for each (color-coded) Hubble type, and the error bars 6.1. Age and metallicity of bulges and disks show the dispersion. While hlog ageiL gives information about the globally “averaged” SFH, hlog agei informs when most of The analysis of SDSS data has generated a general knowledge of M the stellar mass was formed. how the age and metallicity of galaxies change with M?, color, or concentration index (e.g. Kauffmann et al. 2003; Gallazzi Figure 18 shows that the bulges of Sa–Sb and the cores of et al. 2005; Mateus et al. 2006). These studies have confirmed E-S0 formed at a similar epoch; they are very old (∼10 Gyr) and that, in general, early-type galaxies are old and metal rich, while metal rich (∼>1 Z ). Thus, they probably formed through com- late-type galaxies are younger and more metal poor. Numerous mon processes, that occurred rapidly and early on. However, (single spectra or longslit) studies have reported also that ellip- the bulges in Sc–Sd galaxies (shown as the two darkest shade ticals are metal rich, and have a range of ages, 2–10 Gyr, which of blue) are younger and have a more extended SFH (both depend on stellar velocity dispersion (e.g. Trager et al. 2000; hlog ageiM and hlog ageiL are smaller), and have lower stellar Thomas et al. 2005; Sánchez-Blázquez et al. 2006; Graves et al. metallicities. Thus, Sc–Sd galaxies formed in later times and/or 2009; Johansson et al. 2012). with different processes. Our spatially resolved analysis introduces a significant im- Many bulges in late-type spirals are in fact pseudo–bulges. provement in the study of the structure of galaxies. For exam- Unlike true bulges, pseudo–bulges are thought to grow by secu- ple, we compute ages and metallicities of bulges in disk galaxies lar processes, where material from the disk (by the loss of angu- and compare them with elliptical galaxies in a systematic way, lar momentum) is the main source of star formation in the cen- avoiding problems derived from the lack of spatial resolution tral parts of these galaxies (e.g. Kormendy & Kennicutt 2004). prevalent in some of the previous studies. We may see this effect at work in Figs.9 and 12, as a flattening We compute the luminosity-weighted and the mass-weighted of the radial profiles of hlog ageiL and hlog Z?iM, and the posi- age and metallicity: (i) in the central part of galaxies (val- tive hlog ageiL gradient in the core of Sc galaxies. Some effects ues averaged within 0.25 HLR) as representative of the stellar of the secular processes due to the disk may also be present in population properties of bulges and central core of ellipticals; the bulges of Sa–Sb. For example, Fig. 18 shows that bulges of

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bulge disk Fig. 18. Left panel: mass-weighted age at the galaxy center (hlog ageiM ) and at 1.5 HLR (hlog ageiM ) for the diﬀerent Hubble types. Small dots are the individual radial points, while big colored dots represent mean values for each Hubble type, with the error bars indicating the dispersion in the mean. Middle panel: as in the left panel but for the light-weighted age (hlog ageiL). Right panel: as in the left panel for hlog Z?iM. The top horizontal axes in the left and middle panels show the redshift scale. The diagonal line in the three panels is the bisector.

Sa–Sb have hlog ageiL ∼ 6 Gyr, younger than the 10 Gyr epoch of formation derived from hlog ageiM; and this may be understood if some disk stars are rearranged into the bulges or if dis- sipation processes bring gas triggering new star formation in the center. Figure 18 also shows that disks are younger and more metal poor than bulges. Both hlog ageiL and hlog ageiM are lower in disks than in their respective bulges, indicating that disks formed later than bulges, and that star formation continues for a longer time in disks that in bulges, probably as a consequence of a con- tinuing availability of gas in disks (Roberts & Haynes 1994). This indicates a general scenario of inside-out formation.

6.2. Inside-out growth of spheroids and spirals Fig. 19. Radial profiles (in units of HLR) of the mass-weighted age, 6.2.1. Theoretical expectations and recent results from high hlog ageiM, obtained with GMe base. The results are stacked by mor- redshift galaxies phological type as in Fig. 7. Models of galaxy formation predict a common inside-out view for the mass assembly in galaxies (e.g. Kauffmann et al. 1993; Aumer & White 2013). First, the bulge formed at high redshift; other evidence that galaxies rapidly assemble inside-out at z = 1 then, the disk was built around the bulge in the case of spirals. (Nelson et al. 2012; Szomoru et al. 2010, 2012); while Hammer In the case of ellipticals, the central core formed at z ≥ 2, and et al.(2005) find evidence that MW-like galaxies have rebuilt the envelope grew later through minor mergers (e.g. Oser et al. their disk at z ≤ 1 in a major merger epoch that drastically 2010; Hilz et al. 2013). Observational evidence come from the reshapes their bulges and disks, and is consistent with earlier significant size evolution in early-type galaxies (ETG), which cumply evolution. 2 grow in size ∝M? (van Dokkum et al. 2010; Patel et al. 2013). In summary, there is mounting evidence of the major pro- More recently, van Dokkum et al.(2014) find evidence cesses responsible for the assembly and shaping of galaxies at against the inside-out formation scenario for spirals. For a sam- different epochs, and these are complemented with a variety of ple of MW-like spirals at redshift z = 2.5, they estimate the de- processes that modify the inside-out formation scenario: stellar migration, bar induced gas inflows, gas-rich minor merger, angu- pendence of the radius with M?, and find that their size-M? re- lar momentum loss due to reorientation of the disk, infall of gas lation is similar to the size-M? of similar galaxies at z = 0. They conclude that the mass growth took place in a fairly uniform way, with misaligned angular momentum, etc. (Aumer et al. 2014). with the galaxies increasing their mass at all radii, thus, their Reff barely grows. These results seem to be supported by numerical 6.2.2. CALIFA view of the inside-out growth of galaxies simulation by Elmegreen et al.(2008), which find that bulges can be formed by migration of unstable disks. Other observa- Our results favor an inside-out growth of spirals. Pérez et al. tional evidence come from the detection of clumpy star-forming (2013) studied the stellar mass growth as a function of the radius 10 11 disks in galaxies at z ∼ 2 (Genzel et al. 2008; Förster Schreiber and cosmic time in galaxies with 10 . M? . 5 × 10 M , et al. 2011), which may indicate an early build up of bulges by and showed that the nuclei grow faster than the inner 0.5 HLR, secular evolution. Thus, studies at high redshift are providing which, in turn, grow faster than the outer 1 HLR. This conclu- new results that draw a complex landscape of galaxy build up. sion is supported by the stellar age radial profiles presented in For example, Wuyts et al.(2011) also find clumpy disk star for- González Delgado et al.(2014c), and confirmed here in Fig. 10 mation, but at the same time conclude that there is a Hubble se- for most spirals and spheroidals. Further support comes from the quence in place at least since z ∼ 2.5. On the other hand, there is ratio HMR/HLR (Fig.3), a good probe of the spatial variation of

A103, page 16 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence the SFH in galaxies (González Delgado et al. 2014c). This ratio the age values and the gradients) to change more with M? than is lower than 1 (Fig.3), a fingerprint of the inside-out growth with Hubble type. On the contrary, if quenching is driven by found by Pérez et al.(2013). morphology, galaxies of similar stellar mass would have very Figure 19 shows how the radial profiles of hlog ageiM de- different hlog ageiL structure depending on Hubble type. We crease outward for all the Hubble types. Most of the stellar mass explore the relevance of morphology versus M? in Fig. 20: age in the center was formed 10 Gyr ago or earlier (z ≥ 2). However, radial profiles are shown as a function of M? and of morphology, at 1.5 HLR, hlog ageiM ranges from 7 Gyr (z ∼ 1) in E–S0 to in four mass bins (log M? = 11.5−11.2, 11.2−10.9, 10.9−10.6, 4.5 Gyr (z ∼ 0.4) in Sbc, suggesting that, both early type and 10.6−10.1). Clearly, morphology is the main driver: it can ac- MW-like, galaxies have continued accreting or forming in situ count for up to 0.75 dex change in age at a given mass (top pan- stars in their disks until more recent times, thus supporting the els); conversely, at a fixed morphology, mass accounts for less inside-out scenario in these galaxies. than 0.25 dex (bottom panels). Further, morphology accounts not This trend, however, changes beyond 1.5−2 HLR, where only for changes in absolute values, but also for changes in the hlog ageiM and hlog ageiL flatten. This may be interpreted as gradients at a given galaxy mass. indicating that the mass was formed in a more uniformly dis- This confirms the similar result obtained above with log µ?, tributed manner across the outer disk, or that stellar migration and it implies that galaxies of similar M? (equivalent to have shuffles inner born stars to the outer disk, washing out the inside- similar Mhalo) and with a large spheroid have shut down their out formation signs so clearly seen in the inner 1.5 HLR. In the star formation (outside their central core) earlier than galaxies case of E–S0, this may be understood if beyond 2 HLR most of of later morphology. These results indicate that the SFH and the stellar mass in the galaxies was already accreted at z = 1. their radial variations are modulated primarily by galaxy morphology, and only secondarily by the galaxy mass, suggesting that the bulge formation has a relevant role in quenching the star 6.3. Quenching formation in galaxies. Several mechanisms have been proposed to explain the shutdown of star formation in galaxies. Halo mass quenching is one of the most popular explanations of the bimodal distribu- 6.4. Radial structure of the stellar population properties tion of the properties of galaxies, and it is required to explain in ETG and their relation with galaxy formation models the green valley as a pathway towards quenching of star for- 6.4.1. Theoretical predictions from cosmological simulations mation in early and late type galaxies (e.g. Schawinski et al. 2014). In this scheme, galaxies with a halo mass below a critical Classical chemical evolution models of the formation of early- 12 value (a few ×10 M ) accrete cold gas conducive to star for- type galaxies (ETG) are based on two possible scenarios: 1) dis- mation. Above this critical mass, the infalling gas reaches the sipative formation, the well-known monolithic collapse; and sound speed and cannot form new stars (e.g. Cattaneo et al. 2) the non-dissipative collapse. These scenarios produce very 2006; Dekel & Birnboim 2006). The dependence with environ- different radial gradients of ages and abundances, being very ment and clustering strongly supports this quenching mechanism steep in the first case, with values of 5[Fe/H] ∼ −0.5 to (e.g. Weinmann et al. 2006; Peng et al. 2010). −1.0 [dex/dex] (Larson 1974, 1975; Carlberg 1984)8, but (al- The differential dependence of the stellar mass surface den- most) flat when there are pure stellar mergers. This second case sity (Figs.5 and6) with the galaxy stellar mass (a proxy of the may even erase a previously existing radial gradient. halo mass) provides further evidence of the halo quenching (e.g. The most recent cosmological simulations propose a two- Behroozi et al. 2010). Estimating of the properties of the stel- phase formation scenario for ETG’s in which the central core lar populations in SDSS galaxies, Kauffmann et al.(2003) found formed at z ≥ 2, and the envelope grows after this through 10 that there is a critical mass (M? = 3 × 10 M , equivalent to minor mergers (e.g. Naab et al. 2009; Oser et al. 2012; Hilz 10 ∼6 × 10 M for our Salpeter IMF) below which log µ? scales et al. 2012; Navarro-González et al. 2013). As follows: 1) galax- with M?, and above which log µ? is independent of the galaxy ies assemble their mass through dissipative processes and star stellar mass. The right panels of Figs.5 and6 support this sce- formation occurs in situ. Starbursts formed at the center as a nario because the radial profiles of log µ? scale with log M?, consequence, for example, of large major mergers or monolithic and furthermore they do so all along their extent. Our results collapse. The star formation is induced by cold flow of accre- also show that log µ? saturates at high M? because the high- tion or by gas-rich mergers. 2) Galaxies grow in size by mass mass end of the distribution is dominated by early type galaxies assembly through the external accretion of satellites; ex situ star (Sa–S0–E); this suggests that the spheroidal component plays a formation formed by dry mergers of galaxies toward the central significant role in the quenching of star formation in high-mass most massive galaxies. galaxies. Observationally, there is evidence of a significant size evo- 2 The importance of morphology in the quenching of galaxies lution in ETG. The growth of the galaxy size with M? supports has also been reported in the literature (e.g. Bell 2008; Bell et al. this proposal. A transition region is expected between the in situ 2012; Barro et al. 2013; Pan et al. 2014; Woo et al. 2015). Martig central core of ETG and ex situ outer regions. Since the central et al.(2009) found that the dependence of quenching with mor- core of these ETG is enriched very quickly due to major merg- phology is a consequence of the bulge-building mechanism. The ers at high redshift (z ≥ 2), and the satellites that are accreted steep potential well induced by the formation of a large spheroid are less metal rich than the central massive galaxies, a negative component results in the stabilization of the disk, which cuts the radial gradient of the metallicity is expected, even with a change supply of the gas, preventing its fragmentation into bound, star- of the slope in the transition region. Thus, values as 5[Fe/H] = forming clumps. Our results support this scenario, as it is ex- −0.5 [dex/dex] (Pipino et al. 2010) or 5[Fe/H] = −0.3 [dex/dex] plained below, because the dependence of the SFH of galaxies (Kawata & Gibson 2003) are predicted. with the morphology. If the halo mass is the main property responsible for quench- 8 The metallicity gradient measured in spheroids is traditionally calcu- ing, we should expect that the radial structure of hlog ageiL (both, lated as 4[Fe/H]/4log r and expressed in [dex/dex] units.

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Fig. 20. Radial profiles of hlog ageiL (upper panel) in four galaxy stellar mass bins. From left to right: log M?(M ) = 11.2−11.5 (continuum line), 10.9−11.2 (dashed line), 10.6−10.9 (dashed-point line), 10.1−10.6 (dotted line). In each panel, the average profile for each Hubble type is plotted if more than four galaxies have galaxy stellar mass in the log M? bin. Bottom: each panel shows the radial profile of each Hubble type averaged in each of the four log M?(M ) bins.

However, the merger history may change an existing radial are of lower metallicity than in the simulations with no winds. gradient: while dry major mergers can flatten the pre-existing In both cases, however, they predict a positive age gradient of gradient (Kobayashi 2004; Di Matteo et al. 2009; Hopkins et al. ∼0.03−0.04 [dex/dex]. 2009), dry minor mergers can steepen the metallicity gradient. Thus, Kobayashi(2004) SPH chemodynamical simulations of elliptical galaxies which include radiative cooling, star forma- 6.4.2. Implications from this work and comparison tion and feedback from SNII-Ia, and chemical enrichment (but with other results from the literature do not include kinematic feedback), found that the steep negative Following our own results, E and S0 have formed their cen- radial metallicity gradient, established during the initial starburst tral core at similar cosmic time since they have similar central at z ≥ 2, can flatten significantly by later major-mergers in the ages (see Fig. 18). Further they must have formed through sim- galaxy. Following these simulations, the average gradient at the ilar processes since their radial profiles of log µ?, hlog ageiL, present time is 5[Fe/H] = −0.3 [dex/dex], but it may decrease to and hlog Z?iM are remarkably similar. They both show small a value of −0.2 [dex/dex] when major mergers appear. but negative hlog ageiL, and hlog Z?iM gradients in the central core. In the central 1 HLR, E and S0 in our sample have Beside the merger history, feedback can change the inner 5hlog Z i ∼ −0.1 [dex/dex] (std = 0.15)9. Slightly steeper, and outer metallicity gradients. Thus, a strong AGN feedback ? M 5hlog Z i = −0.2 [dex/dex], when CBe models are used. These can stop the star formation in the central core of massive galax- ? M are within the range of values found in other studies based on ies, flattening the inner gradients. Feedback from in situ star long-slit or IFS data up to one effective radius (e.g. Mehlert formation can alter the outer metallicity gradient. Also, the ex- et al. 2003; Sánchez-Blázquez et al. 2007; Annibali et al. 2007; istence of galactic winds may modify the composition of the Rawle et al. 2008; Spolaor et al. 2010; Kuntschner et al. 2010). ISM in a galaxy. Hirschmann et al.(2015) performed cosmo- However, they are shallow compared with theoretical expecta- logical simulations, which include an empirical model for the tions if minor mergers are relevant in growing up the central core momentum driven galactic winds, to investigate the dependence of E and S0 galaxies. This may indicate that major mergers are of the age and metallicity outer gradients with metal cooling more likely the main process building the central regions (up and galactic winds; in principle, required to explain the mass- 1 HLR) of ETGs. metallicity relation, MZR. These simulations including winds predict 5[Fe/H] = −0.33 [dex/dex], steeper than the simula- 9 Our gradients, which are measured in a linear scale, are converted tions without winds that predict 5[Fe/H] = −0.11 [dex/dex]. here to a logarithmic scale to be compared with predictions from simu- The main explanation is that in wind models the stars accreted lations and other works in the literature.

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Between 1 and 3 HLR, the radial profile of hlog Z?iM is of ellipticals. Following these authors, massive galaxies accrete similar slope or slightly shallower than in the inner 1 HLR. We higher mass satellites, and because of their deeper potential do not find any evidence of a transition region where the metal- well they retain their own gas against stellar winds, producing licity radial profile steepens to metallicities below solar. If the a shallower metallicity gradient in the outer regions of massive 1 to 3 HLR envelope of ETG had grown through the accretion ellipticals. of low-mass satellites a steepening of metallicity would be ex- Because the metallicity at 2−3 HLR in E and S0 are simi- pected, as explained before, because the mass-metallicity rela- lar to the metallicity in the bulge of early spirals, and the stars tion implies that low-mass satellites would be of low metallicity. at these distances are as old as the bulges of Sa–Sb galaxies (see In our results there is no evidence either of an inversion of the Fig.9), the 1 −3 HLR envelope of early-type galaxies might have age radial profile toward older ages beyond 1−2 HLR, as ex- built from the centers of early-type spirals. In summary, the neg- pected if these satellites were formed very early on like the core ative but shallow gradients of the metallicity and ages suggest 11 of E and S0 (see Figs.9 and 18). These results are in contrast that massive (M? ≥ 10 M ) early-type galaxies built their in- with recent results by Greene et al.(2012, 2013): for a sample of ner 3 HLR through mergers with massive and metal-rich spirals. ∼30 early type galaxies they find at 2 Reff an old (10 Gyr) stellar population with [Fe/H] ∼ −0.5, and interpret this as the stellar outskirts of these galaxies being built up through minor merg- 6.5. Radial structure of the stellar population properties ers. Coccato et al.(2010), La Barbera et al.(2012), Montes et al. in spirals and their relation with galaxy formation models (2014) observing a few massive ellipticals reported a decline of New insights into the structure of the Milky Way disk, in par- the metallicity to under solar in an old stellar population in their ticular through the measurements of chemical abundances of outskirts (≥10 Reff) suggesting that these galaxies are formed in large sample of stars, are provided by the spectroscopic surveys two phases, the central core through major mergers, and through undertaken in recent times (e.g., SEGUE, RAVE, Gaia-ESO minor mergers farther out. However, other recent works show survey, HERMES, APOGEE, LAMOST, etc.; Yanny et al. examples of ETG with an old and metal-rich stellar population 2009; Steinmetz et al. 2006; Gilmore et al. 2012; Zucker and a very shallow metallicity gradient up to 3 Reff (Trujillo et al. et al. 2012; Majewski et al. 2010; Zhao et al. 2012). Radial 2014) in contrast with the results by Greene et al.(2012, 2013). Velocity Experiment (Steinmetz et al. 2006; Boeche et al. 2014, Our results do not support the minor merger scenario for the RAVE;) is studying the radial and vertical chemical gradients size growth of ETG. Thus, the ages hlog ageiL ∼ 9.7 (yr) and using a very large sample of dwarf stars. Close to the Galactic metallicity hlog Z?iM ∼ Z at 1−3 HLR, and the shallow metal- plane, RAVE shows a negative radial gradient of Fe abun- licity gradient, 5hlog Z?iM ∼ −0.1 [dex/dex], that we obtain are dance, −0.054 dex/kpc10, which becomes flatter or even pos- more consistent with the growth of the 1−3 HLR envelope of itive when measured above the disk. So, the [Fe/H] gradient ETGs through major mergers. ranges from −0.039 to +0.047 dex/kpc when measured at heights Another interesting result reported in the literature that can 0.4−0.8 kpc, or 1.2−2 kpc above the Galactic plane, respectively. be compared with ours is the correlation between the metallicity The radial gradient of abundances in the different regions of gradient and the galaxy mass (or stellar velocity dispersion, σ?) a spiral galaxy are important because they are directly related to found for E and S0. Spolaor et al.(2010) have found that the rela- the formation process. Obviously, not all scenarios of disk/spiral tion between σ? and the metallicity gradient shows two regimes formation are valid, since it is necessary that they produce a ra- −1 of behavior: (i) for galaxies with log σ? ≤ 2.2 km s , the metal- dial gradient of abundances in the disk but not in the halo, as licity gradient steepens with σ?; (ii) galaxies with log σ? > observed in the MW and in M 31. Thus, the formation of the 2.2 km s−1 (the most massive ellipticals), have a metallicity gra- halo from different fragments or minor mergers with very short dient that does not correlate with σ? (or galaxy mass), with free fall times does not create a radial gradient but a dispersion a mean value ∼−0.15 [dex/dex]. On the other hand, Pastorello of abundances, and therefore it was concluded early on that the et al.(2014) derive the [ Z/H] gradient in the outer 1−2.5 Reff of MW halo may be formed from mergers or from the accretion of a sample of ellipticals, and they find that the gradient covers a low-mass galaxies (or part of them). However, disks are more wide range of values, from negative very steep (∼−2) to flat or likely formed from a single cloud falling on and from inside out. even positive; these values correlate with the galaxy stellar mass and stellar velocity dispersion, with galaxies of lower M? (or 6.5.1. Theoretical predictions from “classical” chemical σ?) having the steeper metallicity gradient. However, the most massive galaxies exhibit the flattest gradients and an average evolution models value of 5hlog Z?iM ∼ −0.2 [dex/dex] (std = 0.38) for galax- 11 Most classical chemical evolution models claim that the infall of ies with M? ≥ 10 M . Both works show a significant scatter in 11 gas with a radial dependence, implying an inside-out scenario the 5hlog Z?iM−M? relation for M? ≥ 10 M , and reasonable for the disk formation, is essential to reproduce the observed doubts of the existence of the correlation for high-mass ellipti- radial gradient of abundances. The key ingredient is the depen- cals. dence of the disk infall time scale with the radial distance, which Our results (Fig. 13) indicate that there is no correlation makes the gas to accumulate faster in the inner disk. Since the between the metallicity gradient and M? for the CALIFA early- star formation rate depends on the gas density, these assumptions type galaxies (E and S0). Even so, our results are compati- produce a radial dependence of the star formation rate and a neg- ble with Spolaor et al.(2010) and Pastorello et al.(2014) be- ative radial metallicity gradient (Ferrini et al. 1994; Molla et al. 11 cause E and S0 in our sample are all above 10 M (and 1996; Chiappini et al. 2001; Mollá & Díaz 2005). Thus Molla −1 σ? > 100 km s ), for which no correlation is found between the et al.(1996) give a value −0.08 for a MW-like galaxy, repro- metallicity gradient and M? in Spolaor et al.(2010) or Pastorello ducing the value found by Shaver et al.(1983) and other H  re- et al.(2014). We find that CALIFA ellipticals have a shallow gions studies. Chiappini et al.(2001) models predict a gradient gradient. This behavior is also in agreement with Hirschmann et al.(2013) simulations and their interpretation of the lack of 10 The metallicity gradient in disks is traditionally calculated as correlation between the metallicity gradient and M? in massive 4[Fe/H]/4r, and expressed in dex/kpc.

A103, page 19 of 44 A&A 581, A103 (2015) of −0.04 dex/kpc for a MW-like galaxy. In fact, as theoretical populations. For these galaxies, they find a metallicity gradient equations show (Goetz & Koeppen 1992), the radial gradient of of −0.025 dex/kpc (std = 0.05), equal (−0.027 dex/kpc) to the abundances appears in the disk when there is an adequate ratio gradient that we derive for the same group of galaxies. between star formation rate to infall rate. It also implies, there- To compare the RaDES simulated galaxies (Few et al. 2012) fore, that a dependence of the radial gradient on the morpho- with our results, 19 galaxies of the Pilkington et al.(2012a) logical type of galaxies may exist. Molla et al.(1996) models sample, have been analyzed, measuring the HLR and the gra- already predicted radial gradients for galaxies of different mor- dients in a similar way as we have done with the CALIFA phological types, with values in the range −0.025 (for a M 31- galaxies (Ruiz-Lara, in prep.; Ruiz-Lara et al., priv. comm.). like galaxy) to −0.183 dex kpc−1 (for a late type galaxy like However, these 19 simulated galaxies analyzed cover a nar- NGC 300). More recent works (Mollá & Díaz 2005) calculate row range of morphologies, mainly Sbc–Sc, in comparison with models in which the infall rate was a function of the mass dis- CALIFA. Therefore, these results cannot be extrapolated to the tribution (or rotation curve) of the galaxy, assuming a stronger full work presented here, but they are representative of the state- radial dependence of the infall timescales than in Chiappini et al. of-art of cosmological simulations of disk galaxies, and can (2001). Moreover Mollá & Díaz(2005) models also depend on be used to compare them to similar disks from our observa- an efficiency factor to condense the molecular gas, and to con- tions. Mock B-band images are used to derive the HLR and vert the gas reservoir into stars. The metallicity gradients range to perform a bulge-disk decomposition used as a proxy for the −0.02 to −0.15 dex kpc−1 with flat gradients for galaxies with the morphology. Metallicities are calculated for disk particles us- largest efficiency factor, or the most massive galaxies, although ing Eq. (2) and the gradient is derived between 0−1 HLR and in the extreme end the low-mass and lowest efficiencies mod- also between 1 and 2 HLR. The stellar metallicity and age gra- els also show flat radial distributions. Thus, the steepest gradi- dients of the simulated galaxies are compatible with the results ents appear in the intermediate-mass or intermediate-type galax- presented here. Keeping in mind that the morphological range ies. However, there is no dependence on the morphological type covered by these simulations is rather narrow and that they when the gradient is normalized to a characteristic value, such use B/D as a proxy for a morphological classification, the re- as the effective radius, as recent results by CALIFA have found sults show a slight dependence of 5outhlog Z?iM with B/D ra- based on HII regions abundances (Sánchez et al. 2013). tio, with a steeper slope for B/D = 1 (Sbc galaxies) for which 5outhlog Z?iM ∼ −0.1 dex/HLR. Later type spirals have a flatter gradient of 5outhlog Z?iM ∼ −0.046 dex/HLR. These results go 6.5.2. Theoretical predictions from cosmological simulations in line with those found here, namely, that the metallicity radial Recently, hydrodynamical cosmological simulations have pro- gradient of spirals shows a dependence on morphology, with the vided evidence in support of the imposed inside-out disk growth steepest gradient found in the intermediate Sb–Sbc spirals (see scenario adopted within the “classical” chemical evolution mod- Figs. 12 and 13). However, a larger set of cosmological simulations is required covering from early Sa to late Sd, and a large els. Like spheroidals, spirals are formed in two phases. In the 9 11 first phase, the bulge formed in a similar way as the core of range of galaxy masses (from 10 to 10 M ), to confirm the E-S0. In the second phase, the disk grows by star formation in general trend found here. On the other hand, these results indi- situ from the infalling gas (Kauffmann et al. 1993; Aumer & cate that the feedback recipes used in these simulations are able White 2013). Metal poor gas with higher angular momentum at to recover realistic galaxies with small bulges and are fully in lower redshifts is turned into stars at larger radii. Negative radial agreement with the work presented here. metallicity gradients are expected, as the classical models pre- Furthermore, our results are also compatible with classical dict. This assumption is a natural outcome of the mass, momen- chemical models, and certainly, the CALIFA Sb–Sbc galaxies tum, and energy conservation laws imposed in the simulations have stellar metallicity gradients (−0.025 dex/kpc) in the range of disks in a cosmological context (Brook et al. 2011, 2012; Few observed in the MW disk, but somewhat shallower than the et al. 2012; Pilkington et al. 2012a,b; Gibson et al. 2013). [Fe/H] gradient measured in the Galactic disk. However, it is necessary to take into account that the gradient usually given in Pilkington et al.(2012a) examined a set of 25 simulations, the literature is obtained for young stars or HII regions, while from several groups, using different codes and initial conditions here it is an average value obtained for all stellar populations ex- (Stinson et al. 2010; Rahimi et al. 2011; Kobayashi & Nakasato isting in the studied galaxy or region. Besides that, the number 2011; Few et al. 2012) to predict the present-day metallicity gra- of objects has increased compared with the old studies and, more dient in MW-like galaxies and its evolution. Although the evo- importantly, all of them have been self-consistently analyzed us- lution of the simulated metallicity gradients depends strongly ing the same reduction technique and spectral models. on the choice of the subgrid physics employed, most of the In any case our results favor an inside-out growth of spirals. simulated galaxies tend to a similar present-day gradient of This conclusion is supported by the stellar age radial profiles ∼−0.05 dex kpc−1, in agreement with the Chiappini et al.(2001) presented here: the age decreases outwards for all Hubble types and Mollá & Díaz(2005) models for normal galaxies as the MW. studied11. Beyond 1.5−2 HLR the radial distribution of ages flat- tens, suggesting that the mass forms more uniformly in those 6.5.3. Implications from this work regions, or that the stellar mixing brings stars born in the inner disk to the outskirts. This last possibility has been recently Our findings show that spiral galaxies (excluding Sd) have neg- investigated by Minchev et al.(2013, 2014), who have per- ative radial gradients as indicative of the inside-out growth of formed N-body hydrodynamical models with the chemical evo- the disk (see Fig. 12). The average 5outhlog Z?iM for spirals (ex- lution implementation (Minchev et al. 2013, 2014). They simu- cluding Sd and later type) is −0.08 dex/HLR or −0.02 dex/kpc. lated MW-like galaxies with the aim to investigate whether the These values are compatible with the results obtained by Galactic disk can be understood as a single structure with kine- Sánchez-Blázquez et al.(2014), which have already derived the matic and chemical features that are continuously distributed, metallicity gradients for 62 CALIFA face-on spiral galaxies to study the effect of bars on the properties of the stellar 11 Sd galaxies, however, show a much flatter age gradient.

A103, page 20 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence being the thin and thick disks two extreme cases of these is preserved at any given distance. At constant M?, log µ?(R) structures. Furthermore, they investigate the effect of stellar mi- is higher in early-type than in late-type spirals. E and gration and kinematic heating in the scatter of the age-metallicity S0 show equal log µ?(R) profiles, independently of M?. relation, and how it changes with the Galactic radius. In fact, an The inner gradient, 5in log µ?, correlates with Hubble type. increase of the scatter in the age-metallicity relation and a flat- The negative gradients steepen from late type spirals to tening of the stellar metallicity gradient is produced by the stellar spheroids, as well as with galaxy total mass in galaxies with 11 radial migration, which causes a radial mixing in the older stel- M? ≤ 10 M . At a constant M?, 5in log µ? steepens with lar population and creating the appearance of a flatter gradient morphology, with E and S0 having the steepest gradients. in early times, leading to a decoupling of the stellar population These results indicate that morphology, and not only M?, from their birth interstellar medium (Roškar et al. 2008). These plays a relevant role in defining µ?, and the µ?−M? relation. results also indicate that even though radial mixing has a significant effect in flattening the metallicity gradient, it can not destroy 4. Stellar ages: hlog ageiL(R) shows declining profiles that the negative gradient. scale with morphology; this behavior is preserved at any given distance. Early-type spirals are always older than late spirals. E and S0, although older than spirals, both have 7. Summary and conclusions similar hlog ageiL(R) profiles, indicating that these galaxies have similar SFH. The more massive are also the older We analyzed the stellar population properties of 300 galaxies, galaxies; this “downsizing” behavior is always preserved at observed by CALIFA with the V500 and V1200 gratings and any given distance. The negative 5inhlog ageiL depends on IFU PPak at the 3.5 m telescope of Calar Alto, to investigate Hubble type in different ways: steeper from E and S0 to the trends in the stellar populations properties with radial dis- Sbc, and shallower from Sbc to Sd. Thus, Milky Way-like tance as a function of Hubble type and galaxy stellar mass. The galaxies have the steepest age gradient. A 5inhlog ageiL−M? sample includes ellipticals, S0, and spirals from early (Sa–Sb) relation exists, increasing the gradient from the low mass to late types (Sc–Sd). They cover a stellar mass range from galaxies (which have roughly flat profiles) up to about × 9 × 11 11 0.7 10 to 7 10 M if Salpeter IMF is assumed, and a 10 M , at this point the trend reverses and 5inhlog ageiL de- factor 1.78 (0.25 dex) lower for a Chabrier IMF. A full spec- creases with increasing M?. However, the dispersion in the HLR tral fitting analysis was performed using the  code 5inhlog ageiL−M? relation and hlog ageiL − M? is sig- and a combination of SSP spectra from González Delgado et al. nificant and it is strongly related to the morphology. Even (2005), Vazdekis et al.(2010), or Charlot & Bruzual (2007, priv. more, the dispersion of the hlog ageiL(R) profiles of galax- comm.). Our pipeline y is used to process the spectral ies of equal mass is significant and larger than between the fitting results to produce present day maps of the spatial dis- hlog ageiL(R) profiles of galaxies of different M? but the tribution of the stellar population properties. For each galaxy, same Hubble type. Thus, the SFHs and their radial varia- these maps are azimuthally averaged to produce radial profiles tions are modulated primarily by the Hubble type, with mass L (in units of the half light radius, HLR: a50) of the stellar mass playing a secondary role. surface density (log µ?), stellar ages (light-weighted, hlog ageiL, and mass-weighted, hlog agei ), metallicity (hlog Z i ), and ex- M ? M 5. Stellar metallicity: hlog Z?iM(R) shows mildly decreasing tinction (AV ). The radial profiles are stacked as a function of profiles for most Hubble types, except Sd that show little, if Hubble type and of galaxy mass. Radial gradients of these prop- any, radial dependence. Milky Way-like galaxies (Sbc) stand erties measured within the inner 1 HLR and between 1 and out as those with the steepest radial profiles. hlog Z?iM(R) 2 HLR are also obtained. scales with M? similar to the way as it scales with the mor- Our main results are: phology. This can be understood as a consequence of the global mass metallicity relation, which is a primary depen- 1. Spatially averaged vs. integrated galaxy properties: the dence of the metallicity with M?. The metallicity gradients h i metallicity, log Z? M, galaxy-wide spatially averaged are negative but shallow on average, with 5inhlog Z?iM ∼ matches the metallicity obtained from the integrated spec- −0.1 dex/HLR, and show a small dependence with M? up 11 trum, and the metallicity at R = 1 HLR. This result is equiv- to M? ∼ 10 M , steepening with increasing mass. Above 11 alent to that obtained for the other stellar population prop- 10 M (Sa’s, S0’s and E’s) they have similar metallicity h i erties, log µ?, log age L, and AV , as reported by González gradient. The dispersion in the 5inhlog Z?iM–M? relation is Delgado et al.(2014c,b), proving that e ffective radii are in- significant and a trend with morphology is seen, in the sense deed effective. that, for a given mass, intermediate-type spirals are those with steeper gradients. 2. Mass-weighted size: we confirm our earlier finding (González Delgado et al. 2014c) that galaxies are more com- 6. Stellar extinction: all the galaxies show AV (R) declining pro- pact in mass than in light by ∼20%. The HMR/HLR ratio files, but do not have a clear trend with morphology or with shows a dual distribution with Hubble type, that breaks the galaxy mass. Most spirals show similar radial variations, and galaxies with the smaller HMR/HLR in the Sb–Sbc. This similar average extinction, AV ∼ 0.2 mag at the disk and up ratio also shows a dual dependence with M?: it decreases to 0.6 mag in the center, with the inner gradient 5inAV ∼ with increasing mass for disk galaxies, and becomes almost −0.25 mag/HLR. However, Sd galaxies show a shallow cen- constant in spheroidal galaxies. These results are a signpost tral gradient. E and S0 also show a negative gradient in the of the inside-out growth previously found by Pérez et al. inner 1 HLR (shallower than in early-type spirals), but they (2013). are almost dust free out of the core. On average, 5inAV gets 11 stronger with increasing M? up to 10 M , and weakens to- 3. Stellar mass surface density: log µ?(R) shows declining pro- ward higher mass. However, the dispersion with respect to files that scale with morphology and with M?; this behavior

A103, page 21 of 44 A&A 581, A103 (2015)

the binned mass relation is not related to Hubble type. A – Through the negative metallicity gradients, spirals show evi- major reason for this is that AV (R) profiles are sensitive to dence of growing inside-out. These gradients are flatter than inclination effects, unlike µ?, hlog ageiL, or hlog Z?iM. Thus, the predictions by the classical chemical evolution mod- spirals with larger inclination have larger extinction. This is els (e.g. Chiappini et al. 2001; Mollá & Díaz 2005), but HLR particularly evident in Sb that have the largest AV and the are similar to those measured above the Galactic disk. The steepest central gradient. largest gradient happens in intermediate types and intermediate galaxy mass, as predicted by the Mollá & Díaz(2005) models. However, Sbc galaxies have a hlog Z?iM gradient From these results, we conclude that: similar to the predictions by RaDES simulations (Few et al. 2012; Pilkington et al. 2012a). This indicates that the feed- – Evidence in favor of the inside-out growth of galaxies is back recipes used in these simulations are able to recover found in the negative radial stellar age gradients. Metallicity realistic galaxies with small bulges. However, a larger set of gradients and the fact that galaxies are more compact in mass cosmological simulations is required, covering from early- than in light also support this scenario. On the other hand, the type Sa to late Sd, and a large range of galaxy mass (from 9 11 flattening of the hlog ageiL(R) profiles beyond 1.5−2 HLR 10 to 10 M ), to confirm the general trend with the Hubble may be interpreted as indicative that the mass was formed in type found in this work. a more uniformly distributed manner across the outer disk of Thanks to the uniqueness of the CALIFA data in terms of spa- spirals. In the case of E and S0 this may be understood if be- tial coverage and resolution, large sample spanning all mor- yond 2 HLR most of the stellar mass was accreted at z = 1. phological types, and homogeneity and quality of the spectral analysis, we are able to characterize the radial structure of the – The mean stellar ages of disks and bulges are correlated with stellar population properties of galaxies in the local universe. disks covering a large range of ages, and late-type spirals The results show that the Hubble sequence is a useful scheme hosting the younger disks. The bulges of S0 and early type to organize galaxies by their spatially resolved stellar density, spirals are old and metal rich as the core of E. They formed ages, and metallicity; however, stellar extinction cannot discrim- with similar processes, through mergers. Later-type spirals, inate between the different Hubble types so well. Stellar mass, however, have younger bulges, and larger contributions from although responsible for setting the average stellar population secular evolution are expected. Disks are younger and more properties in galaxies, is less responsible for the quenching pro- metal poor than bulges, as an indicative of the inside-out for- cesses. Morphology is, however, more strongly connected with mation scenario of these galaxies. the shut down of the star formation activity in the bulges and disks of galaxies. – S0 in this sample (all are massive galaxies) act as a transition class between E and spirals, with µ?(R), hlog ageiL(R), and A (R) between massive E and Sa. The gradient in µ and Acknowledgements. CALIFA is the first legacy survey carried out at Calar Alto. V ? The CALIFA collaboration would like to thank the IAA-CSIC and MPIA-MPG hlog ageiL of S0 is so similar to Sa galaxies that they can as major partners of the observatory, and CAHA itself, for the unique access result from the same formation process. to telescope time and support in manpower and infrastructures. We also thank the CAHA staff for the dedication to this project. Support from the Spanish – The Hubble type rather M drives differences in the galaxy Ministerio de Economía y Competitividad, through projects AYA2010-15081 (PI ? R.G.D.), and Junta de Andalucía FQ1580 (PI R.G.D.), AYA2010-22111-C03-03, averaged ages, and radial age gradients. These results indi- and AYA2010-10904E (S.F.S.). We also thank the Viabilidad, Diseño, Acceso y cate that the SFH and their radial variations are modulated Mejora funding program, ICTS-2009-10, for funding the data acquisition of this primarily by galaxy morphology, and only secondarily by project. R.C.F. thanks the hospitality of the IAA and the support of CAPES and the galaxy mass. This suggests that galaxies are morpholog- CNPq. R.G.D. acknowledges the support of CNPq (Brazil) through Programa ically quenched and that the shutdown of star formation oc- Ciencia sem Fronteiras (401452/2012-3). A.G. acknowledges support from EU FP7/2007-2013 under grant agreement n.267251 (AstroFIt) and from the curs outward and earlier in galaxies with a large spheroid EU Marie Curie Integration Grant “SteMaGE” Nr. PCIG12-GA-2012-326466. than in galaxies of later Hubble type. C.J.W. acknowledges support through the Marie Curie Career Integration Grant 303912. E.P. acknowledges support from the Guillermo Haro program at INAOE. Support for L.G. is provided by the Ministry of Economy, Development, From the comparison of the results with the theoretical predic- and Tourism’s Millennium Science Initiative through grant IC120009, awarded tions from cosmological simulations, we conclude: to The Millennium Institute of Astrophysics, MAS. L.G. acknowledges support by CONICYT through FONDECYT grant 3140566. J.I.P. acknowledges financial support from the Spanish MINECO under grant AYA2010-21887- – Major mergers are likely the main process in building the C04-01 and from Junta de Andalucía Excellence Project PEX2011-FQM7058. central regions of ETGs. The metallicity gradient within I.M., J.M. and A.d.O. acknowledge support from the project AYA2013-42227-P. 1 HLR is shallow compared with the theoretical expecta- RAM is funded by the Spanish program of International Campus of Excellence Moncloa (CEI). J.M.A. acknowledges support from the European Research tion of whether minor mergers are relevant in the growth Council Starting Grant (SEDmorph; P.I. V. Wild) of the central core of E’s and S0’s. In our results there is no evidence either of an inversion of hlog ageiL toward older ages beyond 1−2 HLR, or of a steepening of the metallic- References ity if these galaxies were growing in size through minor dry Alongi, M., Bertelli, G., Bressan, A., et al. 1993, A&AS, 97, 851 mergers. Massive galaxies probably accreted massive satel- Annibali, F., Bressan, A., Rampazzo, R., Zeilinger, W. W., & Danese, L. 2007, lites that were able to retain their metal rich gas against A&A, 463, 455 winds, producing flatter metallicity gradients (Hirschmann Arribas, S., Colina, L., Monreal-Ibero, A., et al. 2010, in Galaxy Wars: Stellar et al. 2015). Alternatively, the flattening of the metallicity Populations and Star Formation in Interacting Galaxies, eds. B. Smith, J. Higdon, S. Higdon, & N. Bastian, ASP Conf. Ser., 423, 317 radial profile can result from the quenching of star forma- Asari, N. V., Cid Fernandes, R., Stasinska,´ G., et al. 2007, MNRAS, 381, 263 tion. When this happens, the metal cycle stops and only stars Aumer, M., & White, S. D. M. 2013, MNRAS, 428, 1055 of that last star formation event remain. Aumer, M., White, S. D. M., & Naab, T. 2014, MNRAS, 441, 3679

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Appendix A: Improvements resulting from the new to figure 13 of Cid Fernandes et al.(2014). The broad trough reduction pipeline around Hβ present in the 1.3c spectra, for instance, is much shallower now. In fact, it is confined to late types (compare blue The quality of the  fits to CALIFA version 1.3c spec- and red lines in the lower panel in Fig. A.1), indicating that tra was assessed in Cid Fernandes et al.(2014) by averaging its origin is related to calibration, and also to the SSP spec- 5 Rλ = Oλ − Mλ residual spectra of 107 galaxies (∼10 spectra). tra of young stellar populations (as previously reported by Cid Inspection of these residuals revealed low amplitude (a few %) Fernandes et al. 2005 for SDSS data). Residuals are also visi- but systematic features related to unmasked weak emission lines, bly smaller towards the blue, including the CaII K line, which is SSP deficiencies, and data calibration imperfections. This exer- now well fitted whereas in version 1.3c a small systematic resid- cise needs updating now that the reduction pipeline has changed ual subsisted12. The humps around 5800 Å, on the other hand, to version 1.5. are still present in version 1.5, particularly noticeable for outer Figure A.1 summarizes the results of this re-evaluation. The regions, indicating that further refinement of the sky subtraction plots show stacked Rλ = Oλ − Mλ residual spectra, in units of are warranted. the median flux in the 5635±45 Å window. The top panel shows In short, the spectral fits have improved substantially with results for the nuclear extractions, while the middle and bottom the new reduction pipeline. We attribute this to the updated sen- panels are built using spectra from zones within radial distances sitivity curve used in version 1.5. A more extended discussion R = 0−1 and 1−2 half light radius (HLR, computed in the same of these and other aspects of the data reduction are presented in wavelength range), respectively. Residuals are colored accord- García-Benito et al.(2015). ing with the Hubble type of the galaxies. When all galaxies are Despite these changes, the stellar population properties de- stacked, the residuals are colored according with the spatial zone rived from the spectral fits did not change much in compari- extracted. These subdivisions are presented to get a sense of how son to those obtained for version 1.3c data. The most noticeable the residuals relate to position within a galaxy and its Hubble changes were in mean ages, which become 0.1 dex older, and type, which are two central aspects of this paper. extinction, which is now 0.2 mag smaller on a global average. No matter which panel one looks at, the improvement with respect to version 1.3c is evident to the eye when compared

Fig. A.1. Upper panel: residual spectra averaged for all the spectra (black line), nuclei (grey line) and spectra belonging to zones that are between 0−1 HLR (pink line) and between 1−2 HLR (yellow line). Middle panel: residual spectra averaged for zones inner to 1 HLR and for Hubble type. Bottom panel: as in the middle panel except for spectra of zones located between 1 and 2 HLR.

12 Because of this improvement, our  ﬁts now start at 3700 Å, whereas in previous articles in this series only the λ > 3800 Å range was considered.

A103, page 25 of 44 A&A 581, A103 (2015)

Appendix B: Base experiments and uncertainties due to SSP templates To derive the stellar population properties of these 300 CALIFA galaxies we have ﬁtted ∼253000 spectra with the GMe and CBe using the cluster Grid-CSIC at the Instituto de Astrofísica de Andalucía and the cluster Alphacrucis at IAG-USP Sao Paulo. Examples of the quality of the spectral ﬁts as a function of the Hubble type and radial distance are presented in Fig. A.1.

B.1. Mass-metallicity relation Here we want to find out how well our metallicity definitions follow a MZR which guarantees that galaxies like the MW or Andromeda (log M?(M ) ∼ 11) have solar metallicity at the disk, while LMC-SMC-like galaxies (log M?(M ) ∼ 9) have ∼1/4 Z . We do this with mass-weighted and luminosity- weighted definitions of Eqs. (2) and (3), and with the two sets of SSP models (GMe or CBe). Similarly, the correlation between the galaxy averaged stellar metallicities and the metallicity measured at 1 HLR, and the MZR, guarantee that the metallicity radial profiles scale with the galaxy stellar mass. However, the MZR in González Delgado Fig. B.1. The global stellar MZR for 300 CALIFA galaxies is shown as et al.(2014a) was derived using the galaxy averaged stellar dots, color coded by the morphological type. The metallicity is derived HLR HLR metallicity instead of the metallicity measured at 1 HLR, and us- using hlog Z?iM (upper panels) and hlog Z?iL (lower panels), and ing the mass-weighted definition of the metallicity and only the the total stellar mass, M?, obtained with the GMe SSP models (left results with the base GMe. For these reasons, we derive the MZR panels) and CBe (right panels). A mass-binned smooth mean relation that results from using the mass-weighted and the light-weighted is shown as a solid black or grey-black line. The MZRs obtained for definition of the metallicity, and the GMe and CBe bases. SDSS galaxies by Gallazzi et al.(2005) and Panter et al.(2008) are plotted as brown and magenta lines, respectively, with dashed brown h iHLR Figure B.1 shows the correlation of M? and log Z? M indicating the 16 and 84 percentiles of Gallazzi et al.(2005). HLR (upper panels), and hlog Z?iL (lower panels) for the GMe (left panels) and CBe (right panels) SSP models. The mass- metallicity relation found by Panter et al.(2008) and Gallazzi while CBe goes up to 2.5 Z . Because MILES is built with stars et al.(2005) are the magenta and brown lines, respectively 13. in the solar neighborhood, it does not contain stars as metal rich Note that in the four cases, the metallicities are well in the range as 2.5 Z , so CBe results in over solar metallicity should be inter- of the dispersion given by Gallazzi et al.(2005); brown dashed preted with care. On the other hand, the central parts of galaxies line represent the 16th and 84th percentiles of their distribu- can be as metal rich as 2−3 Z , and the base GMe may be too tion. To compare the general trend of these values, we derived low to fit spectra of these regions, leading to saturation effects. a smoothed mass-binned relation, represented by a solid black To avoid these problems, we also fitted the spectra with another or grey-black line. As expected from the global MZR derived two sets of SSP, are identical to GMe and CBe, but where for in González Delgado et al.(2014a), base GMe and hlog Z?iM each metallicity bin, two SSPs of age 16 and 18 Gyr are added predict stellar metallicities for MW and LMC-SMC with the ex- to our “standard” bases. These extra bases, which we name GMd pected values. But hlog Z?iL gives a MZR that predict higher and CBd, allow galaxies older than the age of the universe if metallicities. The opposite happens for the MZR using the base their bulges are very metal rich. Furthermore, results at very CBe, which gives mass-weighted metallicities higher on average low metallicity also must be taken with care. MILES contain than the SDSS metallicities. But the MZR with hlog Z?iL goes only few stars of metallicity below 1/100 Z . For this reason, close to the Gallazzi et al.(2005) relation, and also predicts stel- Vazdekis et al.(2010) provide a safe age range for each metal- lar metallicities for MW-Andromeda-like and LMC-SMC galax- licity bin, being the models with log Z?(Z ) ≤ −1.7 only valid ies with the expected values. between 10 and 18 Gyr. This safety margin is provided to avoid In summary, hlog Z?iM with GMe and hlog Z?iL with CBe the cases when, because of age-metallicity degeneracy, these old provide a mass-metallicity relation similar to the SDSS MZR, metal poor models fit young metal-rich populations, may hap- and predict metallicities between −0.7 and −0.4 dex for galaxies pen if the base does not include SSP younger than 100 Myr. Our 9 10 with mass between ∼10 and 10 M , the expected values for fits do not suffer this problem because bases GMe and CBe both LMC and SMC-like galaxies, and solar for MW-like galaxies. have spectra of ages as young of 1 Myr.

B.2. Variations over bases GMe and CBe B.3. Global results and uncertainties associated Although the two sets of models are built with the same stel- with SSP models lar libraries (MILES and ), base GMe stops at 1.5 Z , To evaluate to which extent the spectral synthesis results de- 13 The Gallazzi et al.(2005) and Panter et al.(2008) relations have been pend on the choice of SSP models, we now compare the global shifted by 0.25 dex to the right to account for the diﬀerence in IMF properties derived with bases GMe and CBe. Using our pipeline between their results based on models with Chabrier IMF and ours ob- y we obtained the radial distribution of the stellar pop- tained with Salpeter IMF in the two left panels. ulation properties for each galaxy with a spatial sampling of

A103, page 26 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence

Fig. B.2. Correlation between SP obtained with the GMe (x-axis) and CBe (y-axis bottom and middle panels; and GMe (x-axis) and GMd (y-axis). The average diﬀerence between the property in the y-axis and x-axis is labeled in each panel as ∆, and the dispersion as (σ).

0.1 HLR. Here we compare the stellar population properties are always higher with CBe than GMe, reflecting the saturation of the 0.1 HLR radially sampled points, instead of comparing effects in the base GMe due its limitation to Z ≤ 1.5 Z . The shift the results obtained from the individual 253 418 fitted spectra. at under-solar metallicities may be reflecting the age-metallicity Figure B.2 shows the results for a total of 6000 points corre- degeneracy, CBe giving higher metallicity and younger ages. sponding to a maximum of 20 radial points (from nucleus to The upper panels compare the results of GMe with the GMd. 2 HLR) for each of the 300 galaxies analyzed in this work. The Here, we see two relevant effects. The results of GMd also differ figure compares the results for base GMe in the x-axis, with CBe from GMe in the range of extinctions allowed to . in the y-axis, in the bottom and middle panels. The upper panels While with GMe and CBe  always assumes AV ≥ compare the results of GMe with GMd. Each panel quotes the 0, with GMd,  can bluer the SSP spectra by up to mean ∆ and its standard deviation, where ∆ = property(CBe) − AV = −0.5 mag. This is allowed to avoid the effect of saturation property(GMe) or ∆ = property(GMd) − property(GMe). at AV = 0. The global effect is that ages can be 0.11 dex younger GMe GMd GMe-based µ?-values are higher than CBe by 0.27 dex on with than . Metallicity is not affected by this choice of average, reflecting the different IMF used. Apart from this offset, AV ; but the extension to ages older than the age of the Universe in the two stellar mass surface density agree to within 0.08 dex. GMd has some effect on the metallicity above Z , so metallicities Mean extinction is also in good agreement with a dispersion of are slightly lower and ages slightly older. 0.05 mag. Ages are higher in GMe than CBe by 0.14 dex for hlog ageiL and 0.08 dex for hlog ageiM, with dispersion 0.18 dex and 0.12 dex, respectively. This result is expected since base Appendix C: 2D maps and tables GMe also differs from CBe in IMF and isochrones. The dif- C.1. 2D maps of logµ ferences in opacities in the equation of state between Padova ? 2000 (GMe) and 1994 (CBe) tracks produce somewhat warmer Figure C.1 shows the Mr vs. u − r CMD for the stars in the red giant branch in the former. Thus, older ages are 300 CALIFA galaxies of our sample. Each galaxy is represented expected with GMe than with CBe. However, the metallicities by its 2D map of the log µ? located at the position of its inte- are lower in GMe than in CBe by 0.13 dex for hlog (Z?/Z )iM grated Mr and u − r values. In this plot, the SFH is compressed and very similar (on average) for hlog (Z/Z )iL. In both cases, into the present-day stellar mass surface density, which measures the dispersion is similar, 0.11 and 0.13 dex, respectively. Note the end product of the SFH. Because our analysis accounts for that for Z? ≥ Z , the metallicities (weighted in light or in mass) extinction, these log µ? values and their radial variations are free

A103, page 27 of 44 A&A 581, A103 (2015)

Fig. C.1. 2D maps of stellar mass surface density, µ?. Each galaxy is placed in its location in the u − r vs. Mr diagram, where color and magnitude correspond to its global values. The 2D maps are shown with north up and east to the left.

14 from extinction eﬀects. Figure C.1 clearly shows that log µ? relation. Gradients of the stellar metallicities are more clearly correlates with Mr, and spheroids are signiﬁcantly denser than seen in these 2D maps in galaxies with intermediate luminos- late-type galaxies by one to two orders of magnitude at the cen- ity (−22 ≤ Mr ≤ −20). The most luminous galaxies have so- ter, and by one order of magnitude at distances 1−2 HLR. At the lar or over solar metallicity producing a visual saturation in the −2 −2 center, 2.0 ≤ log µ? (M pc ) ≤ 4.7, while ≤ log µ? (M pc ) 2D maps. The stellar metallicities range from hlog Z?iM = –1.4 −2 ≤ 3.4 at 1 HLR, and 1.0 ≤ log µ? (M pc ) ≤ 2.9 at 2 HLR. to 0.22.

C.2. 2D maps of hlog ageiL C.4. 2D maps of AV

Similarly to Figs. C.1, C.2 shows the 2D maps of hlog ageiL. Similarly to Figs. C.1, C.4 presents 2D maps of AV .Effects of It portrays the correlation between the average age of the stel- spatial binning are visible in the AV maps, where all the pixels lar populations and Mr and colors, with the most luminous within a Voronoi zone have the same value. These effects are not and red galaxies being older, while the bluest galaxies are the noticeable in the µ? images (Fig. C.1) because µ? is an extensive youngest. Gradients of the stellar population ages are also clearly property, and the zoning effect was softened by scaling the value detected within each galaxy in these 2D maps, and more re- at each pixel by its fractional contribution to the total flux in the markably in galaxies located in the green valley. At the center, zone; this is not possible for AV (Fig. C.4), hlog ageiL (Fig. C.2), 7.3 ≤ hlog ageiL (yr) ≤ 10.1, while 8.3 ≤ hlog ageiL (yr) ≤ 10.1 or hlog Z?iM (Fig. C.3), because these are intensive properties. at 1 HLR, and 7.5 ≤ hlog ageiL ≤ 9.9 at 2 HLR. Figure C.4 shows how AV changes across the CMD: the most luminous galaxies are little affected by extinction, while AV is higher in spirals of intermediate type and with blue colors. The C.3. 2D maps of hlog Z?iM mean (dispersion) AV values at the nuclei, 1 HLR, and 2 HLR are Similar to Figs. C.1, C.3 presents the 2D maps of hlog Z?iM. 0.47 (0.37), 0.19 (0.16), and 0.13 (0.13), respectively. 2D maps It clearly shows the correlation between the stellar metallic- and the difference in the mean values at different distance indi- ity and galaxy luminosity, equivalent to the mass metallicity cate that stellar extinction shows radial gradients.

14 These 2D maps are the results from the GMe SSP base.

A103, page 28 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence

Fig. C.2. As Fig. C.1 except for images of the luminosity-weighted mean age, hlog ageiL.

Fig. C.3. As Fig. C.1 except for images of the mass-weighted mean metallicity, hlog Z?iM.

A103, page 29 of 44 A&A 581, A103 (2015)

Fig. C.4. As Fig. C.1 except for images of the stellar extinction, AV [mag].

A103, page 30 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence C HLR HMR M i ] pc pc

Z [HLR] Z log h M i ][ [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) [ A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M . CBe Stellar population properties: CALIFA001 Name002003004005 Type IC 5376 UGC007 00005 log 008 NGC 7819 UGC010 00029 3012 A 2 A IC 1528 UGC013 00036 4 A –1 10.84 A014 10.60 NGC 0001016 10.04 NGC 11.16 0036 3.29 1 AB 3.87 UGC017 3 MCG-02-02-030 00148 AB 2.80 3.67019 10.91 3 10.09 A 2 AB UGC020 MCG-02-02-040 00312 2 2.23 UGC B 00335NED 3.89 4 2.41 02023 A 10.68 10.29 3.03 4 AB 1.59 –1 2.49024 A 11.00 ESO 540-G 10.11 003 0.63 5 3.85 3.51025 B 0.42 9.85 10.58 2.70 3.71026 NGC 0.63 1.95 0.44 2 0160 2.97 AB 0.41 9.80027 0.53 NGC 4.06 2.88 0171 2.53 2.39028 0.05 9.95 0.03 NGC 0.73 0177 0.80 2.35 2.93029 NGC 1 0180 2.20 A 9.51 0.19030 3.21 9.97 NGC 2.33 0.43 0.50 1.89 2 0192 0.35 B031 10.87 8.46 NGC 9.89 0.33 1 0216 A 0.99 1.51032 0.63 0.09 10.63 NGC 2 0214 B 0.46 1.24 4.06033 1 9.66 2.21 0.34 10.27 NGC AB 0217 8.97 0.51 3.80035 8.88 10.80 NGC 0.42 5 0237 0.04 9.38 0.24 10.70 A 3.92036 3 MCG-02-03-015 9.55 9.67 NGC AB 0234 0.27 8.86 9.75 2.47 3.46038 0.29 9.97 4.14 1 8.90 10.84 A 1 8.86 AB 2.26039 NGC 10.14 0.16 4 0257 9.51 B 8.63 8.64 1.91 –0.03040 4 10.80 10.69 NGC AB 0.38 3.84 2.53 0364 0.25 2.14041 10.15 NGC 8.59 2.58 0.38 0429 8.95 9.26 10.56 0.13 0.37 4.06 3.78042 0.26 NGC 0.96 9.70 4 0447 9.31 A –1 3.23043 AB 8.33 0.00 NGC 2.45 3.20 –0.04 0.20 1.86 9.78 0444 0.39 1.21 –0.22 8.75 UGC044 0.07 10.79 10.81 00809 –0.17 1 A 2.63 2.43 10.03 UGC045 0.08 00841 –0.27 1 –0.06 0.32 8.17 0.14 B 0.07 0.71 0.29 4.22 6842 3.67 2.05 10.02046 10.71 NGC 2.09 4 0477 A 0.05 8.62 3883 5146 4 –0.62047 A 9.79 –0.29 11.01 0.38 0.70 0.32 0.35 0.20 0.12 2.1 IC 2076 5455 –0.15 3 1683 4.00049 A NGC 9.35 5400 0.09 9.72 0499 3.3 9.19 0.79 4057 –0.02 2.36 9.56 2.31 4.03050 3 9.85 0.05 0.16 NGC –0.06 AB 4372 0496 –1.33 2.1 4348 9.98 3.1 051 9.26 0.42 NGC 4104 –0.76 2.48 0504 0.42 0.09 10.52 9.62 –0.19 8.65 3417 2.89 2.51052 2 –1 –0.10 3142 NGC –0.19 0.59 AB 0.61 3517 A 0517 2.1 9.12 2781 2.7 10.13 2.63 2.39 UGC053 9.65 00987 2782 3.15 4 10.49 9.15 A 11.24 1790 –0.52 4820 9.39 2.2 0.00059 0.39 0.00 NGC 9.43 5717 0.00 0 1.51 0528 3.0 8.56 A 3298 3264 4341 1.61061 10.45 NGC 0.14 0.36 4152 3.42 0.15 0 0529 4.42 2.2 A 1 2179 2435 AB 8.79 1.76 8.54 2.5 062 0.01 10.48 NGC 1.89 3.2 –0.14 5055 0551 2.2 0.38 10.08 2.75 UGC068 9.60 10.71 0.03 10.64 9.41 01057 0.82 3659 0 9.43 A 0.21 –0.37 4.03 UGC 2362 0.12 2.1 –1 01271 0.76 2.38 A 3.08 0.17 8.66 0.46 4.14 8.53 4.07 2035 3 0.18 0.23 10.95 9.97 NGC AB 0681 4 2.6 11.11 –0.01 AB –0.40 1.86 UGC 0.12 0.29 01274 9.90 4165 0.15 10.59 1.12 0.09 2.71 0 4.17 9.68 10.28 B 0.11 NGC 2970 0741 4.37 8.73 0.12 9.12 0.07 –0.08 4523 2.73 2.72 3.0 1 8.75 AB 0.69 3.25 3268 –0.05 0.43 0.03 10.80 –0.09 2979 3.11 1 A 9.04 –0.25 1481 0.01 9.76 2.0 10.11 9.43 0.04 2.87 3.8 0.10 –1 3916 0.34 6498 0.00 A 2.91 3.98 9.47 11.11 0.22 0.00 0.15 2588 4978 2.19 3.89 –0.23 0.01 10.03 8.77 11.48 1254 –0.19 1.94 2.6 8.55 2.2 0.24 0.00 0.00 0.01 4.08 8.57 4346 1351 0.00 8.77 2.5 4.24 8.79 2.65 3270 0.31 0.12 3868 8.91 10.08 0.00 0.44 4431 2.59 2.4 3457 0.00 0.42 3029 –0.32 10.03 10.07 0.19 4954 2616 2558 3.04 3.0 0.14 9.94 –0.34 0.00 9.22 3.1 0.24 2.4 4468 –0.68 2.96 0.61 1.9 9.98 6099 –0.34 0.13 10.04 8.70 0.01 3268 0.36 0.16 9.74 4573 0.24 0.28 9.82 2256 0.14 9.57 2.1 9.74 9.50 –0.38 3.1 0.19 0.38 –0.12 –0.06 2564 10.00 0.00 –0.10 1844 4856 –0.14 9.80 9.75 9.80 –0.09 5690 3.2 0.32 3240 4220 10.00 4555 9.09 3.3 0.09 0.27 3358 4535 0.14 8.69 2.3 10.10 0.31 2.3 6329 4077 9.69 –0.30 2.5 4558 0.35 0.37 2.2 9.53 0.18 3677 3210 9.78 0.19 0.14 0.05 3098 2620 6329 0.28 9.90 2.6 3.0 5510 0.37 1838 2.0 0.18 0.15 2922 1385 –0.05 2323 3.2 2005 0.29 1598 0.14 –0.15 3.0 3.5 0.36 2751 2859 0.25 –0.17 2020 2365 5228 4918 3.1 3.1 0.28 4061 3138 2.1 3013 2.4 0.21 1663 2408 1289 3.1 3791 2.4 2698 4620 3.1 4066 2.8 ID# NED morph. Table C.1.

A103, page 31 of 44 A&A 581, A103 (2015) C HLR HMR M i ] pc pc

Z [HLR] Z log h M i ][ [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) [ A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA069 Name070071072073 NGC Type 0755076 log IC 1755077 NGC 0768081 NGC 4 0774 B087 NGC 0776088 NGC 2 4 9.17 0810 A B100 NGC 0 0825 A UGC101 10.98 10.58 01659 2 2.56 B103 –1 10.77 NGC A 0932 3.87 UGC 3.34104 10.82 01938 1 A 11.39 4 4.00108 B NGC 1.83 1056 3.79115 10.37 NGC 0 1060 3.96 A 3 10.71 AB 2.58 1.93 UGC116 02222 0.37 2.62 4.10 UGC118 10.39 10.97 02229 1 3.07 A 2.05119 –1 NGC 0.33 0.10 A 1093 2.67 0.21 0 AB 3.34 4.13 UGC127 02403 0.35 9.94 11.61 0 2.77 UGC131 0.07 10.71 A 0.17 02405 0.56 1.85 0.58 UGC132 0.02 02465 3 4.01 4.46 B 8.86 10.96 2.04 4.30 2.72 2133 0.17 B NGC 0.00 1167 0.02 0.44 3 10.08134 A 10.61 NGC 9.68 1349 3.72 10.42 0 10.06135 0.25 A NGC 0.54 0.13 2.88 1542 3.20 0.23 10.44 2.74 UGC 3.58141 03038 9.39 0 3.80 A 10.85 9.91 UGC 8.49 0.19143 0.08 –1 03107 A 2.39 2.98 0.96 0.15 10.10144 1 9.75 11.20 NGC 0.24 AB 1645 4.20 8.89 11.12 9.80 1 2.24 UGC146 9.70 A 03151 2.26 0.60 10.36 0.00 0.11 2 4.24 0.01147 9.50 A 9.96 1.68 3.99 9.12 10.70 IC –0.05 2095148 9.80 NGC 3.55 0.14 0 1677 2.71 B 1 0.10 10.73 AB 2.23149 9.66 0.31 10.04 3.50 8.83 0.25 IC 0.35 10.12 2.64 2101 UGC150 10.75 8.99 0.06 0.30 10.84 03253 2.36 3.48 0.50 0.55151 4 4 NGC AB AB –0.08 2.50 2253 8.75 10.08 9.77 0.51 0.38 3.74 UGC 3.98152 0.41 03539 0.17 2.04 0.10 8.60 9.39 0.21 2 0.14153 4 B 0.30 NGC 9.88 AB –0.09 2347 2.32 9.88 0.27 9.03 1638 9.23 UGC 9.77 0.20155 0.00 03899 0.14 3 10.15 B 4 0.01 9.27 10.52 1.61 2.88 1340 AB 0.80 2.61 2.70156 NGC 2410 10.08 0.32 0.11 2.2 0.35 9.65 0.30 5650 6604 0.87 UGC 0.23157 3 10.55 03944 AB 9.79 3.04 3.56 0.39 4152 –0.03 3796 3087 5 10.08 UGC158 A 03969 0.56 0.18 3.0 2.5 10.85 9.28 0.31 9.95 0.39 2265 1.06 1.91 –0.55 UGC 3.65159 2 0.04 9.05 0.32 03995 AB 4943 3.06 3.0 9.66 3 5841 AB 9.35 0.03161 8.00 0.05 –0.23 NGC 2948 1.84 3.91 1806 2449 0.22 2.22 10.89 9.58 9.53 0.21 4624 2 UGC 2.1 165 A 0.00 0.59 04029 1203 9.92 3.0 9.42 8576 2 2.63 2.29 UGC 3.2 171 9.61 B 04054 –0.28 1.97 3.91 0.23 0.25 10.69 0.98 5081 6249 0.40 0.35 0.21 0.12 10.03 –0.07186 1 9.30 9.71 3033 AB 2.56 2.87 2.2 4 IC 3097 11.07 AB 9.34 0480219 0.02 NGC 2341 0.76 2480 2.5 3.29 0.37 10.96 0.14 0.10 1.28 0.39 4 1.68 3.2 UGC 958226 8.87 10.28 A 04132 2.56 4064 4.11 0.29 2826 8.01 8.59 8.90 0.36 0.22 NGC 3624 1.71 3.91 2513 898 0.49 0.15 0.04 2066 3.37 9.33 3.0 4 9.58 0.27 AB 0.43 5458 5 3.2 3.1 9.30 2.33 0.35 9.18 A 3 IC 0.59 9.75 –0.03 AB 2247 0.01 NGC 3813 2639 2.51 10.07 0.25 0.37 2.98 UGC 3.2 0.34 10.83 –1 0.23 04659 9.05 4252 2.75 A 9.26 3747 0.33 9.20 0.88 2.03 0.17 0.13 2838 2.55 2232 8.06 8.63 0.07 9.84 11.22 7902 3.71 0.28 2.6 0.29 2.0 –0.14 0.00 3571 1 1 2.37 0.07 A A 0.58 0.38 4649 5 8.88 1.77 9.06 A 8.61 0.88 2453 2.2 4.23 9.78 0.20 4277 0.04 2.8 10.44 11.03 1.87 0.31 9.28 8.69 2816 4083 –0.25 5199 2.35 8.73 8.83 0.43 –1.04 0.60 1.50 0.40 2.8 2251 9.68 3502 8.94 3.60 4.19 0.16 –0.03 2.5 2.81 2.6 1.19 9.99 –0.17 6289 2.15 –0.04 0.14 1.38 7058 8.15 9.96 4278 0.35 9.31 9.51 4810 –0.14 2.4 0.40 4868 –0.37 3379 0.01 0.05 8.53 2.25 3.07 2.5 0.56 –0.46 0.15 3670 2361 1.28 0.17 9.08 9.50 –0.24 –0.03 2.4 3.2 0.00 9.47 1.12 0.02 –1.16 8.69 2145 1736 –0.13 9.41 0.01 9.22 –0.18 0.48 8.92 1934 1976 7.99 4783 –0.30 3797 0.44 0.20 2.8 2.5 10.01 –0.11 3546 2251 –0.30 0.23 3523 2.5 8.72 2.3 3045 –0.48 0.35 2448 –0.06 2117 8.86 0.23 2.1 9.48 9.95 –0.13 3939 2.8 8.74 0.05 8.60 2454 8.17 –0.08 2253 9.70 2.6 4192 1760 –0.47 3851 2.9 0.21 2945 2534 –0.33 0.14 3.1 6325 –0.24 –0.21 2.2 9.04 9.48 –0.26 5196 8.49 2.4 0.32 –0.26 5286 4073 4598 3309 0.21 2.4 3150 3177 –0.13 –0.14 2.7 2641 –0.01 2.7 0.25 1835 –0.93 0.16 2.7 5759 5756 2662 4705 3925 2397 –0.26 2.3 2.3 3929 0.21 2.6 –0.20 3134 3.1 4250 2830 2935 2189 3.0 2255 1613 2.7 2.2 ID# NED morph. Table C.1.

A103, page 32 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence C HLR HMR M i ] pc pc

Z [HLR] Z log h M i ][ [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) [ A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA231 Name232272273 UGC274 04722 Type275 NGC log 2730277 NGC 2880 5278 A IC 2487279 4 B IC –1 8.56 0540306 AB NGC 2906307 10.44 10.01 NGC 2916 1.53 UGC309 4 05108 AB 4.55 2.67311 1 NGC AB 3 2918 10.43 A UGC312 05358 3 A 9.88 0.99 2 UGC314 10.44 B 05359 3.16 2.99 1.52 UGC318 –1 10.67 05396 A 10.99 3.34 5 3.88319 B NGC 0.32 3106 UGC 11.08 05498NED 2 01 3.58326 B NGC 0.25 2.18 0.53 3057 3.81 3 AB 9.22339 0.14 1 A 3.87 10.87 2.52 2.45 0.00341 10.36 0.23 NGC 1 3158 A 0.53 2.29 10.71 1.83353 NGC 5 3160 2.46 3.30 B 2.98 UGC361 11.23 0.74 7.79 05598 0.58 2.48 0.34 3.49 10.11364 –1 NGC 9.00 0.09 A 8.92 3300 1.04 4.06 0.37 1.28 UGC381 1 0.23 05771 AB 2.12 11.61 0.00 1.97 2386 0.04 A NGC 3381 2.09 10.82 9.59 0.18 2.53 UGC387 05990 0.36 0 4.12 7.91 B 0.00 –1 10.21 0.12 A 2.63 UGC 9.59 9.86388 9.79 06036 0.47 3.46 8.48414 0.53 0.20 9.66 10.61 11.19 5 1.14 3.00 B 0.18 IC 9.55 4 0674 UGC 0.43436 A 06312 0.17 2.88 9.92 0.26 1 9.02 3.96 3.89437 A NGC 2.44 –0.87 9.47 3615 0.25 9.17475 0.02 0.22 NGC 7.97 9.53 3614 –0.17 1.85 9.10 11.09 9.93 0.00 1476 9.50 A NGC 1 3687 2.54 9.07 0.08 B 0.36 2.89 9.79 2.54 2.62479 9.24 –1 NGC A 3811 3.85 0.00 10.88 –0.30 –0.46 1.13 10.02486 3 9.40 11.06 NGC AB 3815 0.01 11.34 –0.03 –0.32489 NGC –0.30 0.19 2 0.01 1.69 3991 3.69 10.13 8.63 B 0.27 0.61 8.49 2.26 3.92500 9.12 NGC 2355 3 3994 8.95 2.88 4.26 B 0.13 9.99502 0.00 0.08 –0.22 10.30 NGC 0.32 1192 1920 9.60 3.15 3 5584 4003 9.92 A 0.28 UGC 2.1 515 9.83 0.17 10.41 07012 0.59 4705 5 2.48 885 A –0.18 8.98 0.06 2.54 3.74518 3 1.9 10.33 NGC –0.01 AB 4047 3.3 2.76 0.14 0.25 6137 UGC 10.00 3.62 8.38 10.07528 0.51 07145 0.13 1.93 9.54 0 0.10 10.41 0.09 B 4 0.09 4994 AB 0.56 3.43548 9.86 0.20 –0.12 NGC 2154 0.17 4149 2.5 2932 0.00 9.76 2.23 0.38569 1812 0.10 11.03 NGC 9.36 4.25 3 2.38 4185 A 0.33 7.99 0.32 2185 2.7 3 2.19580 8.54 A 0.09 8.98 NGC 5141 –0.61 4210 2.1 0.00 9.99 –0.32 4860 –0.46 2.00 3.78588 1 9.54 10.65 9.87 –0.12 3858 AB 2.49 0.06 0.13 –0.14 10.43 IC 3190 4547 0776 2.2 589 3 0.37 NGC AB 2.51 0.56 1.50 4470 2.8 10.26 9.87 2745 –0.02 4017 3.53 10.09 0.14592 0.04 NGC 0.40 2 8278 4644 3.0 10.07 2.88 10.70 B 6220 3462 –0.12 2.37 8.59593 0.18 4434 NGC 1.42 3.68 5819 4711 8.61 0.02 2.4 9.76 5197 0.92 0.12 9.90 2.4 0.24 3654 10.23 NGC 0.05 0.36 3.28 4898 5 4 4816 2.1 A A 3.1 NGC 2.47 10.01 0.18 4841A 4233 0.26 2 1.83 0.51 0.14 A 0.37 0.11 3.42 3.1 9.57 9.68 NGC 9.20 2209 2.49 –0.17 3 9.12 9.88 4874 A 9.72 UGC 0.14 –1 0.01 10.47 9.74 08107 2019 0.37 2.13 A –1.19 –1 0.23 A 5293 0.21 9.12 1.9 0.17 0.14 10.39 0.29 1.79 2.76 9.54 7.52 5313 11.25 1.22 5626 4990 3.50 2.10 11.22 0.40 9.04 0 3.1 A 4260 0.11 0.05 4080 1 3.19 –0.02 A 9.73 0.21 2.7 3.93 0.31 8.97 8.87 2.5 0.37 –0.09 –0.78 4.04 2939 4503 11.49 0.14 1.01 2.01 0.03 8.71 11.09 –0.25 2356 3136 2.38 0.26 9.28 2.7 3.4 9.81 8.43 1.99 3.73 8.06 0.08 0.01 2402 –0.15 1264 2.42 3.46 9.35 0.01 0.22 2.49 0.00 2180 1226 0.05 10.02 –0.01 9.62 0.14 4847 8.03 2.0 2.8 –0.14 0.63 –0.19 0.25 0.31 4367 0.07 4454 2.56 10.01 0.25 0.02 2.18 3.4 8.50 –0.02 4925 3540 0.01 8.97 0.10 0.55 4131 3.0 0.00 3456 3284 –0.14 9.32 0.00 3345 0.00 9.38 2.9 8.58 8.80 0.38 2735 –0.19 0.66 0.03 3.3 9.87 2809 2.2 0.00 8.83 9.48 1905 3827 10.04 0.14 0.45 –0.28 10.00 0.01 –0.27 2.2 2387 –0.26 3764 2.1 0.01 0.30 0.00 2309 7.96 8.51 9.97 2091 2.6 9.15 3613 9.63 –0.10 1117 –0.06 0.08 8.63 9.77 4799 2.7 9.65 –0.08 2864 5263 1.8 2158 3955 0.08 –0.48 –0.04 2.7 7080 2.8 3434 9.79 0.08 0.14 3268 5197 8.99 2943 –0.02 2.3 2397 0.32 2.3 0.36 5782 2.6 5099 –0.06 –0.19 3642 2.0 2795 0.01 0.29 –0.10 1.9 0.11 0.20 3582 2797 0.23 3456 2722 3530 4641 2.0 2.0 2739 4663 6222 0.19 5003 2.8 0.06 2.0 5550 3823 2.9 3.2 7071 8331 6639 6159 2.7 2.5 ID# NED morph. Table C.1.

A103, page 33 of 44 A&A 581, A103 (2015) C HLR HMR M i ] pc pc

Z [HLR] Z log h M i ][ [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) [ A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA602 Name603606607608 NGC Type 4956609 NGC log 4961 UGC610 08231 UGC611 –1 08234 A612 NGC 4 5000 B 5 10.90 AB UGC614 08250 0 UGC630 A 08267 9.11 9.51 4.24633 NGC 3 5016 B 11.15 4634 A NGC 2.25 5029 2.65 2 AB652 10.74 NGC 5056 2.60 4.05 9.93657 10.81 NGC 3 5205 A 3.59663 –1 NGC 1.26 A 1.67 5216 0.00 3.40 2.54664 4 10.28 NGC AB 5218 2.66 11.41 UGC665 08662 3 10.53 B 0.00 0.16 0.36 3.35 2.18 UGC676 08733 0 3.91 A 0.09 2.37 1.84684 3.54 1 9.80 0.07 0.31 B IC 4 0944 UGC708 A 10.40 08778 0.96 0.01 9.98 2.36 5 UGC714 B 10.52 08781 1.10 0.81 3.46 2.57 8.80 3.83715 0.13 NGC 1.82 5378 8.09 8.77 2 9.18 3.65 0.60740 0.49 A NGC 0.47 1 5406 A 9.53 0.00 2 2.10744 B NGC 2.17 5485 10.06 0.57 2.43 UGC 1.88748 9.40 0.21 11.17 09067 9.09 2 B 0.00 11.07 2.45749 9.50 8.63 NGC 2 5520 3.32 0.19 B 0.14 8.20 1.38 3.72 8.17754 –1 10.49 NGC 0.09 A 5614 3.86 3 AB 1.20758 9.33 9.39 11.04 NGC 1.65 5631 10.04 0.14 10.65 3.77764 0.01 10.68 0.24 NGC 0.41 3 5633 2.33 9.84 9.23 A 2.58 3.96768 0.41 NGC 0.39 1 5630 9.19 4.20 8.45 2.25 A –0.27 3.28 –0.18769 0.16 NGC 0 9.81 5657 A 9.99 0.55 10.09 2.19774 0.17 11.21 0.13 NGC 0.41 3 5682 8.77 A 9.88 0.07 2.35 0.23775 10.54 NGC 9.30 5 3.97 5720 2.91 B 0.19 0.29 2.06 8.64 4.45778 0.33 10.10 8.32 –0.41 NGC –0.13 0.02 3 –0.51 5732 0.04 B 0.09 4.44 UGC779 09476 8.36 0.23 4 9.43 B 3026 0.01 3.44 UGC 9.13 0.15780 9.59 0.02 10.33 09537 0.67 2.13 3 –0.10 B 9.61 0.37 3114 2290 2454 2.90 UGC783 0.20 9.82 09542 –0.03 3 9.20 2.32 3.0 10.02 9.10 A 0.01 2596 –0.09 –0.23 2010 3 2.96 3.67 0.34789 A 11.02 NGC 0.07 5784 2.3 8.03 2.5 0.78 3939 2 2.41 UGC 10.05791 A 09598 0.44 9.89 2.14 8.61 10.13 3182 –0.05 5940 –0.09 4 3.60795 A 0.28 NGC 9.87 0.01 5609 1.57 5797 3.3 10.05 0.31 0.07 11.18 4257 2.23 UGC797 9.33 0.20 9.26 –0.04 09665 5862 4712 0.52 0 2.98 –0.33 2.95 9.57 2.2 A 3 9.48 10.28 AB 2.3 798 0.04 4582 NGC 2680 –0.84 1.62 5888 3.81 0.45 2.4 2.14804 0.52 10.48 9.56 –1 11.12 NGC –0.06 2094 6866 1.10 A 5908 3.02 9.35 0.03 6011 2 0.15 2.3 806 A 9.91 NGC 5551 1.75 5930 1.83 9.17 0.25 9.81 10.65 3038 0.56 –0.07 3.03 4.23 –0.10 UGC 3.0 807 0.42 9.98 09873 0.26 2 2.5 8.76 2.37 B 0.19 0.37 –0.40 1814 9.86 UGC809 9.39 09892 1 1.86 3.62 A 0.21 0.63 1254 2207 0.37 1 0.28 11.30 NGC 0.11 AB –0.09 3330 5971 2.4 8.28 1689 0.81 2.11 2 2.73 3.31 8.58 A 10.98 NGC 9.30 2355 2676 5966 3.0 9.50 0.18 10.46 0.04 0.12 0.37 0.24 1.09 3 3.70 2.5 A 2231 0.15 2.60 9.80 0.26 0.16 0.25 8.23 IC 3345 9.99 4566 4.05 10.09 2.2 2 NGC 0.05 AB 0.14 3.79 5987 8.58 0.42 10.24 2810 2.16 0.20 –1 2850 A 2.1 10.38 9.07 9.24 0.00 2.70 2.34 8.03 0.41 0.10 0.15 6635 2258 7114 3.09 0.22 8.92 0.13 0.39 10.95 9.78 2.82 2.4 –0.21 5272 1.04 2.49 3.83 4540 1 2 A B 0.27 0.01 9.12 3334 2.8 8.44 2.5 9.24 0.26 –0.21 4.00 2605 1.60 0.70 9.74 10.89 9.82 10.86 5563 0.73 1914 1.97 2.4 –0.16 1.44 6851 –0.06 8.85 8.69 0.18 3992 1488 2.28 –0.23 0.26 4348 9.18 4.33 2.3 3.68 2.8 9.56 0.38 1.11 2.4 0.02 2.70 0.14 –0.29 0.49 –0.32 8.40 9.29 0.07 1589 0.49 0.37 9.98 3119 –0.34 721 8.82 9.51 0.52 –0.31 –0.10 0.12 2.96 –0.18 2.29 9.84 2209 1498 2214 3.0 0.19 8.88 3.1 9.51 0.20 1269 –0.33 1837 0.00 3530 3.2 –0.28 2.1 0.16 8.77 0.49 0.37 3053 8.75 9.55 3045 –0.10 10.02 –0.17 1915 2.5 9.39 0.06 0.39 2354 3.1 0.48 0.32 10.09 9.69 –0.08 2262 7115 9.43 0.25 –0.30 2.6 4753 3362 4273 –0.32 8.92 2.4 2516 4053 8.70 9.89 9.92 6566 0.06 9.49 0.17 2.3 1.8 0.26 5811 4783 9.77 0.33 –0.11 4686 0.03 3.3 2.3 –0.34 5860 3589 –0.26 –0.20 4869 2243 –0.19 9.75 0.18 9.35 2663 2.2 3.4 –0.03 2437 0.06 0.23 2108 1830 3.1 –0.26 2.4 9098 –0.15 –0.33 4099 2436 7201 –0.04 0.10 3047 2.1 1888 5413 2.7 0.11 2.8 6166 2983 4654 4613 2.4 1950 2.5 3.0 0.16 3597 0.09 2731 2.8 5242 2737 3762 2307 2.2 3.0 ID# NED morph. Table C.1.

A103, page 34 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence C HLR HMR M i ] pc pc

Z [HLR] Z log h M i ][ [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) [ A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA810 Name813814815816 NGC Type 5980818 NGC log 6004 UGC820 10097821 NGC 3 6020 A822 NGC 3 6021 B –1 A UGC823 10.64 10123824 –1 10.57 NGC 11.31 A 6032 3.56826 –1 NGC A 6060 10.87 1 UGC 3.46 4.48827 A 10205 10.89828 NGC 3 6063 3.73 B 10.32 2.46829 2 4.37 A IC 0 1199 2.09 2.87831 A 10.40 NGC 6081 3.51 UGC832 10.91 10297 0.70 3 2.77 A 10.99 UGC 3.41833 10331 0.62 0.09 2.70 3.88834 2 0.61 NGC AB 0 9.98 6125 2.41 3.82 A 0.00 4835 A 0.30 0.00 NGC 6132 10.57 4 AB 0.70 2.01836 11.00 NGC 2.89 6146 0.01 1.11 9.33 2.49837 –1 9.61 NGC 9.93 3.54 A 6154 2.34 0.01 4.21 UGC 10.11838 10380 9.28 0.83 3 A 0.68 11.24 2.65840 –1 NGC 0.67 2.67 A 1.91 6150 9.61 1.46841 0.18 10.16 NGC 2.30 1 6155 10.06 4.25 B 2 11.53 AB 2.74 UGC842 0.52 10384 8.82 0.27 9.42 0.25 1.95 3.16 UGC843 10.90 –1 9.74 11.00 10388 1.84 A 4.19 8.94 0.71844 NGC 9.43 0.90 4 6173 2.89 0.25 A 11.30 3.36 2 3.84845 9.86 A 9.73 NGC 0.59 6168 0.41 1 9.82 0.66 9.64 AB 1.82846 0.14 NGC 9.98 –0.08 6186 2.73 3.78 10.28 0.00 UGC848 9.08 0.26 10.50 –1 10650 0.41 A 0.40 9.61 0.13 2.38 2.20849 4 NGC AB 0.51 3.06 6278 3.69 9.61 9.28 0.00 11.68 0.18 3.86 UGC850 9.99 0.21 10693 2 2.68 9.13 B 0.02 9.64 –0.18 4 UGC851 9.44 0.48 A 9.02 10695 0.66 0.37 4.15 9.01 0.00 –0.07852 0 0.32 10.46 NGC AB 2.12 –1 6310 2.39 10.00 0.08 AB 2.85 0.01 2.42 9.24 8.96 0.35854 0.07 NGC 4507 6301 –0.15 10.69 11.35 –1 A 8.90 3.78 0.08856 9.38 NGC 2957 6314 2.77 10.01 9.81 0.07 0.18 0.00 0.68 1.60 3473 2.35 –0.27 2.3 857 NGC 11.17 0.22 4808 4.47 –0.13 2 6338 4.00 8.78 A 1.96 2295 8.55 UGC 10.01858 9.79 10796 4383 3 0.43 3.3 A 9.90 1.08 2765 –0.07 –0.27 0.18 2.44 UGC 2.0 3.81 0.06859 10.43 10811 2679 1 10.08 A 2195 3266 1.36 0.00860 –1 11.04 0.87 –0.12 1586 2.85 –0.14 A 2.49 3.0 8.52 4 9.80 0.06 0.00 IC AB 2303 1256 3.68 3.3 861 10.92 NGC 0.85 –0.45 5871 6394 2.5 9.06 11.55 9.30 –0.65 10.14 2 2.97 0.56 UGC 2.30862 9.64 B 10905 –0.20 0.51 9.30 4437 9.42 0.00 0.39 5662 0.00 5579 3.84863 2.0 0.30 NGC 6411 10.04 4.11 9.80 10.86 4389 3630 0.05 2.47864 2 0.29 NGC AB 2.40 3 0.27 6427 0.15 2.2 0.00 2.5 0.00 B 0.10 0.27 –0.17 3636 0 2.03 UGC 8.96865 A –0.91 10972 3.45 10.43 2802 2.38 8.75 9.61866 8.50 –1 0.01 10.86 NGC 9.21 0.26 0.43 –0.11 A 6478 2.71 1.9 0.18 4917 11.35 0 7.77 NGC 3500 AB 0.30 1.37 10.09 3.08 2654 6497 3533 9.93 9.89 10.81 0.27 4180 3 3.18 0.30 A NGC 2313 –0.11 2286 0.40 2.1 6515 4.01 2.17 10.67 0.08 3624 0.33 3.4 UGC 2.4 10.07 0.29 11228 3259 4 8.61 2.4 3.75 A 10.41 0.36 –0.19 –0.34 0.16 UGC 0.14 0.37 0.18 11262 2642 2.17 4.49 1 9.30 B 0.00 4043 0.19 2.39 3.2 –1 10.99 9.92 5577 A 9.75 2.69 3.31 0.33 8.12 0.23 0.31 2340 0 9.66 B 9.57 10.89 4113 2.64 2.5 0.35 11.11 0.72 6828 6057 4 3.65 3.4 A 9.54 9.66 –0.81 –0.19 2.95 –0.60 0.43 5696 9.99 10.97 3717 0.24 0.00 3.56 5798 2.8 2.07 4.11 0.26 2.3 8.20 0.09 10.20 0.00 0.38 4977 –0.72 9.31 4.19 0.33 0.16 0.28 0.15 9.97 2491 4188 2.27 2.9 0.31 2.85 9.01 0.00 –0.43 3068 2567 2446 0.37 2.35 0.29 2.38 2.2 0.00 9.35 9.10 2.6 2415 9.86 9.72 0.61 6768 3.1 0.12 2.51 0.12 9.81 –0.44 8.52 3244 5728 0.00 0.21 1.54 0.13 9.52 0.11 0.15 3.1 0.44 9.24 2754 10.08 0.30 2.0 0.12 0.20 0.19 3024 3578 0.21 0.00 9.87 0.47 9.08 0.33 1868 2366 3852 9.16 0.00 6271 –0.47 2.4 2.5 9.76 1412 9.64 0.19 0.59 5424 3.4 6736 –0.02 9.55 9.75 3.3 9.88 0.11 9.95 4124 0.06 10.02 2.8 0.31 0.30 8.84 –0.37 3861 –0.02 9.47 10836 0.22 3220 8.76 0.00 8945 2.3 0.28 4807 2.0 9.31 0.33 3325 6892 9.64 3675 9.75 3467 5783 0.06 3.3 0.04 0.00 7769 2.4 2.8 0.05 8.23 6226 0.10 0.24 2.7 0.07 4814 0.19 6792 0.26 4466 0.36 5410 5489 0.14 2.2 2.1 4081 3196 –0.46 1682 –0.01 3.5 2788 1235 3.0 0.21 4887 3.3 0.05 0.15 3390 7164 –0.23 2.3 4767 5308 2.0 4903 4011 3770 4083 6216 2.7 2648 3.1 7062 3.4 2.2 ID# NED morph. Table C.1.

A103, page 35 of 44 A&A 581, A103 (2015) C HLR HMR M i ] pc pc

Z [HLR] Z log h M i ][ [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) [ A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA867 Name868869870871 MCG-02-51-004 NGC Type 6762872 log 2873 A NGC 6941874 NGC 1 6945 A 10.65875 NGC 6978 UGC 11680NED 01 UGC876 10.16 11649 2 3.62 B877 2 0 B B 3.55878 2 10.99 NGC AB 7025 1 UGC 11.07879 B 10.68 11694 2.22 10.90 3.60880 NGC 7047 10.39 3.68 2.78 UGC 4.36881 MCG-01-54-016 11717 3.76 0 A 0.44 0883 A 3.55 4 2.25 UGC884 11.18 A 11740 0.00 3 B 0.35 10.98 2.45 1 2.68 UGC885 A 11792 2.45 4.46 9.04886 0.16 10.67 NGC 0.18 7194 4.36 2.14 10.81 3 UGC887 A 0.54 11958 0.47 0.65 9.93 3 UGC 3.26 2.24888 A 0.25 11982 3.69 10.36 0.19 2.93 0.24 UGC889 –1 0.00 9.70 12054 A 2.53 0.19 10.26 0890 A NGC 7311 2.78 0.12 11.32 4 2.23 1.40891 A NGC 9.76 0.31 7321 1.94 3.21 11.27 0.82 4 UGC 9.90892 8.86 A 12127 9.90 4.31 9.78 9.64893 0.07 NGC 0.13 0.27 1 7364 1.99 3.76 9.61 A 0.15 9.84 0.62 8.87 UGC894 12185 3 2.38 B –1 A 2.67 UGC895 0.17 0.33 11.09 12224 VV 488NED 2.73 9.31 0.13 02 0.35 10.02 2.86 –0.07896 11.03 11.43 1 9.26 2.30 10.02 9.58 A 1.19 NGC 9.37 2 4.43898 7436B B 2 AB 0.09 0.52 0.07 4 1.67 UGC 3.66 4.01900 9.30 A 10.68 12274 9.90 8.89 0.60 9.76 10.68 10.34 0.17 1.79 UGC901 12308 –0.12 –1 0.01 0.28 A 9.91 2.92 3.87902 9.83 NGC –0.10 0.61 7466 3.83 9.60 3.57 0.03 0.27 9.74 –0.10 1 2.26 2.54903 11.47 A NGC 0.63 7489 0.11 9.23 4 2.83 5738 –0.03904 0.47 A NGC 0.34 7550 10.08 9.15 8.29 10.99 4.16 3751 2.46906 9.26 0.44 NGC 0.33 3 0.00 7549 10.08 0.28 2.33 A 2.31 2.2 0.20 8.71907 0.25 0.35 NGC 1581 3 0.04 7563 3.74 0.14 A 0.25 1.60909 8.58 –1 0.31 10.74 8.94 0.00 NGC 1240 0.11 A 7562 0.16 2.64 0.17 2.9 2.47 0.90910 9.00 10.51 9.21 NGC 6592 3 7591 –1.22 B 9.74 5518 0.11 11.12 3.63911 0.29 9.92 0.30 5325 4942 2442 0.34 1 2.41 4132 B 9.82 0.17 IC 0.49 5309 3.20 2.1 0.23 10.07912 –1 4023 10.50 NGC 9.74 1787 0.08 2.5 A 3470 7608 4.12 0.19 2.4 1.19 –0.17 UGC 3.0 914 0.03 10.93 12519 2354 3 8.38 B 11.10 –0.70 0.23 2.15 –0.23 UGC 3.48 0.04915 2.1 9.59 12518 8.39 0.05 10.04 0.27 3166 9.68 1.76 4.32 0.22916 4 10.74 NGC AB 0.40 3689 3 7619 2.66 4.47 9.50 A 4 0.02 0.39 2368 AB917 9.53 9.75 NGC 1829 0.99 2811 7623 3.3 10.12 9.06 0.09 2 2.06 3.71 2.9 919 0.06 10.07 9.96 A NGC 6172 0.84 –0.18 2303 9.90 –0.31 7631 6527 0.31 2.93920 2.3 0.50 –1 NGC 4959 3.22 A –0.73 7653 3.11 10.09 8.99 0.23 9.27 3180 3.06 9.98 2.0 0.29 NGC 9.06 5392 0.80 0 3.01 7671 2.9 0.23 A 0.28 0.00 11.25 4178 2.46 8.67 NGC 0.30 4449 0.32 2 7683 8.81 A 0.35 3690 0.21 3.57 0.17 2.4 0.47 10.63 9.26 NGC 1.91 4377 2 7684 9.78 4.53 2.4 A –0.14 1.95 0.01 10.45 8.97 NGC 3336 7371 1.45 0 2.10 7691 A 9.71 0.24 0.00 9.79 0.11 4.04 3.3 –0.39 10.51 6020 0 0.09 A 0.82 4393 2.48 3.58 2.8 0.15 0.79 10.90 9.26 0.74 1277 0 –0.36 3.20 7.86 0.07 A 3348 0.61 3.72 10.04 10.96 1006 3 2.3 10.04 0.49 8.96 0.36 B 2.69 4.57 2.7 0.45 10.83 3243 0.85 –0.06 0.12 –0.13 8.21 0.62 0.16 0.03 5868 2.47 4.36 9.53 10.11 2106 8.92 6168 –0.04 5310 2.16 3.1 4.25 0.43 3862 0.09 –0.44 9.14 8.66 3.1 0.00 2.97 4498 2.4 2.77 3545 9.25 9.95 3126 0.35 0.15 9.47 9.80 0.07 2832 3.07 2401 0.01 1981 0.32 4171 –0.28 0.25 2.8 3.1 2.71 3.1 0.26 9.95 0.21 3154 0.23 8.70 9.99 –0.09 1.65 1.8 0.31 5955 0.73 –0.23 8.93 10.06 0.04 5228 0.31 –0.27 0.34 6073 8.56 0.25 –0.02 2.9 8.83 0.07 9.79 4375 1982 0.26 0.31 2.7 0.07 9.72 1226 –0.22 9.53 6059 –0.16 10.09 2.5 9.82 0.11 6686 4131 –0.30 9.78 0.29 9.67 3917 –0.14 2.7 0.27 3694 –0.41 2.3 9.93 4078 9.20 2531 –0.56 9.38 3284 2.9 8.84 0.06 2308 9.84 2.1 0.37 –0.14 2616 –0.08 1926 0.04 –0.18 10.01 2414 3745 3.5 3.1 2813 –0.04 9.76 3694 3.0 4521 –0.11 0.17 8.67 2570 3441 0.30 0.38 3204 2.4 2913 2.3 0.08 0.34 2.3 0.14 2074 0.40 –0.13 2564 1628 –0.25 2118 2126 2.8 0.18 3.3 1562 3602 0.27 3365 3.4 2714 1943 0.02 2.6 2140 –0.34 2.7 1444 2273 3.2 1863 3290 4901 3.2 2330 4146 3.1 2.0 ID# NED morph. Table C.1.

A103, page 36 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence C HLR HMR M i ] pc pc

Z [HLR] Z log h M i ][ [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) [ A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA921 Name923924925 UGC926 12653 Type927 NGC log 7711928 NGC 7716 4929 A NGC 7722 UGC930 –1 12723 A 9.53931 NGC 2 7738 A 10.76 UGC932 12767 1 A 4 2.58 UGC933 A 10.30 12810 4.35 UGC934 10.92 12816 2 B NGC 9.69 2 7783NED 4.10935 B 01 NGC 7782 1.32 3 4.06936 B 11.06 2.77 1 10.75 A 4 2.38937 A NGC 7787 10.78 2.46 UGC 3.94938 12857 0.00 2 11.25 3.61 A 0.30 9.85 2.75 UGC939 MCG-01-01-012 12864 3.38 1.87 1 0.16 11.24 3.64 AB 0.01 0.00 1 AB 3 2.75 2.27 A NGC 0.65 7800 2.07 10.52 4 4.02 0.08 10.62 B NGC 0.79 5947 1.98 NGC 9.46 0.57 4676B 9.47 2.86 3.52 2.01 10.07 0.03 1.46 3.37 9.92 5 0.48 AB 1.19 2.61 10.05 2.86 0.08 3 0.00 B 0.11 8.98 1 2.73 B 9.77 0.28 1.99 0.30 2.45 7.93 10.56 0.38 0.01 8.16 10.92 9.64 2.30 2.01 0.42 8.69 1.58 9.95 1.61 3.57 9.49 0.13 0.77 3.76 9.02 0.57 9.52 0.34 9.49 1.64 0.20 9.13 0.46 –0.24 7.93 0.29 10.07 1.76 0.47 2.58 9.46 9.08 0.24 0.37 0.09 9.00 8.59 9.90 0.39 0.16 0.39 9.86 0.37 0.04 9.23 –0.41 8.24 0.06 9.59 0.17 8.76 –0.32 0.20 –0.15 0.14 4547 9.11 –0.36 8.02 2554 0.25 9.59 3078 –1.04 0.32 9.45 1525 1943 –0.45 2.3 8.69 9.89 3.3 0.30 0.38 1143 8.48 2.7 0.32 3097 5140 –0.53 –0.03 2560 4693 8.23 –0.06 0.15 –0.76 6037 2.6 2.1 –0.03 8.82 3308 5670 8345 9.54 0.26 –0.44 3.0 3800 5908 4353 0.18 5173 2.4 2.0 0.21 –0.37 3003 4188 –0.24 6197 2.4 3.0 0.34 4501 –0.18 0.22 4869 2.2 4977 2723 2041 –0.47 4155 2.6 1783 5227 2.6 2.5 4619 0.07 0.18 2.4 1598 1639 2.4 5225 3912 2688 3032 2.4 2.4 ID# NED morph. Table C.1.

A103, page 37 of 44 A&A 581, A103 (2015) C HLR HMR M i ) pc pc

Z [HLR] Z log h M i )( [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) ( A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M . GMe Stellar population properties: CALIFA001 Name002003004005 Type IC 5376 UGC007 00005 log 008 NGC 7819 UGC010 00029 3012 A 2 A IC 1528 UGC013 00036 4 A –1 11.16 A014 10.82 NGC 0001016 10.39 NGC 11.41 0036 3.60 1 AB 4.05 UGC017 3 MCG-02-02-030 00148 AB 3.37 3.90019 11.13 3 10.42 A 2 AB UGC020 MCG-02-02-040 00312 2 2.54 UGC B 00335NED 4.12 4 2.64 02023 A 10.96 10.57 3.49 4 AB 1.91 –1 2.73024 A 11.21 ESO 540-G 10.31 003 0.66 5 4.13 3.79025 10.21 B 0.41 10.77 2.92 3.88026 NGC 0.62 2.29 0.43 2 0160 3.18 AB 0.36 10.11 3.38027 0.49 NGC 4.20 0171 10.16 2.80 2.71028 0.06 0.05 NGC 0.77 0177 2.98 0.79 2.62029 NGC 1 0180 2.39 A 3.42 9.58 2.22 10.00 0.21030 NGC 2.50 0.47 0.50 2 0192 0.27 B031 11.09 8.59 NGC 9.96 0.38 1 0216 1.87 A 0.90032 0.61 0.07 10.88 NGC 1.21 2 0214 B 0.51 2.44 4.26033 1 9.74 0.29 10.49 NGC AB 0217 9.20 0.39 0.37 4.01035 9.50 9.00 0.19 11.06 NGC 5 0237 0.04 10.99 A 4.09036 0.28 3 MCG-02-03-015 9.57 9.71 NGC AB 0234 9.12 9.82 0.26 2.66 3.77038 9.90 4.43 1 9.24 11.15 A 1 9.01 0.09 AB 2.54039 8.92 NGC 10.05 4 0257 9.59 B 8.91 2.20040 4 11.03 10.95 0.09 0.00 NGC AB 0.44 4.15 2.78 0364 8.56 2.41041 10.52 NGC 2.86 0.40 0429 9.09 9.36 10.92 0.06 0.17 9.74 4.23 4.02042 0.27 NGC 1.01 4 0447 9.46 A –1 3.50043 AB 8.76 0.01 NGC 2.71 3.68 0.21 2.24 9.74 0444 9.09 0.16 1.22 –0.26 –0.05 UGC 0.01044 0.08 10.99 11.02 00809 1 A 2.86 2.70 UGC045 8.63 0.09 9.99 00841 –0.33 –0.01 1 –0.00 –0.08 0.31 B 0.77 0.29 4.38 3.86 2.43 10.02046 10.91 NGC 9.11 2.45 3883 6842 4 0477 –0.08 A 4 –0.63047 A 9.72 11.22 2324 5175 0.43 0.70 –0.08 0.10 0.34 0.02 0.13 –0.37 IC 5455 5400 3 1683 4.12 3.3 2.1 049 A 10.12 NGC 9.48 0.09 0499 9.34 0.75 3561 –0.18 –0.08 4542 2.60 9.84 2.46 4.19 –1.14050 3 0.06 0.15 NGC AB 0496 2.1 9.82 3.1 10.29 –0.10 2.90051 9.35 –0.34 0.36 NGC 4348 4104 –1.00 0504 0.32 10.80 9.71 8.67 –0.35 3.10 2.64052 2 –1 –0.19 3204 3355 NGC 0.64 AB 0.66 2781 3517 A 0517 3.10 9.22 2.1 2.7 10.09 2.60 UGC053 9.66 00987 1871 2872 3.47 4 –0.77 5717 10.83 9.22 A 11.41 4820 9.48 1.91 3.0 2.2 0.00059 0.30 NGC 9.46 0.00 4621 3264 0 –0.07 0528 –0.06 8.83 A 3298 4739 1.89061 2.5 10.73 NGC 0.04 2187 3.99 0.20 0 0529 2.10 4.58 2.2 A 1 2205 AB 9.01 5055 8.81 2.2 062 0.01 10.69 NGC 0.32 2.18 3.2 –0.15 0551 10.00 4256 3.13 UGC068 9.67 10.91 2362 0.03 10.84 9.43 01057 0.67 0 9.51 A 2.1 0.07 –0.38 0.71 4.21 UGC 0.02 1935 0.08 –1 01271 2.68 A 3.25 8.84 0.48 2.6 4.31 8.82 4.20 3 0.12 0.23 11.12 9.90 NGC AB 0681 4 0.27 11.29 –0.01 AB –0.45 2.11 UGC 0.01 01274 9.89 4165 0.12 10.87 1.05 0.14 2.93 0 4.30 –0.04 9.69 10.58 B 9.37 NGC 2985 0741 4.54 9.05 0.01 0.09 –0.14 4523 2.97 2.97 3.0 1 8.83 AB 0.61 3.51 –0.03 0.42 0.03 9.26 10.99 –0.02 3344 3268 3.39 1 A –0.16 –0.35 0.01 9.76 2.0 10.39 9.52 1503 3.04 0.04 –1 3916 0.36 6498 0.06 A 3.08 3.8 4.14 9.56 0.15 11.28 0.00 0.03 2513 –0.08 5109 2.49 4.17 –0.11 10.01 9.02 9.13 11.64 4346 1254 –0.14 2.27 2.6 2.2 0.25 0.00 –0.10 0.00 4.27 8.76 3307 1360 0.01 8.98 9.07 3868 2.4 2.5 4.40 2.84 0.19 0.13 4431 9.09 10.10 0.00 0.45 3128 2.85 3457 0.00 0.22 4954 2530 –0.55 0.16 3.0 9.99 9.97 2542 –0.42 3.20 3.1 0.12 9.95 4090 0.00 9.35 0.24 2.4 –0.15 1.9 3.11 0.67 –0.26 9.94 6099 0.06 10.03 9.12 0.01 3268 0.39 0.06 9.73 4623 0.22 0.13 9.85 2477 –0.76 0.16 2.1 9.54 3.1 9.62 9.73 0.20 0.21 –0.16 –0.10 2564 –0.23 10.00 0.00 1999 4856 5690 –0.25 9.76 9.76 9.73 –0.17 3.2 0.13 3509 4866 4220 9.98 4535 9.26 3.3 2.3 3474 –0.06 0.01 8.85 10.07 0.11 0.10 3450 2.3 6329 2.5 9.70 –0.51 4599 0.22 0.21 2.2 9.54 0.14 3677 3210 0.06 0.02 9.80 2616 2728 6329 0.03 0.07 9.82 2.6 3.0 4849 0.17 1838 2.0 0.14 0.20 2922 1471 –0.07 2323 3.2 –0.02 2061 0.09 –0.24 1780 3.0 0.20 3.5 2751 2859 0.14 –0.16 5228 2046 2505 4918 3.1 3.1 4260 2965 2.1 0.06 3013 2.4 0.20 1663 2543 1174 3.1 2.4 3791 4620 2787 4159 3.1 2.8 ID# NED morph. Table C.2.

A103, page 38 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence C HLR HMR M i ) pc pc

Z [HLR] Z log h M i )( [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) ( A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA069 Name070071072073 NGC Type 0755076 log IC 1755077 NGC 0768081 NGC 4 0774 B087 NGC 0776088 NGC 2 4 9.54 0810 A B100 NGC 0 0825 A UGC101 11.19 10.90 01659 2 2.94 B103 –1 10.99 NGC A 0932 4.03 UGC 3.60104 11.10 01938 1 A 11.58 4 4.16108 B NGC 2.24 1056 4.07115 10.60 NGC 0 1060 4.18 A 3 10.98 AB 2.79 2.34 UGC116 02222 0.42 2.85 4.26 UGC118 10.69 11.17 02229 1 3.25 A 2.43119 –1 NGC 0.35 0.12 A 1093 2.87 0.21 0 AB 3.60 4.29 UGC127 10.20 02403 0.37 11.79 0 3.03 UGC131 0.08 10.93 A 0.17 02405 0.61 2.15 0.56 4.24 UGC132 0.02 02465 3 4.65 B 8.90 11.19 2.34 4.43 2.93 2133 0.16 B NGC 0.00 1167 0.02 0.42 3 10.05134 A 10.87 NGC 9.76 1349 3.86 10.81 0 3.14 10.01135 0.24 A NGC 0.58 0.14 1542 3.38 0.24 10.69 2.98 UGC 3.79141 03038 9.48 0 4.25 A 11.04 9.99 UGC 8.74 0.15143 0.08 –1 03107 0.91 A 2.61 3.32 0.17 10.08144 1 9.76 11.44 NGC 0.31 AB 1645 4.33 9.04 11.35 9.83 1 2.53 UGC146 9.71 0.54 A 03151 2.62 10.67 0.00 0.14 2 4.42 0.01147 9.52 A 9.92 2.04 4.17 9.27 10.99 IC –0.08 2095148 9.81 NGC 3.86 0.16 0 1677 2.94 B 1 0.08 10.97 AB 2.21149 9.70 9.38 0.13 –0.01 10.07 3.75 IC 0.30 2.91 2101 UGC150 11.03 9.99 9.09 0.08 0.13 11.07 03253 2.63 3.68 0.42 0.62151 4 4 NGC AB AB –0.41 2.80 2253 9.12 9.80 0.41 0.19 9.99 4.06 UGC 4.20152 0.20 03539 0.20 2.41 0.12 8.98 9.77 –0.11 0.22 2 0.04153 4 B 0.15 NGC 9.85 AB 2347 2.53 9.40 9.86 0.31 9.19 1638 UGC 0.10155 0.00 03899 0.01 3 9.80 10.47 B 4 0.01 9.72 10.79 2.10 3.09 1304 AB 0.79 2.89 2.94156 NGC –0.01 6604 2410 0.27 2.2 0.16 0.18 9.98 5650 0.81 UGC 0.13157 3 9.70 10.11 10.82 03944 AB 3.39 4297 3.86 0.40 4535 –0.10 3087 5 10.05 UGC158 2.5 A 03969 0.57 0.21 –0.91 3.0 11.08 9.33 0.21 9.94 0.17 2397 1.48 2.31 –0.02 4943 3.37 UGC 4.00159 2 9.27 03995 AB 3.0 9.71 0.15 3 3152 5841 AB 9.55 0.03161 8.55 0.06 –0.50 NGC 2.14 4.09 1806 2449 0.07 2.53 2.1 11.17 9.76 4699 2 UGC165 9.63 10.21 A 0.12 0.59 0.03 04029 1289 8576 3.0 9.68 –0.43 2.30 2 2.73 2.57 UGC 3.2 171 9.64 B 04054 4.19 6395 0.14 0.13 10.97 0.99 5081 3.19 0.42 0.26 0.24 10.02186 1 9.46 2.2 9.80 0.21 3033 AB 2.84 0.03 4 IC 3052 11.27 AB 9.51 0480 –0.13219 NGC 2544 1.27 0.76 2480 2.5 3.50 0.35 11.17 958 0.07 4 1.92 3.2 UGC 0.03226 9.05 10.57 A 0.21 04132 2.85 4064 4.28 0.30 2.00 807 8.42 8.61 0.36 9.34 0.29 0.06 NGC 2826 3692 4.09 2513 0.07 –0.08 3.56 9.64 3.2 3.0 4 9.63 2206 0.10 AB 0.31 –0.27 5 9.41 2.57 0.27 9.21 A 3 IC 0.61 9.81 –0.10 5458 AB 3.1 2247 NGC 0.32 2639 2.77 10.45 3907 3.26 –0.11 UGC 0.12 0.29 11.16 9.32 –1 04659 9.33 4252 2.96 A 9.71 3747 3.2 0.29 0.88 2.32 7902 –0.10 0.12 0.04 3129 3.04 1852 8.49 8.95 9.76 11.42 4987 4.08 0.33 2.6 –0.03 0.11 2.0 –0.22 1 1 3.01 0.03 2.2 A A 0.53 0.41 3571 5 8.98 2.05 9.18 A 8.62 0.86 –0.25 4.39 9.79 0.16 2530 4277 9.48 10.72 11.25 2.19 –0.52 9.08 2.8 0.30 8.92 2816 4315 5199 2.67 9.11 0.35 –0.68 –1.38 0.56 2.00 2.8 0.06 2146 9.65 3946 6289 9.14 3.95 4.34 –0.05 2.5 3.04 2.6 1.15 9.99 4937 7058 2.16 –0.08 0.07 0.15 1.37 8.44 9.91 2.4 0.28 4807 9.50 9.77 –0.29 3379 –0.26 8.79 0.50 2.5 –0.81 –0.17 4868 0.02 2.57 3.32 0.45 2618 0.13 3653 1.56 0.06 9.28 3.2 9.58 –0.08 2.4 0.00 –0.32 9.61 2145 1736 1.14 0.05 –0.61 8.91 –0.35 –0.33 9.49 0.01 9.36 1659 1654 0.31 –0.31 9.07 8.59 2.8 2.5 3797 0.43 0.22 4783 –0.23 9.98 2282 –0.27 0.31 3253 3045 3523 8.94 2.3 –0.68 2.5 0.12 2117 2413 –0.16 –0.31 8.99 2.8 0.07 2.1 9.59 9.87 3939 –0.25 9.10 9.04 2627 8.49 –0.30 2253 2.6 9.70 4192 3851 1805 –0.52 2.9 0.01 2913 2526 –0.42 –0.29 0.02 3.1 6325 2.2 9.38 9.53 0.04 –0.58 4899 8.82 2.4 –0.39 5286 0.19 4598 4073 3636 3155 0.04 2.4 3343 2.7 –0.29 –0.30 2.7 2641 0.02 0.13 1802 –0.97 0.12 2.7 5759 5756 2662 4316 3763 2319 2.3 –0.39 2.3 0.06 2.6 –0.41 3929 3508 3.1 4250 2830 2189 2858 2486 1657 3.0 2.7 2.2 ID# NED morph. Table C.2.

A103, page 39 of 44 A&A 581, A103 (2015) C HLR HMR M i ) pc pc

Z [HLR] Z log h M i )( [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) ( A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA231 Name232272273 UGC274 04722 Type275 NGC log 2730277 NGC 2880 5278 A IC 2487279 4 B IC –1 8.84 0540306 AB NGC 2906307 10.63 10.35 NGC 2916 2.00 UGC309 4 05108 AB 4.70 3.06311 1 NGC AB 3 2918 10.67 A UGC312 05358 3 10.14 A 1.38 2 UGC314 10.72 B 05359 3.45 3.20 1.85 UGC318 –1 10.96 05396 3.63 A 11.23 5 4.14319 B NGC 0.25 3106 UGC 11.31 05498NED 2 01 3.83326 B NGC 0.27 2.38 0.52 3057 4.04 3 AB 9.57339 0.10 1 2.76 A 3.99 11.05 2.80 0.00341 10.65 0.22 NGC 1 3158 A 0.57 2.57 10.97 2.38353 NGC 5 3160 2.68 3.48 B 0.78 3.25 UGC361 11.45 8.38 05598 0.61 2.78 0.27 3.75 10.05364 –1 NGC 9.15 0.10 A 9.26 3300 0.33 1.06 4.27 1.50 UGC381 1 0.21 05771 AB 2.33 11.77 0.00 2.23 2386 0.05 A NGC 3381 2.49 11.05 9.62 0.17 2.75 UGC387 05990 0.21 0 4.30 8.39 9.60 B 0.00 –1 10.56 0.09 A 2.82 UGC 9.83388 9.81 06036 0.50 3.74 8.76414 0.60 0.16 9.68 10.84 11.39 5 1.54 3.50 B 0.08 IC 9.58 4 0674 UGC 0.35436 A 06312 0.19 3.04 9.81 0.28 1 9.23 4.20 4.04437 A NGC 2.65 –0.73 9.85 3615 0.19 9.59 9.42475 0.02 0.08 NGC 8.99 3614 –0.16 2.19 9.20 10.07 11.30 0.01 1476 9.49 A NGC 1 3687 3.09 9.18 0.13 B 0.35 3.03 9.75 2.79 2.82479 9.40 –1 NGC A 3811 4.06 0.00 11.10 –0.22 –0.51 1.10 10.04486 3 9.49 11.26 NGC AB 3815 0.01 11.53 –0.08 –0.22 –0.49489 NGC 0.20 2 0.06 2.10 3991 3.89 10.37 8.87 B 0.12 0.42 8.91 9.60 2.52 4.08500 NGC 2355 3 3994 9.10 3.06 4.43 B 0.06 9.99502 0.01 0.08 –0.37 10.55 NGC 0.18 1192 1852 9.64 3.39 3 5584 4003 9.99 A 0.28 UGC 2.1 515 9.83 0.22 –0.30 10.67 07012 0.43 4505 5 2.67 946 A 9.28 0.07 2.74 3.93518 3 1.9 10.78 NGC –0.09 AB 4047 3.3 –0.74 2.98 –0.11 0.15 6137 UGC 10.03 3.95 8.61 10.04528 0.34 10.03 07145 –0.00 2.22 0 10.75 0.07 B 4 0.10 4340 2154 AB 0.62 3.74548 9.77 0.00 –0.03 NGC 0.17 4149 2.5 1851 2932 0.00 9.75 2.74 2.50 0.22569 0.01 11.26 NGC 9.83 4.49 3 4185 2.7 A 0.33 8.22 0.33 2261 5141 –0.69 3 –0.60 2.48580 8.84 A 0.09 9.14 NGC 4210 2.1 10.02 0.00 4291 4860 –0.61 2.43 3.99588 1 9.57 10.98 9.88 –0.39 AB 3.01 0.08 2.2 0.08 –0.15 10.70 IC 3245 4547 1.90 0776589 3 0.19 NGC AB 3.06 0.54 4017 4470 8278 2.8 10.58 9.77 3315 –0.10 3.84 10.07 0.02592 0.04 NGC 0.43 3325 2 4644 5839 3.0 10.10 3.17 10.94 B 6220 –0.15 2.61 8.75593 2.4 0.13 4434 2.4 NGC –0.37 0.25 1.91 4.03 4711 9.37 9.76 9.90 4919 0.94 –0.03 0.17 3505 10.50 NGC 0.14 3.49 4898 5 4 4816 2.1 A A 3.1 NGC 2.80 0.21 0.12 4841A 9.97 4461 0.32 2 2.12 0.41 A 0.56 2209 0.04 3.64 3.1 9.61 9.67 10.28 NGC 9.37 2.80 –0.24 3 9.47 4874 A 9.72 2011 –0.08 UGC –1 0.01 10.74 9.84 08107 0.38 2.41 A –1.42 –1 0.33 1.9 0.09 A 0.10 5293 0.23 –0.02 9.31 3.16 10.75 7.33 2.13 9.51 5313 11.43 1.27 5447 4990 3.71 2.40 11.43 0.38 0 3.1 9.20 A 4282 0.08 0.04 3516 1 3.52 –0.04 A 9.74 0.01 –0.16 2.7 2939 4503 4.11 0.30 8.89 9.14 2.5 0.13 –1.13 4.22 2.40 11.69 2430 0.15 1.35 3261 0.02 9.06 0.11 11.39 –0.27 2.67 2.7 3.4 9.40 8.25 9.83 2402 8.77 2.32 3.90 0.09 0.29 –0.20 1264 –0.10 2.60 3.78 9.41 0.05 2046 2.74 –0.04 1283 0.02 10.04 4847 –0.07 2.0 9.68 0.10 8.43 2.8 0.37 –0.27 4483 0.65 –0.14 0.24 0.08 4454 2.77 10.00 3.4 0.23 0.04 4925 2.49 8.87 0.14 –0.12 3509 0.12 9.14 0.44 3618 4131 3.0 0.00 3456 –0.33 9.39 0.00 2.9 8.93 3481 0.00 9.45 9.01 0.14 2899 –0.16 0.69 0.06 3.3 9.83 2809 2.2 0.00 9.08 9.54 –0.25 2090 3827 10.05 0.05 0.39 –0.18 10.04 –0.21 2.2 2330 3764 0.02 2.1 3458 –0.13 8.64 0.17 –0.22 3613 8.38 2.6 9.95 2091 9.30 9.73 5230 –0.15 1318 –0.18 –0.21 1.8 8.93 9.74 2.7 9.70 2864 5263 –0.12 –0.03 2014 4032 0.02 3434 –0.67 2.7 7080 2.8 –0.05 9.78 0.12 2819 5345 9.25 3268 5782 –0.01 2.3 2.3 0.10 0.13 2168 5233 2.6 3642 2.0 –0.42 –0.22 2990 –0.07 0.16 1.9 –0.21 0.04 0.16 3582 2797 0.12 3019 2460 3530 4641 2.0 2.0 2896 4861 6222 0.13 2.8 5003 –0.05 2.0 5842 4254 2.9 3.2 7071 8331 6771 5755 2.7 2.5 ID# NED morph. Table C.2.

A103, page 40 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence C HLR HMR M i ) pc pc

Z [HLR] Z log h M i )( [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) ( A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA602 Name603606607608 NGC Type 4956609 NGC log 4961 UGC610 08231 UGC611 –1 08234 A612 NGC 4 5000 B 5 11.14 AB UGC614 08250 0 UGC630 A 08267 9.38 9.96 4.33633 NGC 3 5016 B 11.40 4634 A NGC 2.48 5029 3.21 2 AB652 10.99 NGC 5056 2.90 4.26 10.28657 11.05 NGC 3 5205 A 3.88663 –1 NGC 1.48 A 2.11 5216 2.79 0.00 3.62664 4 10.53 NGC AB 5218 2.93 11.61 UGC665 08662 3 10.86 B 0.00 0.03 0.36 3.63 2.45 UGC676 08733 0 2.20 4.07 A 0.15 2.63684 10.06 3.74 1 0.06 0.29 B IC 4 0944 UGC708 A 10.61 08778 0.90 0.01 9.85 0.74 2.59 5 UGC 3.65714 B 10.85 08781 1.09 2.77 9.10 3.99715 0.12 NGC 2.22 5378 8.97 0.34 9.04 2 9.57 3.98 0.59740 A NGC 0.44 1 5406 A 9.50 0.02 2 2.38 2.46744 B NGC 5485 10.35 0.59 2.68 UGC 2.30748 9.46 0.18 11.41 09067 9.30 2 B 0.00 8.90 11.29 2.80749 9.60 NGC 0.17 2 5520 3.63 0.20 B 8.36 1.64 4.01 8.50754 –1 10.72 NGC 0.12 A 5614 4.03 3 AB 1.62758 9.43 0.16 9.49 11.28 NGC 1.63 5631 10.04 10.86 3.96764 0.01 11.02 0.16 NGC 0.31 3 5633 2.60 9.79 9.38 A 2.82 4.17768 9.08 0.37 NGC 0.47 1 5630 9.28 4.36 2.49 A –1.17 3.62 –0.17769 0.14 10.15 NGC 9.86 0 5657 A 0.55 10.09 2.46774 0.22 11.40 0.05 NGC 0.46 3 5682 8.98 A 9.86 0.12 4.19 2.69 0.15775 10.76 NGC 9.45 5 5720 3.14 B 0.21 0.03 2.32 8.80 4.63 –0.61778 0.35 10.50 8.61 –0.52 NGC –0.42 0.03 3 5732 0.01 B 0.08 4.58 UGC779 09476 8.58 0.25 4 9.86 9.22 B 3026 0.03 2.58 3.74 UGC 0.01780 9.63 0.03 10.66 09537 0.71 3 –0.21 B 9.70 0.17 3114 2556 2454 3.15 UGC783 0.18 9.87 09542 –0.12 –0.48 3 9.50 2.77 3.0 10.02 9.22 A 0.01 2396 –0.18 1891 3 3.18 4.02 0.20789 A –0.05 11.24 NGC 0.82 5784 2.3 8.55 2.5 3939 2 2.87 UGC 10.04791 A 10.24 09598 0.48 2.47 –0.10 8.84 10.42 3398 –0.26 5940 5609 4 3.79795 0.24 A 0.02 NGC 9.89 0.04 2.08 5797 3.3 10.01 0.01 11.38 5862 4242 5107 3.33 2.60 UGC797 9.39 0.21 9.39 –0.07 09665 0.54 0 –0.20 3.24 9.63 2.2 2.3 A 3 9.53 10.58 4611 AB798 0.05 NGC 2680 –0.93 1.90 5888 2.4 4.07 0.41 –0.14 2.40804 0.44 10.75 9.62 –1 11.32 9.33 NGC 2019 6866 1.08 A 5908 –0.08 3.38 6011 2 2.07 0.06 2.3 806 A 9.92 NGC 6033 5930 2.16 9.30 0.25 9.79 10.92 3467 0.47 –0.03 3.24 4.38 –0.14 UGC 3.0 807 0.37 9.90 09873 0.29 2 1814 2.5 9.13 2.53 B 0.03 10.20 0.11 –0.47 UGC809 9.41 09892 0.59 1447 1 2.20 3.87 A 0.15 2207 –0.24 0.39 –0.02 1 2.4 0.26 11.55 NGC AB 3330 5971 3.70 8.89 8.38 1801 0.85 2.38 2 2.94 0.14 A 11.25 NGC 9.44 –0.07 2285 2676 5966 3.0 9.55 10.75 0.09 0.19 0.10 1.03 3 3.92 2.5 A 2274 3345 0.03 2.89 9.78 0.13 0.01 0.27 8.70 10.28 IC 4566 4.28 10.06 2.2 2 NGC 0.06 AB 0.19 4.28 2884 5987 2.56 8.86 0.07 0.32 10.52 2850 2.1 –1 9.14 A –0.03 3.20 10.57 9.31 2227 0.00 2.63 8.50 0.30 0.13 6635 7114 3.41 0.07 9.10 0.09 2.4 0.17 11.15 9.79 0.95 3.11 –0.35 5305 2.77 3.97 5011 1 2 3334 A B 0.16 0.00 9.53 2.8 8.86 2.5 9.40 2769 0.28 –0.10 1.89 4.16 0.64 –0.20 2.4 9.74 11.13 9.78 11.12 5563 0.74 1914 2.25 8.98 1.36 6851 –0.26 8.98 0.17 4457 1558 2.45 –0.14 0.12 4124 9.23 1.01 4.54 2.3 3.94 2.8 9.49 0.37 –0.05 1589 2.4 2.92 0.11 9.61 –0.23 0.55 –1.14 –0.06 9.26 839 0.31 0.46 10.00 3119 –0.65 –0.13 9.02 3.0 9.57 0.36 1498 –0.12 0.14 3.20 –0.20 2.56 9.88 2495 2214 0.12 9.31 1360 3.1 9.56 0.04 –0.70 1837 –0.04 3.2 0.00 9.24 3530 –0.39 2.1 9.01 0.52 0.42 3053 9.59 –0.11 3056 9.48 –0.05 –0.40 1902 2.5 9.96 0.20 2354 7115 3.1 0.49 0.29 10.04 9.76 –0.18 2147 5424 9.58 0.20 –0.79 3362 2.6 2.4 4273 –0.36 9.07 2412 –0.02 3748 2.3 9.08 9.84 9.88 0.09 6566 0.04 1.8 5811 9.54 4555 9.74 0.14 –0.20 4418 0.01 3.3 5860 –0.35 2.3 –0.44 3589 4847 –0.12 2.2 0.03 2659 9.78 9.43 –0.40 2663 2437 3.4 –0.11 –0.10 0.08 2176 1619 –0.46 3.1 2.4 9098 –0.34 2436 7510 4099 –0.46 –0.05 2.1 0.10 1725 3172 5413 2.8 2.7 0.11 3943 6166 2.4 2983 4541 2.5 1987 0.06 –0.00 3597 3.0 2943 2.8 2737 5242 2392 3897 3.0 2.2 ID# NED morph. Table C.2.

A103, page 41 of 44 A&A 581, A103 (2015) C HLR HMR M i ) pc pc

Z [HLR] Z log h M i )( [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) ( A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA810 Name813814815816 NGC Type 5980818 NGC log 6004 UGC820 10097821 NGC 3 6020 A822 NGC 3 6021 B –1 A UGC823 10.98 10123824 –1 10.88 NGC 11.49 A 6032 3.87826 –1 NGC A 6060 11.08 1 UGC 3.82 4.63827 A 10205 11.08828 NGC 3 6063 3.92 B 10.65 2.81829 2 4.56 A IC 0 1199 2.42 3.03831 A 10.63 NGC 6081 3.86 UGC832 11.15 10297 0.74 3 2.97 A 11.24 UGC 3.72833 10331 0.65 0.18 2.90 4.09834 2 0.49 10.29 NGC AB 0 6125 2.75 4.08 A 0.00 4835 A 0.25 0.00 NGC 6132 10.93 4 AB 0.73 3.16 2.27836 11.22 NGC 6146 0.01 1.13 9.58 2.73837 10.23 –1 9.66 NGC 3.84 A 6154 2.56 0.01 4.47 UGC 10.10838 10380 9.41 0.80 3 A 0.61 11.41 3.02 2.26 2.96840 –1 NGC 0.71 A 6150 9.59 1.50841 0.16 10.46 NGC 2.74 1 6155 10.04 4.43 B 2 11.73 AB 2.94 UGC842 0.50 10384 0.25 9.07 0.25 9.55 2.13 2.23 3.47 UGC843 11.13 –1 9.75 11.19 10388 A 4.37 9.12 0.76844 0.21 NGC 9.55 0.93 4 6173 3.06 A 11.48 3.65 2 4.02845 9.82 A 9.74 NGC 0.60 0.53 6168 0.34 1 9.81 9.59 AB 2.16846 0.16 10.32 NGC –0.05 6186 2.96 3.95 10.63 0.00 UGC 0.22848 9.38 0.19 10.71 –1 9.64 10650 0.12 A 0.11 2.61 3.52 2.42849 4 NGC AB 0.50 6278 4.01 9.67 9.45 0.00 11.86 0.21 4.00 UGC850 9.98 0.22 10693 2 2.87 –0.08 9.21 B 9.88 –0.13 4 UGC851 9.51 0.51 9.23 A 9.32 10695 0.72 0.40 4.31 0.00 2.46 –0.02852 0 0.18 10.75 NGC AB –0.07 –1 6310 2.68 10.03 9.16 AB 3.10 0.00 2.62 9.47 0.37854 0.06 NGC 4507 6301 –0.19 10.89 11.57 –1 –0.01 A 9.16 4.10 0.01856 9.55 NGC 3045 0.73 6314 2.97 9.80 0.07 0.00 9.99 4808 1.60 3473 2.71 –0.11 2.3 857 NGC 11.36 0.27 4.61 –0.24 2 6338 4.15 9.26 9.01 A 2.18 4213 2410 UGC 10.01858 0.47 9.86 10796 3 –0.14 2679 2.0 3.3 A 9.85 0.90 2765 –0.38 0.22 2.72 UGC 3.96 0.07859 10.70 10811 1 1683 10.08 A 2312 3266 1.57 0.01860 3.3 –1 11.35 0.84 –0.22 3.07 –0.22 A 2.75 3.0 8.61 4 0.04 0.00 IC AB 2206 1256 3.91861 9.80 11.21 9.12 NGC –0.67 0.86 –0.68 5871 6394 2.5 11.72 9.40 –0.32 10.03 2 3.24 0.44 UGC 2.54862 9.67 B 10905 0.44 9.55 4509 9.45 0.07 0.18 5662 0.02 5579 4.09863 2.0 0.29 NGC 6411 10.01 4.28 9.80 –0.01 11.11 4337 3500 2.78864 2 0.28 NGC AB 2.57 3636 3 –0.02 0.30 6427 0.01 2.2 0.00 2.5 0.00 B –1.23 –0.41 0 2.39 UGC 9.06865 A 0.08 10972 2902 8.88 3.71 10.71 2.70 9.64 1.9 866 8.82 –1 0.01 11.14 NGC 4917 9.35 0.06 0.46 –0.10 A 6478 2.87 0.10 11.55 0 3834 8.46 NGC 3500 AB 4180 0.33 1.62 10.01 3.50 2654 6497 9.94 9.85 11.07 0.06 –0.00 2.1 3 3.47 0.29 A NGC 2342 3470 2328 0.42 6515 4.19 2.45 10.84 0.36 3.4 UGC 2.4 0.01 2.4 10.07 0.27 11228 3259 4 8.88 –0.07 3.99 A 10.64 0.36 –0.21 0.00 UGC 0.03 0.37 0.03 11262 2775 2.41 4.59 4043 1 9.40 B 0.01 0.27 2.67 3.2 –1 11.25 9.93 2356 A 9.74 2.92 3.52 0.31 8.58 0.14 0.11 5577 0 9.69 2.5 B 9.61 11.13 2.91 0.30 11.34 0.72 6828 4509 –0.18 6057 4 3.86 A 9.56 9.70 –0.80 3.16 –0.58 0.45 3.4 10.01 5439 –0.09 11.17 4066 0.26 3.77 5798 2.8 –0.01 2.35 4.27 0.21 2.3 8.17 10.39 0.00 0.34 5082 2491 –1.22 9.45 4.35 0.14 0.01 0.28 0.23 9.94 4188 2.58 2.9 0.19 3068 2158 3.07 9.15 0.00 –0.57 2380 0.35 2.60 2.2 2280 0.07 2.64 0.00 9.40 9.33 2.6 9.84 –0.03 3.1 9.84 0.66 6768 2.72 0.12 9.78 –0.69 8.65 3244 6044 0.00 –0.00 0.18 1.66 9.56 0.14 0.07 3.1 0.35 9.41 2549 3024 9.92 0.14 2.0 0.17 0.05 0.16 2252 3578 0.09 0.00 9.91 0.11 0.47 9.19 2.4 1868 2917 9.35 0.00 6271 –0.18 0.00 2.5 9.80 1491 9.60 0.26 5659 3.4 6736 –0.10 9.59 9.73 –0.06 3.3 9.91 4807 9.85 0.02 10.00 0.16 2.8 3861 0.16 9.21 –0.46 –0.12 10836 9.56 3313 0.08 9716 –0.15 2.3 9.05 2.0 0.22 4807 6892 9.41 3325 0.19 9.65 3807 6039 –0.09 9.73 3673 –0.11 3.3 –0.15 7769 2.8 2.4 –0.04 6024 8.79 0.10 2.7 0.14 4814 6792 0.12 0.03 3805 0.10 –0.20 5471 5489 0.17 2.2 2.1 4385 3196 –0.48 3.5 –0.21 1682 2925 4887 3.0 0.04 1386 3704 0.01 3.3 2.3 0.12 7164 –0.76 5127 5308 2.0 4903 4345 3770 4519 6216 2.7 2876 3.1 6729 3.4 2.2 ID# NED morph. Table C.2.

A103, page 42 of 44 R. M. González Delgado et al.: The CALIFA survey across the Hubble sequence C HLR HMR M i ) pc pc

Z [HLR] Z log h M i )( [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) ( A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA867 Name868869870871 MCG-02-51-004 NGC Type 6762872 log 2873 A NGC 6941874 NGC 1 6945 A 10.92875 NGC 6978 UGC 11680NED 01 UGC876 10.38 11649 2 3.85 B877 2 0 B B 3.78878 2 11.22 NGC AB 7025 1 UGC 11.32879 B 10.92 11694 2.48 11.15 3.84880 NGC 7047 10.67 3.93 3.00 UGC 4.49881 MCG-01-54-016 11717 3.98 0 A 0.46 0883 A 3.85 4 2.52 UGC884 11.38 A 11740 0.00 3 B 0.30 11.20 2.70 1 2.94 UGC885 A 11792 2.69 4.62 9.25886 0.16 10.94 NGC 0.21 7194 4.53 2.44 11.06 3 UGC887 A 0.59 11958 0.52 0.72 9.93 3 UGC 3.48 2.07888 A 0.24 11982 3.98 10.60 0.18 3.14 0.23 UGC889 –1 0.00 9.68 12054 A 2.77 0.20 10.56 0890 A NGC 7311 3.03 0.11 11.53 4 2.51 1.69891 A NGC 9.79 0.35 7321 2.16 3.63 11.44 0.84 4 UGC 9.90892 9.07 A 12127 9.79 4.47 9.74 10.01893 0.06 NGC 0.15 0.06 1 7364 2.30 3.92 9.65 A 0.14 9.93 0.63 9.18 UGC894 12185 3 2.66 2.84 B –1 A UGC895 0.13 0.21 11.31 12224 VV 488NED 2.93 9.41 0.12 02 0.33 10.01 3.12 –0.10896 11.28 11.62 1 9.41 2.49 10.02 9.60 A 1.20 NGC 9.39 2 4.61898 7436B B 2 2.08 AB 0.07 0.41 0.11 4 UGC 3.89 4.18900 9.39 A 10.98 12274 9.88 8.94 0.48 9.91 10.91 10.61 0.22 2.07 UGC901 12308 –0.28 –1 0.01 0.12 10.20 A 0.51 3.19 4.11902 9.85 NGC –0.15 7466 3.92 9.65 3.79 0.03 0.14 9.74 –0.12 1 2.59 2.79903 11.68 A NGC –0.03 0.50 7489 3.08 0.46 9.44 4 5738 –0.07904 A NGC 0.36 7550 10.04 9.31 8.83 11.19 4.34 3735 2.84906 9.41 0.38 NGC 0.36 3 0.01 7549 10.08 0.08 2.60 A 2.63 2.2 1581 0.03 9.12907 0.27 0.16 –0.01 NGC 3 7563 1.89 3.93 0.09 A 9.30 0.04 1359909 8.80 –1 0.27 11.05 0.00 NGC 0.16 A 7562 2.9 0.04 2.86 0.23 2.83 0.93910 9.31 10.85 9.31 NGC 6592 3 7591 –1.18 B 9.75 5518 0.02 11.36 5325 0.34 3.97911 0.28 9.89 0.08 5207 2442 1 2.59 4081 B 9.84 0.13 IC 4111 0.45 5309 3.53 2.1 0.08 10.08912 –1 10.78 NGC 9.78 1949 0.11 2.5 A 3470 7608 2.4 4.36 0.04 0.01 1.66 –0.21 UGC 3.0 914 8.88 11.09 12519 2449 3 B 11.30 –0.93 0.25 2.41 –0.12 UGC 3.81 0.05915 2.1 9.56 12518 –0.07 8.83 0.07 9.84 3166 9.67 2.09 4.47 0.03916 4 11.03 NGC AB 0.26 3689 3 7619 2.93 4.65 9.51 A 4 0.03 9.52 0.15 2621 AB917 9.77 NGC 2043 1.00 –0.01 2811 7623 3.3 10.46 9.23 6172 2 –0.60 2.36 4.04 2.9 919 0.03 10.40 9.98 A 10.24 NGC 0.86 –0.30 2758 7631 5106 6527 0.32 3.09920 2.3 0.44 –1 NGC 3.61 A –0.70 7653 3.31 10.09 9.12 2.0 0.16 10.23 3102 3.43 5392 3.25 0.22 9.45 NGC 9.15 0.84 0 7671 2.9 0.08 A 0.11 0.00 11.43 4369 4178 2.76 9.04 9.23 NGC 0.06 –0.17 0.38 2 7683 3.78 2.4 A 0.12 3216 0.20 0.45 10.84 9.37 NGC 2.20 4377 2 7684 9.80 4.69 2.4 A –0.54 2.28 2.41 0.02 10.71 9.11 NGC 3412 7371 1.43 0 7691 A 9.73 0.14 0.00 9.84 4393 4.24 –0.06 3.3 –0.36 10.75 6319 0 0.07 2.73 A 0.84 3605 3.88 –0.37 2.8 0.07 0.70 11.10 9.34 0.76 –0.06 0.60 1277 0 3.37 8.54 2.3 A 3.97 11.15 9.98 1006 3 10.06 0.43 9.16 0.14 B 0.85 2.90 4.73 2.7 0.38 0.54 11.02 3243 –0.20 0.03 –0.03 –0.13 8.68 6168 0.05 5868 2.73 4.58 9.58 10.37 2430 9.39 –0.16 0.43 4331 5296 2.41 3.1 4.39 0.17 –0.92 2.4 9.28 8.77 3.1 0.00 3.17 4498 3.00 3545 9.35 9.47 3126 0.39 0.05 9.90 9.79 0.04 3199 4171 3.23 2525 0.01 0.09 2132 0.37 –0.06 2.8 9.93 3.1 3338 2.93 3.1 0.24 0.08 0.28 9.10 1.8 9.98 –0.71 1.98 0.26 5955 0.76 –0.05 6073 9.06 0.04 5822 9.96 –0.39 0.38 8.79 9.03 0.14 0.09 –0.22 4522 2.9 0.09 9.84 1982 0.28 2.7 0.19 9.64 0.08 9.67 1439 –0.61 6059 –0.10 10.08 2.5 9.81 0.04 6686 4081 –0.15 9.67 3974 –0.07 –0.42 2.7 9.76 0.17 0.19 3694 2.3 9.91 4078 9.33 2866 –0.75 –0.09 9.37 2965 2.9 9.00 9.84 2.1 0.18 –0.21 2308 2616 –0.19 –0.20 1969 10.00 2439 –0.00 3745 3.5 3.1 2632 –0.11 9.74 0.02 3694 3.0 4521 3441 –0.08 8.96 2196 0.19 0.17 3147 2889 2.4 2.3 0.05 2.3 0.18 2074 –0.01 1660 0.22 –0.22 2564 2.8 –0.31 2192 2126 0.08 3602 3.3 1659 0.07 2780 3365 3.4 2.6 –0.02 2035 2140 –0.51 2.7 1562 2273 3.2 1875 3290 4901 3.2 2442 4184 3.1 2.0 ID# NED morph. Table C.2.

A103, page 43 of 44 A&A 581, A103 (2015) C HLR HMR M i ) pc pc

Z [HLR] Z log h M i )( [0]

Z Z log h L i log age[HLR] h L i log age[0] h [HLR] V A [0] V mag mag (yr) (yr) ( A 2 pc [HLR] / ?

µ M 2 [0] log ? pc / µ

M log ?

M M continued. CALIFA921 Name923924925 UGC926 12653 Type927 NGC log 7711928 NGC 7716 4929 A NGC 7722 UGC930 –1 12723 A 9.89931 NGC 2 7738 A 10.98 UGC932 12767 1 A 4 2.81 UGC933 A 10.52 12810 4.49 UGC934 11.19 12816 2 B 10.07 NGC 2 7783NED 4.25935 B 01 NGC 7782 1.71 3 4.32936 B 11.34 3.03 2.67 1 11.04 A 4937 A NGC 7787 11.00 2.72 UGC 4.38938 12857 0.00 2 11.47 3.82 A 10.12 0.33 3.05 UGC939 MCG-01-01-012 12864 2.27 3.72 1 0.15 11.42 3.87 AB 0.02 3.02 0.01 1 AB 3 2.49 A NGC 0.66 7800 2.41 10.85 0.71 4 4.17 0.07 10.88 B NGC 5947 2.15 NGC 9.78 0.53 4676B 9.53 3.05 3.83 1.99 1.69 0.24 9.98 10.22 0.03 3.67 5 AB 1.25 2.81 3.14 0.08 9.90 3 3.21 0.00 B 0.10 9.31 1 0.26 B 9.78 2.29 0.16 9.07 2.67 10.81 0.43 0.01 8.57 11.11 0.26 2.64 2.33 9.63 8.98 1.86 1.58 9.98 3.77 0.13 0.80 3.98 9.54 9.32 0.61 9.56 0.27 9.63 9.29 1.99 0.44 9.36 0.20 –0.25 10.09 2.10 0.39 0.16 2.80 9.53 0.18 9.20 0.34 9.44 0.15 9.08 9.95 0.45 0.12 0.40 9.86 –0.16 –1.07 0.38 8.68 9.17 8.94 0.08 9.65 0.14 –0.10 0.19 0.02 –0.20 4547 9.31 –0.29 8.12 3468 –1.37 0.12 9.65 2554 0.22 –0.50 2.3 9.46 8.99 9.94 1698 1943 0.10 8.82 0.09 3.3 1238 5140 0.19 3097 –0.79 2.7 –0.24 4808 2555 –1.00 8.61 –0.08 2.1 0.09 6037 2.6 8.96 0.07 2750 5670 8345 9.56 –0.37 0.06 3.0 4376 5480 4353 –0.01 5173 2.4 2.0 2984 0.02 –0.34 4350 2.4 6197 3.0 –0.29 –0.50 0.18 4869 4866 0.01 2.2 2635 4977 2.6 –0.69 4080 2041 5227 2.6 1602 3944 –0.03 2.5 2.4 0.03 1598 1632 5225 2.4 3912 3149 2.4 2902 2.4 ID# NED morph. Table C.2.

A103, page 44 of 44