Astronomy & Astrophysics manuscript no. SFR_CALIFA c ESO 2018 November 7, 2018
Star Formation in the Local Universe from the CALIFA sample I. Calibrating the SFR using IFS data
C. Catalán-Torrecilla1, A. Gil de Paz1, A. Castillo-Morales1, J. Iglesias-Páramo2, 3, S.F. Sánchez4, R. C. Kennicutt5, P.G. Pérez-González1, R. A. Marino6, C.J. Walcher7, B. Husemann8, R. García-Benito2, D. Mast9, R. M. González Delgado2, J.C. Muñoz-Mateos10, J. Bland-Hawthorn11, D. J. Bomans12, A. del Olmo2, L. Galbany13, 14, J.M. Gomes15, C. Kehrig2, Á. R. López-Sánchez16, 17, M. A. Mendoza2, A. Monreal-Ibero18, M. Pérez-Torres2, 19, P. Sánchez-Blázquez20, J. M. Vilchez2, and the CALIFA collaboration
(Affiliations can be found after the references)
ABSTRACT
Context. The Star Formation Rate (SFR) is one of the main parameters used to analyze the evolution of galaxies through time. The need for recovering the light reprocessed by dust commonly requires the use of low spatial resolution far-infrared data. Recombination-line luminosities provide an alternative, although uncertain dust-extinction corrections based on narrow-band imaging or long-slit spectroscopy have traditionally posed a limit to their applicability. Integral Field Spectroscopy (IFS) is clearly the way to overcome such limitation. Aims. We obtain integrated Hα, ultraviolet (UV) and infrared (IR)-based SFR measurements for 272 galaxies from the CALIFA survey at 0.005 < z < 0.03 using single-band and hybrid tracers. We aim to determine whether the extinction-corrected Hα luminosities provide a good measure of the SFR and to shed light on the origin of the discrepancies between tracers. Updated calibrations referred to Hα are provided. The well-defined selection criteria and the large statistics allow us to carry out this analysis globally and split by properties, including stellar mass and morphological type. Methods. We derive integrated extinction-corrected Hα fluxes from CALIFA, UV surface and asymptotic photometry from GALEX and integrated WISE 22µm and IRAS fluxes. Results. We find that the extinction-corrected Hα luminosity agrees with the hybrid updated SFR estimators based on either UV −1 or Hα plus IR luminosity over the full range of SFRs (0.03-20 M yr ). The coefficient that weights the amount of energy pro- duced by newly-born stars that is reprocessed by dust on the hybrid tracers, aIR, shows a large dispersion. However, it does not became increasingly small at high attenuations, as expected if significant highly-obscured Hα emission would be missed, i.e. after a Balmer decrement-based attenuation correction is applied. Lenticulars, early-type spirals and type-2 AGN host galaxies show smaller coefficients due to the contribution of optical photons and AGN to dust heating. Conclusions. In the Local Universe the Hα luminosity derived from IFS observations can be used to measure SFR, at least in statistically-significant, optically-selected galaxy samples, once stellar continuum absorption and dust attenuation effects are ac- counted for. The analysis of the SFR calibrations by galaxies properties could be potentially used by other works to study the impact of different selection criteria in the SFR values derived and to disentangle selection effects from other physically motivated differences, such as environmental or evolutionary effects.
Key words. galaxies: evolution, galaxies: spiral, galaxies: star formation, techniques: photometric, techniques: spectroscopic
1. Introduction Until now, the study of the evolution of the SFR has focussed on the analysis of the integrated SFR in galaxies, with little at- The measurement of the star formation rate (SFR) is crucial for tention being paid to where in galaxies (nuclei, bulges, disks) understanding the birth and evolution of the galaxies (Kennicutt SFR takes places and how the SFR in each of these components 1998a) as it provides information on the amount of gas in galax- evolves separately with redshift. It is remarkable that the use ies and the efficiency in the formation of stars inside them, which of NIR integral field spectroscopy on 8-10 m class telescopes is depends strongly on the conditions of the interstellar medium arXiv:1507.03801v1 [astro-ph.GA] 14 Jul 2015 now allowing us to measure the SFR in these different compo- in which they are formed (Kennicutt & Evans 2012, and ref- nents in distant galaxies, up to z=1-3 (e.g., Genzel et al. 2008; erences therein). The SFR is, together with galaxy mass, one Förster Schreiber et al. 2009, 2011a,b; Nelson et al. 2012, 2013; of the most important parameters that define galaxies and their Wuyts et al. 2013; Lang et al. 2014) while the local benchmark evolution across cosmic times (Somerville & Davé 2014; Madau for these and possible future studies is still missing except for & Dickinson 2014). Several authors have tried to quantify the a few studies rather limited in number and completeness (e.g., rate of on-going star formation and its evolution with redshift Pérez-González et al. 2006; Kennicutt et al. 2007; Leroy et al. (e.g., Madau et al. 1996; Lilly et al. 1996; Pérez-González et al. 2008; Bigiel et al. 2008; Blanc et al. 2009; Schruba et al. 2011; 2008; Bouwens et al. 2007, 2011, 2014) using different tracers. Leroy et al. 2012). In this regard, a correct determination of These works have shown that the SFR density has declined by the calibrators we use to calculate the spatially-resolved SFR is roughly a factor of six from z=2 to present day (Hopkins & Bea- essential in order to compare how the star-formation of these dif- com 2006).
Article number, page 1 of 36 ferent spatial components behave at different wavelength ranges balance approach (see Gordon et al. 2000; Inoue et al. 2001; Hi- and/or redshifts. rashita et al. 2003; Iglesias-Páramo et al. 2006; Calzetti et al. Although SFR calibrators have existed for almost 30 years, 2007; Kennicutt et al. 2007, 2009; Hao et al. 2011, for more de- the last decade has been particularly fruitful thanks to the multi- tails) wavelength surveys of nearby and distant galaxies. The devel- In this work, we derive integrated extinction-corrected Hα- opment of the Integral Field Spectroscopy (IFS) technique has based SFRs from the analysis of CALIFA IFS data and com- allowed us to combine the advantages of both imaging and spec- pare them with measurements from other SFR tracers. We pro- troscopy at optical and near-infrared (NIR) wavelengths. In this vide new single-band and hybrid updated SFRs tracers (with paper we make use of a large and well-characterized sample and without type-2 AGN being considered) using our integrated of nearby galaxies from the Calar Alto Legacy Integral Field extinction-corrected Hα SFR as a reference, thanks to the qual- Area (CALIFA) survey (Sánchez et al. 2012) that spans the en- ity of our attenuation correction via Balmer decrement. We pay tire color-magnitude diagram to address this fundamental issue. special attention on the hybrids ones, providing for the first time, CALIFA allows us to properly determine the Hα and Hβ fluxes a set of hybrid calibrations for different morphological types and using IFS spectroscopic data. This is particularly important in stellar masses. We also analyze the dependence with the color the case of galaxies with low equivalent widths in emission, spe- (SDSS g − r), axial ratio and ionized-gas attenuation. This anal- cially in Hβ, like many of the objects in the CALIFA sample and ysis is the starting point for a series of papers in which we will in the Local Universe in general (Gallego et al. 1995; Brinch- study how the SFR in the Local Universe is distributed across mann et al. 2004), where narrow-band imaging is not feasible. galaxy components (bulge, disks, nuclei) and bidimensionally. Furthermore, using these data we can separate the Hα and [NII] Ultimately, we are interested in knowing how the local SFR den- flux while in narrow band imaging is only feasible if a [NII]/Hα sity is spatially distributed over galaxies and how these results ratio is assumed. In this paper, the use of the uniqueness IFS data would compare to similar future studies at high redshift. will allow us to obtain precise Balmer-decrement measurements This paper is organized as follows: in Section 2 we describe to compute Hα extinction-corrected luminosities. Various stud- the reference sample used in this article, in Section 3 we describe ies have shown the importance of computing the extinction using the data and the analysis apply to the data, in Section 4 we dis- IFS data in nearby regions where the line ratios obtained from cuss our results, and finally, in Section 5 we summarize the main the integrated spectra are dominated by regions of lower surface conclusions. Throughout this paper we use a cosmology defined −1 −1 brightness rather than by the brighter ones (Pellegrini et al. 2010; by H0 = 70 km s Mpc , ΩΛ = 0.7 and a flat Universe. Relaño et al. 2010; Monreal-Ibero et al. 2011). Other advantage related with the IFS data is that we can cover the whole galaxy avoiding problems associated with the limited spatial coverage 2. The Sample of long-slit spectroscopy. From these Balmer-corrected Hα lu- minosities, we compute their corresponding SFRs which will be The galaxies studied in this work are part of the CALIFA survey use as a fiducial measure of the current SFR. However, it is crit- (Sánchez et al. 2012). The CALIFA mother sample includes 939 ical to first determine that at least, in a statistically sense, no galaxies of all types. A total of ∼600 galaxies will be observed as significant fraction of the SFR is being missed when using the part of CALIFA using the Potsdam Multi-Aperture Spectropho- extinction-corrected Hα luminosity as SFR estimator. This re- tometer (PMAS, Roth et al. 2005) in the PPak mode (Kelz et al. quires of a combined analysis of this estimator with other SFR 2006) mounted at the Calar Alto 3.5 m telescope. The survey estimators, including the continuum ultraviolet (UV) emission, is described in detail in the presentation article (Sánchez et al. recombination lines of hydrogen and other atomic species to- 2012). As a summary, the CALIFA mother sample (Walcher gether with other estimators less affected by dust attenuation et al. 2014) includes all galaxies in the DR7 SDSS photomet- ◦ such as total infrared (TIR) luminosity, monochromatic infrared ric catalog (Abazajian et al. 2009) with declinations above 7 , (IR) emission or radio emission. The combination of different spectroscopic redshifts (from the SDSS spectroscopic catalog SFR estimators is also needed to evaluate the potential differ- or elsewhere) in the range 0.005 Article number, page 2 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample 350 300 300 whole sample whole sample whole sample this work this work this work 300 250 250 250 200 200 200 150 150 150 100 100 Galaxy number 100 Galaxy number Galaxy number 50 50 50 0 0 0 23.5 22.5 21.5 20.5 19.5 18.5 17.5 16.5 1 3 5 7 9 11 13 15 17 19 8.2 8.6 9.0 9.4 9.8 10.2 10.6 11.0 11.4 11.8 Mr Half-light radii (kpc) log(M /M ) ¯ Fig. 1: From left to right. Distributions of the r-band absolute magnitude, half light radius and log(M?/M ). The white histograms correspond to the complete CALIFA sample (939 galaxies), and the grey-filled areas correspond to our sample (380 galaxies). A visual inspection of these histograms together with the perform of the K-S test probability show that our sample is representative in terms of galaxy properties of the entire CALIFA sample. (white areas in Figure 1) and the galaxies involved in this study minor-to-major axis ratio listed in NED1 for each object. The (grey-filled areas). We use a K-S test to check whether the two previous values and the position angle (PA) are measured at the data samples come from the same distribution. The K−S test 25.0 mag/arcsec2 isophote at B-band provided by the RC3 cat- probability is computed using the limits where the mother sam- alogue (de Vaucouleurs et al. 1991). When this information is ple is representative of the general galaxy population as men- not available we use the SDSS g or r-band isophotal photom- tioned before. The values of the probabilities found by the K−S etry. As the extracted aperture is significantly larger (∼4000×) test are 40.25% in half light radius, 70.95% in r-band absolute than the CALIFA pixel size (1 arcsec2), effects associated with magnitude and, finally, 75.55% in log(M?/M ). From these val- the treatment of fractional pixels are negligible. ues and from the visual inspection in Figure 1, we conclude that the subsample we are using is representative of the mother sam- ple, except for a marginal deficiency of intermediate luminos- 3.1.3. Continuum subtraction and line-flux measurements ity objects in the range Mr=(−20.5,−21.2) that might explain the low K−S values but that certainly does not bias the results In order to minimize systematics associated to the stellar contin- against these systems. uum subtraction at low-S/N regimes, we decide to first spatially integrate the datacube within these apertures. This is a partic- ularly interesting use of the IFS data that allows both covering 3. Data and Analysis the whole galaxy and having a high-S/N in the integrated spec- trum. Then, we carry out the necessary corrections to derive total 3.1. CALIFA Integral-Field Spectroscopy extinction-corrected Hα luminosities. The use of the Hα/Hβ ra- 3.1.1. CALIFA Survey tio derived from the integrated spectra is justified instead of cor- recting for extinction spaxel to spaxel and then coadding the flux The CALIFA spectra cover the range 3650-7500 Å in two over- in order to minimize systematics when adding up signal from lapping setups, one in the red (3745-7500 Å) at a spectral resolu- very noisy individual spaxels, as is shown below. Thus, while tion of R∼850 (V500 setup) and one in the blue (3650-4840 Å) at Sánchez et al. (2011) and Marino et al. (2012) show that the R∼1650 (V1200 setup), where the resolutions quoted are those at dust attenuation from individual spaxels is a little larger than the the overlapping wavelength range (λ∼4500 Å). For the purpose one derived from the integrated spectra (1.24/1.04 and 1.19/1.03, of deriving extinction-corrected Hα luminosities we make use of respectively), Castillo-Morales et al. (2011) obtain rather simi- the V500 setup as we are interested in having both Hβ λ4861 Å lar values in each case. Nevertheless, the interesting point here is how the luminosity-weighted attenuation compares when us- and Hα λ6563 Å emission lines in the same observing range. ing individual spaxels with the one from the integrated spectra. ∼ The spectral resolution (FWHM 6 Å) is sufficient to deblend This question is more relevant as we are analyzing attenuation- the Hα emission line from the nearest [NII]λλ6548,6584 Å dou- corrected Hα luminosities rather than attenuations themselves. blet lines. We are using the v1.3c data products which yield the For this matter, we select the galaxy NGC 5668 in Marino et al. measured flux densities corrected for Galactic extinction. The (2012) as it is a nearby spiral galaxy similar to the ones used in data reduction is explained in detail in Sánchez et al. (2012) and this work. We find that the difference between computing the Husemann et al. (2013). average luminosity-weighted attenuation from individual spax- els and that derived from the integrated spectrum is less than 1%. 3.1.2. Aperture spectrophotometry From this result, we conclude that we can safely use the Hα/Hβ ratio derived from the integrated spectra to correct the Hα flux For each galaxy for which the CALIFA V500 observations in each galaxy. Besides, this way of obtaining the extinction- reached the full depth planned (3×900 seconds exposures in a corrected Hα luminosity would actually mimic what one could three-point dithered scheme), we generate an integrated spec- measure in more distant systems for which this work is intended trum within the largest common aperture possible between the to provide a local benchmark. CALIFA and the other complementary data (UV, IR). This aper- ture has an elliptical shape with a major axis radius of 36 arcsec and the corresponding ellipticity of the galaxy, as given by the 1 https://ned.ipac.caltech.edu Article number, page 3 of 36 The first of those corrections to be applied to our data is to carefully remove the stellar continuum underlying the Hβ and Hα lines. This is done by means of adjusting a linear combina- tion of two single stellar population (SSP) evolutionary synthe- sis models of Vazdekis et al. (2010) based on the MILES stel- lar library (Sánchez-Blázquez et al. 2006) to the spectrum ob- tained for each aperture. Two set of models with a Kroupa IMF (Kroupa 2001) are combined. One set contains models (consid- ered as a young stellar population) with ages of 0.10, 0.50 and 0.79 Gyr. A second set (considered as an old stellar population) involves ages of 2.00, 6.31 and 14.13 Gyr. For each age we considered five different metallicities with [M/H] values equal to 0.00, 0.20, −0.40, −0.71 and −1.31 dex offset from the solar value. Different wavelength ranges corresponding to the emission lines from the ionized gas and sky-lines are masked and not in- cluded in the fit. The basic steps applied in this method are the following: (1) shift the SSP templates to match the systemic ve- locity of the integrated spectrum, (2) convolve each stellar pop- ulation model with a Gaussian profile so the absorption features could be broadened to match those of the integrated spectrum, −0.7 (3) redden the spectrum using a k(λ) = RV (λ/5500 Å) power law, where RV = 5.9, as given by Charlot & Fall (2000) (4) Fi- nally, the best linear combination of SSPs is determined by a χ2 minimization. Once we obtain the best underlying continuum of the stel- lar population, we subtract it from the original integrated spec- trum to derive the pure emission line spectrum. The emission line fluxes are computed from this residual spectrum. As some residual continuum could still be present in some cases, we do not simply add all the flux in fixed windows in wavelength. In- stead, we compute the Hβ and Hα emission line fluxes by fitting Gaussian functions plus a low order polynomial function. Fig- ure 2 shows the original integrated spectrum for three galaxies with different levels of emission-line strength, IC 4215, NGC 2906 and NGC 5630 in black. The best fit to the spectrum of the underlying stellar population is shown in red and the emission- line spectrum produced by the ionized gas is shown in blue. Grey-coloured wavelength ranges correspond to the emission- lines and sky-lines masked out in the fitting procedure. A proper estimation of the Hβ emission line flux is crucial to obtain a reliable Balmer decrement and, from it, the correc- tion for extinction of the Hα-based SFR. The method applied here is expected to be a robust procedure as long as a relatively wide wavelength coverage is available (see Mármol-Queraltó Fig. 2: Top panel: The original integrated spectrum for galaxy IC 4215 et al. 2011) and the models contain an extensive range of ages is shown in black, the best fit to the stellar population appears in red and metallicities. However, when the whole spectral range (using the 3745−7500 Å entire spectral range) and the pure emission (3750−7000) Å is used for the stellar continuum fitting we still line spectrum (after the subtraction of a residual continuum shape ap- detect systematic residuals around the Hβ absorption line. The plying a smooth function) is shown in blue. This galaxy shows a small equivalent width in the Hβ line. Center panel: Same as top panel for treatment of these spectral features is particularly critical. They the galaxy NGC 2906. Note that in this case the spectrum shows more could be real, due to the limitation of the models in reproducing prominent Hα and Hβ emission lines. Bottom panel: In this case NGC simultaneously a broad wavelength range and the Hβ region, or 5630 shows a typical emission line dominated spectrum with very high introduced during the data reduction. We have also checked that EWs of Hα and Hβ emission lines. These spectra show the variety in adding an intermediate age population in the linear combination levels of emission line strength in our galaxies. A proper subtraction of of the SSPs does not change the overall results. For that reason, the underlying stellar population is required to obtain the estimation of the stellar continuum fitting around the Hβ line is done for other the emission line fluxes. wavelength ranges using the method explained before. The new spectral ranges used are (3700−5500) Å, (4100−5500) Å and (4800−5500) Å. Given that the residual continuum around Hα and Hγ emission lines does not show systematic uncertainties, pared with the values obtained when different spectral ranges we determine the Hβ flux by anchoring to Hα and Hγ fluxes for the stellar continuum fitting are used. Finally, we obtain based on theoretical line ratios and extinction coefficients. In that the Hβ emission fluxes calculated using the spectral range high S/N spectra, this Hβ emission line flux estimation is com- (4800−5500) Å are in best agreement with the theoretical ones. Article number, page 4 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample We impose a minimum S/N for both Hα and Hβ emission dust screen approximation. Although this could be a possible lines fluxes in order to obtain a precise measurement of the ex- source of systematic error in the analysis, some models have tinction using the Balmer decrement. The S/N emission line es- shown that when applied to normal star-forming galaxies it does timation is done using a formal method calculating the ratio be- not introduce significant systematic errors (Jonsson et al. 2010). tween the Gaussian amplitude at Hβ and the root mean square See a detail discussion about the use of attenuation corrections in the near featureless continuum. A visual inspection of the based on Balmer decrements with a Galactic extinction curve continuum-subtracted spectra at Hβ is performed for all the an- and a foreground screen dust geometry in sections 3.3 and 6.3 in alyzed galaxies and a minimum S/N> 5 is considered for Hβ Kennicutt et al. (2009). For the attenuation correction we use an emission line detection. The number of galaxies with detected intrinsic Balmer ratio of 2.86 for case B recombination (Oster- Hβ emission is 272 over the initial 380 galaxies. This is the sam- brock 1989) at electron temperature Te = 10,000 K and density −3 ple (listed in Table 1 for reference) that will be used in the rest ne = 100 cm (Hummer & Storey 1987) using the following ex- of the analysis. pression 1, where KHα = 2.53 and KHβ = 3.61 are the extinction The spectrophotometric accuracy in CALIFA DR1 was coefficients for the Galactic extinction curve from Cardelli et al. checked using SDSS g and r-band photometry which are both (1989). entirely covered by the V500 setup. Husemann et al. (2013) found a systematic offset of ∆(g −r) = −0.06 mag (median) with K FHα/FHβ a scatter of only 0.05 mag. This means that the spectrophoto- A(Hα) = Hα × log (1) − × − 10 metric accuracy across most of the covered wavelength range is 0.4 (KHα KHβ) 2.86 6% for the CALIFA data. This value is included in our error As an example of how little this attenuation correction would estimation for the emission line fluxes. vary among extinction curves and dust-to-stars geometries, we compare the ratio between the A(Hα) attenuations for the same A(Hβ) using the Cardelli et al. (1989) (R = 3.1) law above and 3.1.4. Flux corrections and uncertainties V the Calzetti et al. (2000) (RV = 4.05) attenuation law. We obtain We also apply some corrections to our data such as aperture cor- A(Hα)Calzetti/A(Hα)Cardelli = 1.03. rections and those associated with the spatial masking of field Note that the standard Hα/Hβ ratio used in equation 1 is only stars and background galaxies applied to the datacubes. We valid for the particular ionization conditions indicated above, but use aperture corrections for the galaxies whose line emission values below 2.86 are also physically possible in HII regions, de- is expected to be more extended that the CALIFA field-of-view pending on the electron density, effective temperature and there- (FOV). Our main criterion is to select the band that would first fore on the chemical abundance. This leads to a number of galax- trace the SFR, and second, that it would do it with the best spa- ies for which we assumed A(Hα) = 0. The computed extinction tial resolution possible. Also, we want to have them available for values A(Hα) are listed in Table 1. The Hα luminosity corrected most galaxies in our sample. Here, we assume that the Hα light by attenuation and by the effects mentioned along this section distribution beyond the FOV is similar to that of the ultraviolet will be referred to hereafter as Hαcorr. light in each galaxy individually. Besides, the UV band is the We test if the foreground dust screen approximation has an one with highest spatial resolution within the bands we are using effect on edge-on galaxies. For that purpose, we plot the dif- in this work that is closely related to the SFR. Thus, for that pur- ference between A(Hα) derived from the Balmer decrement and pose we use the GALEX NUV images (deeper and available for the A(Hα) values derived from the ratio of IR/Hα as a function a few more objects than the FUV ones). We compute the differ- of galaxy axial ratio (see top panel in Figure 3). The expression ence between the NUV magnitudes obtained for the 36-arcsec- used to derive the attenuation from the ratio of IR/Hα is A(Hα) aperture and the asymptotic magnitudes. We fit the aperture cor- = 2.5 × log[1+aIR×L(IR)/L(Hαobs)] (see Kennicutt et al. 2009, rection data for the whole mother sample (those galaxies having equation 2). L(IR), in this case, corresponds to L(22µm) avail- NUV data) as a function of the galaxy size, given by the isoA able for a larger number of galaxies in our sample than L(TIR). +0.018 in the r-band from the SDSS. The correlation between aperture The coefficient aIR is equal to 0.015−0.006 (being this coefficient correction and isophotal diameter is the strongest of all those an- the average value derived from our sample in Section 4.5.2). The alyzed and it will be used for galaxies without NUV magnitudes. value given by Kennicutt et al. (2009) is aIR = 0.020 ± 0.001r ± The resulting median correction is around 1.4. The observed Hα 0.005s which is in good agreement with ours even taking into ac- luminosities already corrected for aperture effects are listed in count that they are obtained from different samples. Finally, the Table 1. difference between both A(Hα) estimations yields a mean and The spatial masking is applied over the datacubes before per- 1σ values of (0.05 ± 0.43) mag after doing a rejection of 4σ. forming the stellar continuum fitting. That means that the light This value shows that both methods produce compatible results. from spaxels contaminated by field stars and background objects As we do not see systematic residuals against the axial ratio pa- is not summed up at this stage. Then, we correct the emission- rameter, we conclude that we do not find a different behaviour in line fluxes for the flux coming from those missing spaxels. The the case of highly-inclined galaxies. mean value for the correction factor is 2.2 %. The corrections Finally, the uncertainty in the Hα flux is estimated doing are applied only over 44.7 % of the galaxies, the ones that have a random redistribution of the residuals obtained after fitting a contaminating sources. These correction factors are obtained Gaussian function to the pure emission line spectrum in the spec- comparing the aperture fluxes between two sets of synthetic tral range around Hα emission. The new residual spectrum is continuum-subtracted narrow-band images. One of them with- added to the pure emission line spectrum and a new Gaussian out the flux from the corresponding contaminated pixels and the fit is performed. This procedure is repeated 1000 times and the other where the flux from those pixels is obtained by local inter- standard deviation of the computed Hα fluxes is considered as polation. the error in the Hα flux. On the other hand, the comparison Once Hα and Hβ emission line fluxes are computed, we cor- between the measured Hβ line fluxes and those expected from rect the Hα flux for dust-attenuation assuming that the relation the Hα/Hγ Balmer decrements for the same ionized-gas physi- between Hα reddening and extinction follows the foreground cal conditions gives us an estimation for the Hβ flux uncertainty. Article number, page 5 of 36 The typical background in the GALEX UV images is very 1.5 low so the distribution of count rates in each image typically fol- lows a non-Gaussian distribution. Because of this peculiarity, we 1.0 ) [Balmer] estimate the background using the mean instead of the median or α 0.5 the mode used at high background levels, such as ground-based optical or NIR imaging (see Gil de Paz et al. 2007). Surface and ] - A(H 0.0 α aperture photometry was then carried out for each galaxy, using the IRAF task ELLIPSE as described in Gil de Paz et al. (2007), 0.5 ) [IR/H − within elliptical isophotes with fixed ellipticity and position an- α 1.0 gle (the same ones used for the extraction of the spectra from A(H − the CALIFA datacubes). In addition to the 36-arcsec aperture 2.5 0.0 0.2 0.4 0.6 0.8 1.0 mentioned above we also extracted UV photometry in other con- 2.0 centric elliptical apertures until the error in the surface photome- 1.5 ) [Balmer] 1.0 try reached 0.8 mag (including both background-subtraction and α 0.5 photon noise). From each set of concentric elliptical apertures A(H 0.0 0.0 0.2 0.4 0.6 0.8 1.0 we finally obtained asymptotic magnitudes for the whole sample Axial ratio (Gil de Paz et al. 2007). These asymptotic magnitudes are the ones applied previously in Section 3.1.4 to obtain the aperture Fig. 3: Top panel: Difference between Hα attenuations derived from corrections. the ratio of IR/Hα and those obtained using the Balmer decrement as As the UV luminosity suffers from severe attenuation by dust a function of galaxy axial ratio. Solid and dashed black lines corre- this has to be corrected in order to properly estimate the SFR. spond to the mean and 1σ values, (0.05 ± 0.43) mag, after applying The most commonly accepted method to estimate the dust at- a rejection of 4σ. Due to the good agreement between both attenua- tenuation at UV wavelengths is to use the ratio between the IR tions computed from different methods, we can safely assume that our Balmer decrement attenuations yield proper results. Besides, we do not (22-25 µm MIR, FIR or TIR) and the UV flux (also known as find any systematic residuals against the axial ratio parameter associ- infrared excess or simply IRX). This is equivalent to the use of ated with highly-inclined galaxies. Bottom panel: A(Hα) derived from hybrid SFR estimators that include information from these two the Balmer decrement as a function of the axial ratio (used as a proxy wavelengths, which is the approach used later in this work. The for inclination). Grey shadow corresponds to the 1σ intervals around IR/UV ratio is almost independent of the dust properties and the the mean value showed in black solid line. relative distribution of dust and stars (Buat et al. 2005). How- ever, it depends on the age of the dust-heating populations (see Cortese et al. 2008). In the context of this section, we analyzed A dispersion of σ = 7 % centered around unity is obtained across only the case when no IR data is available. Should that be the the whole sample. This method provides much larger uncertain- case, a relation between the FUV−NUV color and the infrared ties compared with the one using the redistribution of the residu- excess could be used instead (see Muñoz-Mateos et al. 2009, als around the Hβ emission line. The reported error includes the equation 2). For the sake of simplicity, and given the intrinsic potential uncertainties in the modeling of the stellar continuum large dispersion of the IRX-β relationship (see Section 4.5.2 for and it is taken as a conservative upper limit for the error in the Hβ the IRX-β relationship in our sample), we make use of the fol- flux. This Hβ flux uncertainty propagates to a much larger one lowing linear relation between A(FUV) and FUV−NUV and the in the corrected Hα flux. A standard error propagation method corresponding ±1σ prediction intervals: is used to compute the uncertainties in other quantities such as extinction or luminosity. A(FUV) = 0.556 + 2.292 × (FUV − NUV) (2) This linear empirical relation is based on the analysis of UV 3.2. GALEX UV imaging and infrared surface photometry of the SINGS sample (Ken- For nearly two thirds of the galaxies in the CALIFA mother nicutt et al. 2003) carried out by Muñoz-Mateos et al. (2009). sample, we could collect UV observations available from the These authors use FUV, NUV and TIR luminosity profiles with − Galaxy Evolution Explorer (GALEX2) archive (see Martin et al. the same spatial resolution to compute both, FUV NUV colors 2005). In most cases (655 out of the 663 objects with GALEX and A(FUV) attenuations via the L(TIR)/L(FUV) ratio using the UV data) this includes both far-ultraviolet (FUV, effective wave- expression given by Buat et al. (2005). This is similar to the IRX- β relationship first studied by Meurer et al. (1995) and calibrated length λeff∼1516 Å) and near-ultraviolet (NUV, λe f f ∼2267 Å) bands, 200 of them included in the sample of the 272 galaxies for starburst galaxies. However, Muñoz-Mateos et al. (2009) use analyzed in this work. star forming galaxies that have lower values of the extinction for a given FUV-NUV color. We emphasize that such relations The GALEX archive provides simultaneous co-aligned FUV (based on the UV color alone) should be used only as a rough and NUV images with a pixel scale of 1.5 arcsec per pixel and estimate of the UV light attenuation with some (limited) statisti- − a spatial resolution (FWHM) of 4 5 arcsec. We have selected cal meaning but of very little use in a case-by-case basis. More galaxies located within the central 0.5-degree radius of the 1.2- recently, Hao et al. (2011) provided a physical motivation for degrees circular GALEX FOV. We also imposed that the whole such linear relationship between UV color and attenuation and galaxy is included in the GALEX FOV. In order to calculate the yielded a y-intercept of −0.084 mag (that corresponds to a FUV- integrated mean flux of the galaxy, foreground stars and other NUV color in the absence of dust of 0.022 ± 0.024 mag) and a targets in the field were identified and removed by averaging the slope of 3.83. Taking into account that the intrinsic FUV−NUV interpolation along rows and columns from the GALEX images. color for zero attenuation is different in both cases (due to the noisy relation between A(FUV) and FUV−NUV color), we de- 2 http://galex.stsci.edu/GalexView/ cide to use Equation 2 (J. C. Muñoz-Mateos, priv. comm.) as in Article number, page 6 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample this case we have prediction intervals as a function of the UV 1045 Whole sample color. This work We apply Equation 2 to galaxies that have a FUV-NUV color less or equal to 1 mag. Galaxies with colors FUV−NUV>1 1044 ) 1 could correspond to either red galaxies with old stellar popu- − lations or galaxies with large amounts of dust reddening. In our 43 sample the mean value of the dust attenuation in the FUV is 1.73 10 (erg s magnitudes and vary from 0.81 to 2.80 magnitudes, as it is found by other authors (e.g., Buat et al. 2005; Burgarella et al. 2005). PSF 42 −1 10 The FUV−NUV colors, L(FUV) and L(NUV) in ergs s for 200 m) galaxies over the 272 galaxies analyzed in this work are listed in µ Table 1. L(22 1041 3.3. WISE MIR imaging 1040 The Wide-field Infrared Survey Explorer (WISE, Wright et al. 1040 1041 1042 1043 1044 1045 1 2010) surveyed the entire sky at MIR wavelengths 3.4, 4.6, 12 L(22µm)AP ERT URE (erg s− ) and 22 µm (W1 through W4 bands) with 5 σ point-source sensi- tivities of ∼ 0.08, 0.11, 0.8, and 4 mJy, respectively. The WISE Fig. 4: Comparison between the WISE 22 µm PSF and aperture pho- All-sky Data Release is available through the Infrared Science tometry for the entire CALIFA mother sample (white points) and for 3 the galaxies used in this work (black points). Solid line corresponds to Archive (IRSA ). It includes imaging (Image Atlas) and PSF- the 1:1 line and it is given for reference. Although the WISE Source photometry source catalogs (Source Catalog) for all four WISE Catalog is ideal for point sources and the resolution of the WISE 22 µm bands for the entire CALIFA mother sample. For the purpose of band is wide enough (FWHM ∼ 11 arcsec) it seems not to be suitable this work we make use of the WISE 22 µm data (W4-band) in or- for the CALIFA sample. der to have information on the amount of (mainly UV) photons being processed through dust absorption and re-emitted. The WISE Source Catalog is optimized for point sources and in spite The FSC is at least 98.5% reliable at 12 and 25 µm and ∼94% at of the resolution of the 22 µm data (FWHM ∼ 11 arcsec) and 60 µm . the size of the CALIFA galaxies (limited in diameter to ∼ 1 ar- We have performed a cross-match of the CALIFA mother cmin), this photometry catalog might not be appropriate for our sample with the IRAS Faint Source Catalog (closest IRAS sample (see Figure 4). Therefore, we decide to perform aperture source within 40 arcsec) finding 488 galaxies in common for the photometry using the Image Atlas in order to calculate the in- four IRAS bands. Within the IRAS FSC catalog a flux density tegrated 22 µm fluxes and magnitudes. We obtain aperture pho- measurement can be either high quality (FQUAL=3), moderate tometry in circular apertures that enclosed the entire flux from quality (FQUAL=2) or just an upper limit (FQUAL=1) (Moshir the source. A circular annulus around this aperture is used to et al. 1992). In this work we make use of only high and moderate compute the sky. We derive L(22 µm) in ergs s−1 for 265 objects quality measurements available for the CALIFA mother sample out of the 272 galaxies with detected Hβ emission included in galaxies, that yielded 12, 25, 60 and 100 µm detections for 200, this work (see Table 1). It has been pointed out by several au- 203, 486 and 443 sources, respectively. Note that poor spatial thors (Wright et al. 2010; Jarrett et al. 2013; Brown et al. 2014) resolution of IRAS in any of these bands ensures that the flux that star forming galaxies measured with the WISE 22 µm fil- measurements in the FSC are accurate for the CALIFA objects as ter are systematically brighter by ∼ 10 % than what is inferred long as the object is relatively isolated, but could have an impact from Spitzer IRS and 24 µm data. This factor has been applied on the TIR measurements of galaxies in pairs or close groups. in our 22 µm luminosities along the article. Values of L(22 µm) As noted above, the fraction of galaxies with 25 µm mea- in Table 1 should be multiplied by this correction factor. surements is significantly lower than that of galaxies with 60 µm and 60+100 µm measurements. This is due to the comparable detection limit of IRAS at 25 and 60 µm but larger flux densi- 3.4. TIR fluxes: WISE, IRAS and AKARI ties of nearby star-forming galaxies at these latter wavelengths. Although the longest WISE band already provides relevant infor- In order to recover a larger fraction of galaxies with total in- mation on the attenuation of the UV light associated to regions frared flux densities (TIR, i.e. 8-1000 µm) we decide to combine of star formation, a significant fraction of the energy re-radiated the WISE 22 µm photometry with that from IRAS to determine in the IR by dust emerges at longer wavelengths. In order to the galaxies TIR luminosity. The reliability of this procedure is account for possible differences in the dust temperature or grain- demonstrated by the tight correlation between our WISE 22 µm size distribution that could hamper the use of WISE 22 µm alone, luminosities and those detected at 25 µm with high and mod- we have also collected IRAS photometry for the entire CALIFA erate quality flux by IRAS (blue and red points in Figure 5). mother sample. Therefore, we can confidently use our WISE 22 µm photometry The Infrared Astronomical Satellite (IRAS, Neugebauer to increase the number of CALIFA galaxies with TIR measure- et al. 1984) surveyed 96% of the sky in four wavelength bands ments. In addition, we are going to use the WISE 22 µm mea- at 12, 25, 60 and 100 µm. The main data source for this paper is surements instead of the IRAS 25 µm in the corresponding IR the IRAS Faint Source Catalog v2.0 (FSC) (Moshir et al. 1990) SFR tracers as both are found compatible and there are signifi- with a detection limit about one magnitude deeper than the Point cantly more measurements from WISE 22 µm. Previous studies Source Catalog (PSC) (Beichman et al. 1988), reaching a depth (Kennicutt et al. 2009; Calzetti et al. 2010) show that the average of ∼0.2 Jy at 12, 25 and 60 µm and greater than 1.0 Jy at 100 µm. ratio between 24 µm and 25 µm luminosities is 0.98 ± 0.06. In our case, we find that the average ratio between 25 µm and 22 3 http://irsa.ipac.caltech.edu µm luminosities is 1.05 ± 0.22 when using high quality 25 µm Article number, page 7 of 36 1045 luminosities obtained for each galaxy is listed in Table 1. We High quality (IRAS) adopt these values as the best measure of the TIR luminosity. Moderate quality (IRAS) ) 1 44 Upper limit (IRAS) − 10 4. Results 1043 Our aim is to determine the different SFR estimators (single- m) (erg s µ band and hybrids) for the CALIFA sample and, in particular, en- 42 10 sure that Hαcorr can be safely used for future statistical spatially- L(25 resolved studies, at least in the Local Universe. 1041 There are three different ways to carry out this analysis: by 1041 1042 1043 1044 1045 comparing fluxes, luminosities (or SFRs) and surface bright- ) ) 3 Article number, page 8 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample 4.1. Recipes for determining the SFR in galaxies −1 −44 The SFR indicators considered are of two types: single-band and SFR (M yr ) = 4.6 × 10 [L(FUVobs) + 0.46 × L(TIR)] (8) hybrid. In the case of the recipes based on a single photometric band we use the extinction-corrected UV (from the UV slope), −1 −42 extinction-corrected Hα (from the Balmer decrement) and the SFR (M yr ) = 5.5 × 10 [L(Hαobs) + 0.021 × L(22µm)] (9) observed MIR or TIR luminosities. The hybrid tracers combine luminosities measured directly (observed UV or Hα) with that of the light emitted by dust after being heated by young massive −1 −42 SFR (M yr ) = 5.5×10 [L(Hαobs)+0.0024×L(TIR)] (10) stars (see Gordon et al. 2000; Inoue et al. 2001; Hirashita et al. 2003; Iglesias-Páramo et al. 2006; Calzetti et al. 2007; Kennicutt The original recipes for the hybrid tracers make use of the et al. 2007, 2009; Hao et al. 2011; Kennicutt & Evans 2012; 25 µm luminosity but we are interested in using our 22 µm lu- Calzetti 2013; Domínguez Sánchez et al. 2014, for more details). minosities instead as we have a large number of these measure- The most widely-used recipes for SFR tracers are included ments. The reader is referred to Section 3.4 where we justify in Calzetti (2013) and are listed here for convenience. These the use of L(22) instead of L(25) after a 1.05 ± 0.22 conversion expressions are used to compute the SFR from different data, factor is applied. This factor is computed as the average ratio both for single-band and hybrid recipes, scaled to the same IMF between 25 µm and 22 µm luminosities when using high quality (Kroupa 2001). The mass range varies from 0.1 to 100 M . The 25 µm IRAS measurements. value of the timescale over which the star formation must re- main constant depends on each tracer, being up to 100 Myr for 4.2. SFR surface density the UV, MIR or TIR and having a lower value for the Hα tracer, equal or larger than 6 Myr. The expressions listed below are As mentioned before, the only tracers with enough spatial res- for global scales as we are using integrated fluxes for the whole olution to compute SFR surface densities across the CALIFA galaxy in each case. Also, recipes for determining the SFR at sample are the UV (FWHM ∼ 4.5 arcsec) and Hα (FWHM ∼ local scales could be found in the review of Calzetti (2013). It is 2.5 arcsec) measurements. We calculate the SFR surface density worth noting that for the case of Equations 6, 7 and 9, we have in both as the SFR per unit area measured in the largest ellip- rescaled the coefficients that multiply L(22 µm) taking into ac- tical apertures (semi-major axis = 36 arcsec) fitting the PPaK count the L(25 µm)/L(22 µm) ratio obtained for our sample and FOV with the ellipticity and PA of the corresponding galaxy. the average ratio between L(24 µm) and L(25 µm) derived in pre- The Hα data are corrected for extinction using the Balmer decre- vious studies (Kennicutt et al. 2009; Calzetti et al. 2010) as ex- ment measured within these apertures. In the case of the UV, we plained in Section 3.4. The non-linear behavior for galaxies with use the hybrid tracer (Equation 7). This tracer combines UV- L(22 µm) > 5×1043 erg s−1 present in the original recipe (see observed luminosities with 22 µm luminosities (FWHM ∼ 11 Calzetti 2013) is not included here as we only find four galaxies arcsec). Because the negative dust extinction gradients found in in that range. L(TIR) is the total infrared emission in the range star-forming galaxies virtually all the flux at 22 µm was found to 8−1000 µm. come from inside these elliptical apertures (see Muñoz-Mateos First we list those based on single-band, where all the lumi- et al. 2009). Nevertheless, in order to avoid systematic offsets we nosities are in units of ergs s−1 : de-corrected for aperture the total 22 µm fluxes using the same aperture correction as described in Section 3.1.4. This means −1 −44 that now all the fluxes, 22 µm, Hα and FUV, are calculated for SFR (M yr ) = 4.6 × 10 × L(FUVcorr) (3) the same area. Thus, the SFR surface density is computed using the follow- ing expression: −1 −42 SFR (M yr ) = 5.5 × 10 × L(Hαcorr) (4) SFR Σ = (11) SFR 2 πa2 d −1 −44 206265 SFR (M yr ) = 2.8 × 10 × L(TIR) (5) Where the expressions used for estimating the SFR values are Equations 4 and 7 given in Section 4.1. The parameter a corresponds to the semi-major axis set to 36 arcsec in all cases SFR (M yr−1) = 2.10 × 10−43 × L(22µm) (6) as described in Section 3.1.3 and d is the distance in Mpc to the The hybrid tracers are obtained assuming an approximate galaxy calculated from its redshift (listed in Table 1). energy-balance approach. The expressions for the hybrid trac- Figure 6 compares the hybrid star formation surface density ers are shown below where the luminosities are observed and using the observed FUV and 22 µm fluxes with their correspond- are in units of ergs s−1. The global coefficients, 4.6 × 10−44 and ing Hα attenuation-corrected star formation surface density. We −42 −1 −1 5.5 × 10 [M yr /ergs ], correspond to the calibration of the have excluded elliptical and lenticular galaxies in this plot where single-band or monochromatic indicators shown before, UV and part of the UV emission could come from HB stars responsible Hα respectively. On the other hand, the coefficients that multiply for the UV-upturn (Brown et al. 1997; Yi et al. 1997). the IR luminosity, either L(22 µm) or L(TIR), are dependent on We found a good linear correlation between both mea- this tracer and on the one used for the direct stellar light emis- surements in a wide range of values of ∼2 dex, specially, −9 −1 −2 sion. We are going to calibrate empirically these coefficients in at ΣSFR[Hαcorr(36)]>10 M yr pc . The mean value of Section 4.5.2 to create dust-corrected SFRs: Article number, page 9 of 36 SF (N=104) ) SF (N=164) 1 ] Type-2 AGN (N=70) − 101 Type-2 AGN (N=99) µm yr ) 8 2 10− − pc (36)+22 1 100 − obs m] (M yr 9 µ FUV 10 [ − M 1 ( 10− SFR Σ ] SFR[22 µm 10 2 10− 9 8 10− ] 1 10 10 ] 1 (36)] − −1σ = 0.24dex 2 1 0 1 10 ± 10 10− 10− 10 1σ =10 0.29dex corr (36)+22 ± corr µm 0 obs 0 Hα 10 Hα [22 [ 10 [ FUV [ 1 SFR SFR 9 8 Σ 10− 10− 2 1 0 1 1 2 10− 10− 10 10 ΣSFR[Hαcorr(36)] (M yr− pc− ) 1 SFR[Hαcorr] (M yr− ) Fig. 6: Star formation rate surface density derived using a combi- Fig. 7: Comparison between the MIR (22 µm) and the Balmer-corrected nation of observed-FUV and 22 µm luminosities compared to Balmer Hα SFR tracers (recipes from Calzetti (2013); see Equations 6 & 4). attenuation-corrected Hα star formation rate surface density. The val- Solid line corresponds to the 1:1 line. Orange points show type-2 AGN ues are obtained within an area of semi-major axis equal to 36 arcsec. galaxies and blue points represent star-forming galaxies. The legend The solid line corresponds to equal Σ on both axis. Bottom part of SFR gives the number of objects available in both cases. Residuals appear this plot represents the residuals as a function of the Balmer-corrected in the bottom part of this figure as a function of the Balmer-corrected Hα star formation rate surface density. The mean value is shown with Hα SFR tracer. Solid line shows the mean value of -0.08 when all the the solid line while dashed lines correspond to the 1σ dispersion in dex galaxies are included and the dashed lines are referred to the 1σ disper- around the mean value. Black, blue and orange numbers correspond to sion (±0.29 dex) around it. For SF galaxies the mean values is -0.15 the mean values for the whole sample, SF and type-2 AGN host galax- while for type-2 AGN host galaxies is 0.03. ies, respectively. The tight relation found for these two SFR density tracers shows that there are not systematic differences between them. Then, we can safely use the SFR instead of these measurements along this work. 4.3.1. Single-band SFR tracers Comparison between Mid-IR and extinction-corrected Hα The consistency between the two star formation surface den- In this section we compare the SFRs using warm-dust sity values and the large range involve shows that there are no sensitive 22 µm WISE luminosities with Balmer attenuation- systematic differences between the two tracers when SFR sur- corrected Hα SFRs (Equations 6 and 4, respectively) in Fig- face densities are used or, at least, these are of the order of the ure 7. We find that at high luminosities 22 µm reproduces object-to-object variation. Thus, we can use the SFR instead of the SFR measured with extinction-corrected Hα. Neverthe- SFR density surface from now on in this work, which will allow less, L(22 µm) underestimates the SFR expected at low lumi- us to use all TIR measurements confidently. nosities, where these galaxies are supposed to have very little dust and consequently weak L(22 µm) emission. This could be the reason why the mean value of the residuals, expressed as 4.3. Comparison of the different SFR tracers Article number, page 10 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample tion given by Equation 2 mentioned in Section 3.2 to correct the from higher attenuation. Dark-green and black dashed lines in FUV luminosity. the same figure correspond to values of Hα attenuation of 1 and 2 We find a rather noisy relation of ±0.36 dex around the mean magnitudes, respectively. It seems like values of A(Hα) around 1 value, Article number, page 11 of 36 SF (N=106) ) SF (N=117) 1 Type-2 AGN (N=71) − 1 Type-2 AGN (N=63) A(Hα) = 0mag, E(B V ) = 0.44E(B V ) ) s g 10 1 − − yr A(Hα) = 1mag, E(B V )s = 0.44E(B V )g − 100 − − A(Hα) = 2mag, E(B V )s = 0.44E(B V )g yr − − ] (M 100 corr 1 10− 1 (36)] (M 10− obs SFR[FUV ] ] 1 0 1 1 10− 10 10 corr 10 1 σ = 0.36dex corr 2 ± 10− 0 Hα [ 10 FUV SFR [FUV [ 1 10− 3.0 E(B-V)s = 0.44 E(B-V)g Hao+11 E(B-V)s = E(B-V)g 5 Buat+05 9 1 2 ΣSFR[Hαcorr] > 3 10− (M yr− pc− ) 2.5 × 4 2.0 3 1.5 ) [Balmer]1 (mag) .0 α 2 A(H 0.5 A(FUV) [FUV-NUV] (mag) 1 0.0 0 0 1 2 3 4 5 0 1 2 3 4 5 A(FUV) [IRX] (mag) A(FUV) [IRX] (mag) Fig. 8: Top left panel: FUV-corrected SFR as a function of Balmer-corrected Hα SFR, using Equations 3 & 4, respectively. Color-coding, solid and dashed lines have the same meaning as in Figure 7. At high L(Hα)corr the SFR derived from the UV alone is underestimated. It might be that the extinction correction using the FUV−NUV color traces only the most superficial and less extinct part of the SFR. Consequently, the higher SFRs associated with higher values of the extinction are being underestimated. Top right panel: Relation between observed-FUV SFR and observed-Hα SFR. Light-green, dark-green and black dashed lines corresponds to values of A(Hα) equal to 0, 1 and 2 magnitudes, respectively. All cases are based on the assumption that the relation between the color excess of the stars and the gas is E(B-V)s=0.44E(B-V)g. Bottom left panel: Comparison between A(Hα) from the Balmer decrement (Equation 1) and A(FUV) from IR/FUV flux ratio using the expression by Hao et al. (2011). Blue dashed line represents a relation between color excess of E(B-V)s=0.44E(B-V)g while red dashed line assumes that the relation is E(B-V)s=E(B-V)g. Grey points show cases of A(Hα) equal to zero magnitudes while light-green points show galaxies with the highest values of the SFR surface density. This plot suggest that we might apply a higher value for the relation between the color excess of the gas and the stellar continuum that the one found by Calzetti et al. (2000) for our sample, although galaxies with higher values of the SFR surface density are more similar to this previous relation. Bottom right panel: Comparison between A(FUV) derived using the FUV−NUV color (Equation 2) and A(FUV) from IR/FUV flux ratio. A discrepancy between these two expressions is found for the lowest and highest values of the attenuation. Star-like symbols show the values when the expression used to compute the A(FUV) is the one from Buat et al. (2005) while open circles show the values from Hao et al. (2011). The solid line shows the 1:1 line for reference. to the current SF we want to trace. Nevertheless, optical photons thors derive that if τ< 5 Gyr the UV light contributes less than from old stars contribute to the heating of the dust (see Johnson 50% to the TIR emission. et al. 2007) and thus, to the TIR luminosity. Indeed, based on constant star formation (CSF) models, Calzetti (2013) found a The comparison between the SFR[TIR] and SFR[Hαcorr] reduction in this constant of almost a factor of 2 from models (Figure 9) shows that at low TIR luminosities the SFR[TIR] is with a CSF lasting for 100 Myr compared to those CSF mod- being underestimated (in the line of the results of Rieke et al. els lasting for over 10 Gyr. Besides, according to Cortese et al. 2009). We find that for values of the SFR[TIR] below 0.3 M −1 (2008), for star formation timescales (equivalently the age of the yr the average value of AHα is 0.28 ± 0.04 mag. On the other −1 Universe at which the SFR peaks in their ’a la Sandage’ SFH) hand, at high luminosities (SFR[TIR] > 1 M yr ) the TIR larger than ∼6-7 Gyr the UV radiation dominates the dust heat- seems to provide SFRs somewhat higher than those obtained ing with a contribution of >75% to the total energy absorbed and from Hα. In fact, a large number of galaxies appear in this then re-emitted in the infrared. On the other hand, the same au- regime making the mean value of the ratio between these tracers larger than zero in the residuals, Article number, page 12 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample Once we have explained the assumptions imposed to the use SF (N=135) ) 1 of the hybrid tracers, we compare them with our Hα extinction- 101 Type-2 AGN (N=83) − corrected SFR tracer. In the first place, we examine the be- yr haviour using Hα observed luminosity combined with 22 µm and 0 TIR luminosity. Figure 10 shows that applying the method ex- 10 plained in Kennicutt et al. (2009) we now obtain very similar results to theirs but using a larger sample and for the first time with IFS data. Secondly, we replace Hα observed luminosity 1 10− with FUV observed luminosity combined with the IR luminosi- SFR[TIR] (M ties (Figure 11). In both cases we find a very good correlation ] 1 0 1 across 2.5 dex in SFR but with an offset in the mean ratio of ] 1 10 10− 10 1σ10= 0.27dex corr ± SFRs of 25 per cent. This offset goes in the sense that SFR de- TIR [ Hα 0 [ 10 rived from the hybrid Hα+IR and FUV+IR SFR tracers is larger than for the extinction-corrected Hα one. As explained before, SFR SFR 1 10− 1 0 1 one possibility could be the presence of optical photons from old 10− 10 10 1 stellar populations heating the dust specially at long IR wave- SFR[Hαcorr] (M yr− ) lengths or the effects of AGN. In particular, if we discriminate between star-forming and type-2 AGN galaxies when computing Fig. 9: Comparison between the SFR tracer using the TIR luminosity and the extinction-corrected Hα SFR tracer, after applying Equations 5 these ratios, type-2 AGN host galaxies yield larger offsets than and 4, respectively. Color-coding, solid and dashed lines have the same the ones reported for SF galaxies. We refer the reader to Section meaning as in Figure 7. 4.5.2 for an extensive discussion on this issue. We conclude that when comparing the hybrid calibrators with Balmer decrement attenuation-corrected Hα SFR tracer, we find tighter correlations than those obtained with single-band = 0.11. This is either because the contribution of heating due tracers (see the 1σ dispersions around the mean values in Fig- to optical photons or nuclear activity becomes relevant at those ures 7, 8, 9, 10 and 11.) luminosities and/or because a fraction of the Hα recombination line luminosities are not recovered when correcting for dust at- tenuation using the Balmer decrement. The analysis of the hy- brid calibrations (see section below) favors the former scenario. 4.4. Origin of the discrepancies among SFR tracers 4.3.2. Hybrid SFR tracers As we have seen in the previous section, there is a general good agreement between the SFR tracers considered, single-band and A powerful way of determining the SFR is to combine a photo- especially hybrids, compared to the attenuation-corrected Hα metric band that is sensitive to the light directly emitted by young SFR tracer. Nevertheless, we can appreciate some differences if stars (i.e. observed UV or Hα fluxes) with that reprocessed by we take a closer look at these relations. In the case of the single- dust, usually in the mid- or far-infrared (or, alternatively, the TIR band tracers the main problems appear when using FUVcorr lu- emission). This is an alternative approach to correct the UV or minosities as the extinction correction is a big problem to deal Hα fluxes for attenuation. with, specially for high SFR values. However, we can mitigate In both cases, the validity of these hybrid tracers is that the this effect using hybrid tracers combining the FUVobs luminos- observed IR emission comes from light whose optical depth (or ity with the IR luminosity, both 22 µm and TIR. In the latter attenuation) is of the order of that in the UV or Hα, otherwise it case we are assuming that we can recover all the light that it would not be possible to write the total SFR as a sum of the two has been re-emitted by the dust. Similar cases appear when us- luminosities, observed and dust-processed (see Kennicutt et al. ing the single-band tracers for 22 µm or TIR luminosities, where 2009). This assumption could not be valid if (1) the heating of we apparently lose some SF in galaxies with low values of the the dust is dominated by optical photons particularly important at SFR. Again, when using hybrid tracers the agreement between long IR wavelengths where the contribution of low-temperature calibrators improves. dust emission is most relevant or (2) by UV photons more en- One of the main reasons behind these discrepancies is the ergetic than those observed directly (e.g. if the bluest observed different selection criteria used in the process of determining the band is in the NUV) or (3) in the case of a significant AGN con- SFR calibrators in the literature. Now, we have the opportu- tamination, where any of these bands could be actually tracing a nity to re-calibrate these tracers for a diameter-limited sample of UV radiation field that is not merely due to recently-formed mas- 380 galaxies. Moreover, we are able to use integral field spec- sive stars. If there is a significant difference between the τFUV troscopy data to assure a proper determination of the attenua- (or τHα) and the opacity of the photons that lead to the IR emis- tion using the Balmer decrement avoiding the problems associ- sion used in the corresponding tracer, a linear relation between ated with narrow-band imaging. Thus, we are going to provide the SFR and the two (emitted and dust-absorbed) luminosities updated SFR tracers based on our state-of-the-art attenuation- should not be present. In the particular case of UV and Hα the corrected Hα luminosities. τ are similar to the one that comes from the dust component so the approach of using a linear relation between L(FUVobs) 4.5. Updated SFR tracers for the diameter-limited CALIFA or L(Hαobs) combined with the L(IR) luminosity can be safely sample done. That implies that for the IR tracer both bands should suf- fer the same attenuation (see a detail analysis in Kennicutt et al. We now provide updated calibrations for the global current SFR 2009; Hao et al. 2011) in external galaxies by means of anchoring the different trac- Article number, page 13 of 36 SF (N=164) SF (N=113) Type-2 AGN (N=99) 1 Type-2 AGN (N=74) 1 m] 10 m] 10 µ µ ) ) 1 1 + 22 − − + 22 0 yr yr 0 10 obs 10 obs α (M (M 1 1 10− 10− SFR[H SFR[FUV ] ] µm ] ] µm 1 1 0 1 1 0 1 10 10− 10 1σ10= 0.18dex 101 10− 10 1σ10= 0.23dex corr corr ± +22 ± +22 0 obs 0 obs Hα Hα 10 10 [ [ Hα 1 1 [ FUV 10− 10− 10 10 SFR 10− 10 10 1 1 SFR[Hαcorr] (M yr− ) SFR[Hαcorr] (M yr− ) SF (N=135) SF (N=94) Type-2 AGN (N=83) Type-2 AGN (N=62) 101 101 ) ) + TIR] 1 1 + TIR] − − obs yr yr obs 0 0 α 10 10 (M (M SFR[H SFR[FUV 10 1 10 1 − ] − ] 1 1 ] ] 10 1 0 1 1 0 1 10− 10 10 TIR 10 10− 10 10 TIR 1σ = 0.19dex 1σ = 0.22dex + corr corr ± ± + 0 0 obs obs 10 10 Hα Hα [ [ Hα 1 [ 1 FUV [ 10− 10− SFR 1 0 1 1 0 1 10− 10 10 SFR 10− 10 10 1 1 SFR[Hαcorr] (M yr− ) SFR[Hαcorr] (M yr− ) Fig. 10: Top panel: Hαobs + 22 µm hybrid tracer as a function of Fig. 11: Top panel: Comparison between FUVobs + 22 µm hybrid tracer Balmer-corrected Hα SFR, using Equations 9 and 4, respectively. and the Balmer-corrected Hα SFR, using Equations 7 and 4, respec- Color-coding and lines have the same meaning as in Figure 7. Bot- tively. Color-coding, lines and residuals have the same meaning as in tom part shows the residuals as a function of Hα-corrected SFR being Figure 10. Bottom panel: Same as the top panel this time showing the mean value 0.06, 0.03 and 0.13 for all the galaxies, SF galaxies and FUVobs + TIR hybrid tracer as a function of Balmer-corrected Hα SFR type-2 AGN host galaxies, respectively. Dashed lines represent the 1σ (Equations 8 and 4). The hybrid tracers shown here and in Figure 10 dispersion in dex around the mean value. Bottom panel: Same as the reduce the dispersion when compared with single-band tracers. top panel but showing the Hαobs + TIR hybrid tracer as a function of Balmer-corrected Hα SFR instead (Equations 10 and 4). clear region that could expand until 1kpc from the central region. We estimate the level of contamination of the emission from the ers (single-band and hybrid ones) to the SFR derived from the AGN host galaxies to the total SFR. We find that the contribution extinction-corrected Hα luminosity measured in our sample of of the attenuation-corrected Hα luminosity in the nucleus (mea- CALIFA galaxies. Seminal works in this context include Kenni- sured in a 3"-diameter aperture) over the total one for galaxies cutt (1998a), Kennicutt et al. (2009) and Hao et al. (2011). classified as type-2 AGN is 8.3% while for the purely SF galax- As we are interested in calibrating the SFR tracers we need ies this contribution is 5.1%. Galaxies classified as type-2 AGN to exclude galaxies that have type-1 AGN signatures to avoid are shown in our plots as orange points. contamination of sources that are not star-forming (only galax- We first provide updated calibrations in the case of the ies UGC 00987 and UGC 03973 are classified as type-1 AGN single-band tracers. We do not perform this analysis in the case within our sample). As explained in Section 4.3, the information of the SFR[FUVcorr] because, as we have explained before, the regarding the optical AGN classification can be found in Walcher attenuation correction is highly uncertain and the SFR tracer pro- et al. (2014). We provide separate calibrations for the sample posed would not be reliable. On the other hand, the estimation when type-2 AGN galaxies are included and when they are not. of the hybrid tracers using FUVobs luminosity will be pursued. The reason for this is that, despite numerous efforts (Alonso- Herrero et al. 2006a; Díaz-Santos et al. 2008, 2010; Castro et al. 4.5.1. Single-band tracers 2014), the fraction of UV or line emission arising from circum- nuclear star formation in type-2 AGN is still highly uncertain. In this section we provide calibrations for the observed 22 µm We remind the reader that the nuclear emission in type-2 AGN and TIR luminosities as tracers of the SFR anchoring them to the includes the contribution of both a dusty torus (external radius of SFR given by the extinction-corrected Hα luminosity according a few parsecs, see Ramos Almeida et al. 2009) and a circumnu- to Equation 4. Figure 12 shows the relation between L(Hαcorr) Article number, page 14 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample ) 43.0 SF (N=164) ) 43.0 SF (N=164) 1 1 − 42.5 Type-2 AGN (N=99) − 42.5 Non-linear fit (log scale) Non-linear fit (log scale) Non-linear fit (curve fit) 42.0 Non-linear fit (curve fit) 42.0 Linear fit Linear fit 41.5 41.5 )] (ergs s )] (ergs s 41.0 41.0 cor cor α 40.5 α 40.5 40.0 40.0 39.5 39.5 log [L(H log [L(H 41 42 43 44 40.5 41.0 41.5 42.0 42.5 43.0 43.5 44.0 1 1σ = 0.30dex 1 1σ = 0.27dex ± ± 0 0 1 1 Residuals − ) SF (N=135) ) SF (N=135) 1 43.0 1 43.0 − Type-2 AGN (N=83) − Non-linear fit (log scale) 42.5 Non-linear fit (log scale) 42.5 Non-linear fit (curve fit) Non-linear fit (curve fit) Linear fit 42.0 Linear fit 42.0 )] (ergs41 s .5 )] (ergs41 s .5 cor cor α 41.0 α 41.0 40.5 40.5 40.0 40.0 log [L(H log [L(H 42.0 42.5 43.0 43.5 44.0 44.5 45.0 45.5 42.0 42.5 43.0 43.5 44.0 44.5 1 1σ = 0.29dex 1 1σ = 0.25dex ± ± 0 0 1 1 Residuals − Fig. 12: Updated calibrations for the 22 µm and TIR single-band SFR tracers anchoring them to extinction-corrected Hα luminosity (Equation 4). Orange points in left panels correspond to type-2 AGN while blue points refer to star-forming galaxies. Linear fits are shown in dashed lines while non-linear fits are shown in dotted lines. The results for these calibrations appear in table 2. The residuals are computed as the average value of the −42 log[5.5×10 ×L(Hαcorr)/a×L(IR)] where L(IR) could be 22 µm or TIR for the case of the linear fits, after applying a 4 σ rejection. These values are computed for all galaxies (black), star forming (blue) and type-2 AGN host galaxies (orange). and the observed infrared luminosities. We include in these plots values for these fitting parameters with 22 µm and TIR luminosi- h i h 0 i 4 L(IR) a ties. both non-linear , log[L(Hαcorr)]=b×log 43 +log × −42 10 h 5.i5 10 a Figure 12 also shows the results of this analysis after in- and linear fits, log[L(Hαcorr)]=log[L(IR)]+log 5.5×10−42 ; (see Calzetti 2013). We warn the reader that the use of non-linear cluding (left panels) or excluding (right panels) type-2 AGN calibrations should be restricted to studies using similar (1) se- from the sample. In all cases, with and without type-2 lection criteria, (2) apertures and (3) corrections or the resulting AGN and using either the 22 µm or TIR luminosity, a non- SFRs could be affected by severe systematics. Linear fits are linear behaviour is clearly present, especially at low lumi- shown in dashed line while non-linear fits are shown in dotted nosities (log[L(22µm)]<41.8 or log[L(TIR)]<43.3), where most lines. The coefficients for these fits are given in Table 2. Note galaxies are located above the best linear fit (see Figure 12). that in the case of the linear fit we name the constant a which is On the other hand, galaxies with high 22 µm luminosities −1 −1 (log[L(22µm)]>43.4) are all found below the linear fit, as also expressed in units of M yr /ergs s and for the non-linear fit we use a0 because it lacks the physical meaning of a. We are found by Rieke et al. (2009) for 24µm luminosities above 5×1043 ergs s−1. The best-fitting global slope for our non-linear going to use aIR for the hybrid tracers as in this case it is dimen- sionless and has a different physical meaning than the previous SFR calibrations based on 22 µm luminosity, 0.733 (0.702) when constants (see Section 4.5.2 for more details). The values for type-2 galaxies are (not) included in the sample, is somewhat a0 and the exponent b are obtained by two different methods, a smaller (less linear) than the local value (500 pc scale) of 0.885 least-squares linear fit in log scale and a non-linear least squares obtained by Calzetti et al. (2007) and than the value of 0.82 given fit using the Python task curve_fit. Both methods yield similar by Cluver et al. (2014) from the analysis of the GAMA survey. Regarding a0, the values can be very different from those in the 4 We normalize the luminosities to 1043 ergs s−1 to ensure that the y- literature (even their units are different, obviously) but the un- intercept for these non-linear fits is located near the values covered by certainties are of the order of 5-18%, similar to the value quoted our sample by Calzetti (2013), when the log is computed in normalized lu- Article number, page 15 of 36 minosities. It is worth noting that the 22 µm or TIR luminosities 4 explored by our sample are significantly lower than those of the SF Type-2 AGN sample studied by Rieke et al. (2009), which explains why these Hao+11 authors only needed to make use of a non-linear fit at their high- 3 This work luminosity end. For the case of the linear fit, the difference between the a coefficients with and without type-2 AGN is very small, lead- 2 ing to smaller a coefficient when these objects are included by roughly 7 per cent in the case of the 22 µm and 18 per cent in the IRX TIR. This is likely due to the enhanced contribution of an AGN 1 or, alternatively, obscured circumnuclear star formation to the infrared emission compared to Hα. In the former case, the use of this calibration would remove, statistically speaking at least, 0 part of the AGN contamination, although some fraction of the Hα could still arise from the AGN. Regarding the latter possibil- ity, the use of a calibration anchored to the extinction-corrected 1 Hα luminosity would slightly underestimate the total SFR, as − 0.0 0.5 1.0 1.5 2.0 2.5 the star formation due to highly-obscured circumnuclear regions FUV - NUV (mag) in type-2 AGN could be missed. In general, independently of Fig. 13: IRX-β relation for the galaxies that have FUV and TIR mea- its origin (star formation or not), the total UV light emitted in surements in the CALIFA mother sample. Blue points represent star- these regions is hardly recovered using the UV slope or even forming galaxies while orange points correspond to type-2 AGN host the Balmer decrement as a measure of its dust attenuation, espe- galaxies. Black line shows the fit from Hao et al. (2011) while green cially in type-2 AGN where the BLR is completely hidden from line shows our fit. We obtain a value of aIR = 0.33 ± 0.08 for our fit. us. The difference here is that if the emission is due to either a BLR or NLR is something that should not be accounted for in terms of the SFR anyway. Finally, we can not rule out at this stage that, since we are CHα dealing with single-band tracers, this difference arises from a L(Hαcorr) − L(FUVobs) a = CFUV (17) dependence of the attenuation with the level of nuclear activity IR L(IR) at a given SFR. The analysis of the hybrid tracers and the dependence of the Where CHα and CFUV are the constants that multiply the −42 aIR coefficient with the attenuation presented later in this work L(Hα) and L(FUV) in Equations 4 and 3 (5.5×10 and −44 −1 −1 favors the scenario where it is the contribution of the AGN itself 4.6×10 [M yr /ergs ]), respectively. what leads to these small changes in the single-band SFR cali- Histograms in Figure 16 show the distribution of the aIR co- brations. This will be studied in more detail in a future (spatially- efficient for different hybrid SFR tracers. We refer the reader to resolved) analysis. Although, as mentioned before, we will still the end of this section and Sections 4.5.3 through 4.5.7 for an be unable to disentangle the relative contribution of AGN or extensive analysis on the nature of the variation of aIR. circumnuclear star formation to the nuclear emission of type-2 In the case of the combined UV + IR SFR tracers, there are AGN host galaxies using these data (see Alonso-Herrero et al. several ways of estimating the calibration. The most common 2006a; Díaz-Santos et al. 2010; Castro et al. 2014, for alterna- methods are (1) using an energetic balance approach once we tive approaches). have corrected for attenuation in the UV or (2) anchoring our data to other SFRs measurements. With regard to (1), we must keep in mind that the estima- 4.5.2. Hybrid tracers tion of the FUV attenuation is challenging so several methods In the case of the hybrid indicators we assume a simple energy have been put forward for that purpose. The most common ap- balance (see Kennicutt et al. 2009, for more details). proaches are the use of the β slope of the UV continuum (similar but not identical to the FUV−NUV color) or the ratio of IR to UV luminosity. We obtained the FUV−NUV-corrected FUV lu- −1 −42 SFR (M yr ) = 5.5 × 10 [L(Hαobs) + aIR × L(IR)] (14) minosity (using the β slope) in Section 4.3.1, but the results show that with this method recovering the SFR is complicated, espe- cially for the highest values of the SFR where the attenuation appears to be underestimated. Other expressions for the FUV −1 −44 SFR (M yr ) = 4.6 × 10 [L(FUVobs) + aIR × L(IR)] (15) attenuation using the FUV−NUV color can be found in the lit- erature (e.g., Kong et al. 2004; Seibert et al. 2005; Salim et al. where L(FUVobs) and L(Hαobs) are the observed luminosi- 2007; Hao et al. 2011). As an example, Hao et al. (2011) use ties in ergs s−1 and L(IR) could be either L(22 µm) or L(TIR), their own expression for the A(FUV) (equation 16 in their pa- also in ergs s−1 . per) to calibrate the TIR + FUV hybrid tracer obtaining that the − We calculate the value of the dimensionless aIR coefficient in FUV NUV-corrected FUV luminosity also underestimates the the previous hybrid relations as the median of the following ratio highest SFRs. As explained before, the other way of deriving the for the L(Hα) case (Equation 9 & 10) and the same with L(FUV) attenuation is using the ratio of IR to UV luminosity. We can use (Equation 7 & 8): the IRX-corrected FUV luminosity to calibrate the TIR + FUV and the 22 µm + FUV hybrid tracers as done in sections 4.1 and L(Hα ) − L(Hα ) 4.3 in Hao et al. (2011). The authors obtain a relation (equation a = corr obs (16) IR L(IR) 13 in their paper) between the IRX and the FUV−NUV observed Article number, page 16 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample ) ) 1 1 ) ) 1 1 SF (N=113) − SF (N=113) − − 42.5 1.2 − 42.5 1.2 Type-2 AGN (N=74) yr yr 42.0 0.7 42.0 0.7 41.5 0.2 )] (M 41.5 0.2 )] (M )] (ergs s )] (ergs s corr corr α α corr 41.0 -0.3 corr 41.0 -0.3 α α 40.5 -0.8 40.5 -0.8 40.0 aIR = 3.55 -1.3 40.0 aIR = 4.52 -1.3 log [L(H log [L(H log [SFR(H log [SFR(H 42.0 42.5 43.0 43.5 44.0 44.5 45.0 42.0 42.5 43.0 43.5 44.0 44.5 45.0 1 1σ = 0.23dex 1 1σ = 0.19dex ± ± 0 0 1 Residuals − Residuals − 42.0 42.5 43.0 43.5 44.0 44.5 45.0 42.0 42.5 43.0 43.5 44.0 44.5 45.0 1 1 log [ L(FUV ) + a L(22µm) ] (ergs s− ) log [ L(FUV ) + a L(22µm) ] (ergs s− ) obs IR · obs IR · ) ) 1 1 ) SF (N=94) ) SF (N=94) 1 1 − − − 42.5 1.2 − 42.5 1.2 Type-2 AGN (N=62) yr yr 42.0 0.7 42.0 0.7 41.5 0.2 )] (M 41.5 0.2 )] (M )] (ergs s )] (ergs s cor cor α α cor 41.0 -0.3 cor 41.0 -0.3 α α 40.5 -0.8 40.5 -0.8 40.0 aIR = 0.33 -1.3 40.0 aIR = 0.40 -1.3 log [L(H log [L(H log [SFR(H log [SFR(H 1 42.0 42.5 43.0 43.5 44.0 44.5 45.0 1 42.0 42.5 43.0 43.5 44.0 44.5 45.0 1σ = 0.22dex1 1σ = 0.19dex1 log [ L(FUVobs) + aIR L(TIR) ] (ergs s±− ) log [ L(FUVobs) + aIR L(TIR) ] (ergs s±− ) 0 · 0 · 1 Residuals − Residuals − 42.0 42.5 43.0 43.5 44.0 44.5 45.0 42.0 42.5 43.0 43.5 44.0 44.5 45.0 1 1 log [ L(FUV ) + a L(TIR) ] (ergs s− ) log [ L(FUV ) + a L(TIR) ] (ergs s− ) obs IR · obs IR · Fig. 14: Updated calibrations for the hybrids tracers that combine FUV observed luminosity and IR luminosity. Top panels show the FUVobs + 22 µm hybrid tracer while the FUVobs +TIR hybrid tracers appear on the bottom panels. Galaxies hosting type-2 AGN (orange points) are included on the left panels. Blue points refer to star-forming galaxies. Dashed lines correspond to the 1:1 line taking into account the CHα and −42 −44 CFUV constants (5.5×10 and 4.6×10 ) given in Equations 3 and 4, respectively. The best fitting aIR coefficients calculated as the median value of the expression 17 are shown for clarity. These aIR values and their corresponding errors appear in table 3. The residuals are computed as the average value of the log[CHα×L[Hαcorr]/(CFUV ×(L[FUVobs] + aIR×L[IR]))], where L(IR) could be 22 µm or TIR, after applying a 4 σ rejection. These hybrid tracers show a trend with the aIR coefficient, so when type-2 AGN host galaxies are included the value of aIR decreases. color. An important parameter that appears in the former equa- Finally, method (2) relies on anchoring the data we want to tion is the aIR value linked directly to the IRX by A(FUV) = calibrate to other SFRs measurements, i.e. we establish a refer- IRX 2.5log(1+aIR × 10 ). (See Hao et al. 2011, equation 2, where ence SFR against which we can compare the hybrid tracers. One the authors name it aFUV instead of aIR; we used aIR for consis- possibility would be to use the SFR provided by the extinction- tency along this work). For the IRX, they use the definition of corrected Paα line emission (see Calzetti et al. 2007). This line Meurer et al. (1995): IRX = log[L(TIR)/L(FUV)obs]. In order to is only moderately influenced by dust extinction and gives us a see the differences from Hao et al.’s sample and this work, we good measure of the current SFR. The problems related with this have derived our own IRX−β relation for the galaxies that have emission line are due to its faintness and the difficulty of observ- FUV and TIR measurements in the CALIFA mother sample in- ing a large number of nearby galaxies as it is only accessible cluding SF and type-2 AGN host galaxies (see Figure 13). We from space. use the intrinsic FUV−NUV color obtained by Gil de Paz et al. (2007), 0.025 ± 0.049 mag, which is very similar to the one ob- In our case, we are going to use the extinction-corrected Hα tained by Hao et al. (2011), 0.022 ± 0.024 mag. Black line shows SFR tracer measurements obtained for the first time from IFS Hao et al.’s fit and the green line is our own fit. This fit gives us data as a reference. These data are required to obtain a proper a value for the coefficient aIR of 0.33 ± 0.08, in comparison with estimation of the stellar continuum and therefore, to estimate their value aIR = 0.46 ± 0.12. Notice that there is a large disper- a reliable measurement of the ionized-gas dust attenuation via sion in the previous figure, so even taking into account that our the Balmer decrement. Besides, we count with an homogeneous measure of the aIR coefficient is in good agreement with the one large survey that provides us with good statistics on the proper- found by Hao et al. (2011), we trust more on method (2) as we ties of nearby star-forming galaxies. For these reasons (and oth- discuss next. ers explained in Section 4.4), we consider this tracer as a robust estimator of the SFR. Using this method we obtain the updated Article number, page 17 of 36 ) ) 1 1 ) SF (N=164) ) SF (N=164) 1 1 − − − 42.5 1.2 − 42.5 1.2 Type-2 AGN (N=99) yr yr 42.0 0.7 42.0 0.7 41.5 0.2 )] (M 41.5 0.2 )] (M )] (ergs s )] (ergs s cor cor α α cor 41.0 -0.3 cor 41.0 -0.3 α α 40.5 -0.8 40.5 -0.8 40.0 aIR = 0.015 -1.3 40.0 aIR = 0.018 -1.3 log [L(H log [L(H log [SFR(H log [SFR(H 1 40.0 40.5 41.0 41.5 42.0 42.5 43.0 43.5 1 40.0 40.5 41.0 41.5 42.0 42.5 43.0 1σ = 0.17dex1 1σ = 0.14dex1 log [ L(Hαobs) + aIR L(22µm) ] (ergs± s− ) log [ L(Hαobs) + aIR L(22µm) ] (ergs± s− ) 0 · 0 · 1 Residuals − Residuals − 40.0 40.5 41.0 41.5 42.0 42.5 43.0 43.5 40.0 40.5 41.0 41.5 42.0 42.5 43.0 1 1 log [ L(Hα ) + a L(22µm) ] (ergs s− ) log [ L(Hα ) + a L(22µm) ] (ergs s− ) obs IR · obs IR · ) ) 1 1 ) ) 1 SF (N=135) 1 SF (N=135) − − − 42.5 1.2 − 42.5 1.2 Type-2 AGN (N=83) yr yr 42.0 0.7 42.0 0.7 41.5 0.2 )] (M 41.5 0.2 )] (M )] (ergs s )] (ergs s cor cor cor 41.0 -0.3 cor 41.0 -0.3 Hα Hα Hα 40.5 -0.8 Hα 40.5 -0.8 40.0 aIR = 0.0015 -1.3 40.0 aIR = 0.0019 -1.3 log [L( log [L( log [SFR( log [SFR( 1 40.0 40.5 41.0 41.5 42.0 42.5 43.0 43.5 1 40.0 40.5 41.0 41.5 42.0 42.5 43.0 43.5 1σ = 0.17dex1 1σ = 0.15dex1 log [ L(Hαobs) + aIR L(TIR) ] (ergs s±− ) log [ L(Hαobs) + aIR L(TIR) ] (ergs s±− ) 0 · 0 · 1 Residuals − Residuals − 40.0 40.5 41.0 41.5 42.0 42.5 43.0 43.5 40.0 40.5 41.0 41.5 42.0 42.5 43.0 43.5 1 1 log [ L(Hα ) + a L(TIR) ] (ergs s− ) log [ L(Hα ) + a L(TIR) ] (ergs s− ) obs IR · obs IR · Fig. 15: Updated calibrations for the hybrid tracers that combine Hα observed luminosity and infrared emission. The ones with 22 µm luminosity appear at the top while the tracers that used TIR luminosity are shown at the bottom. Galaxies hosting type-2 AGN (orange points) are included on the left panels. Blue points refer to star-forming galaxies. Dashed lines correspond to the 1:1 line. The best fitting aIR coefficients calculated as the median value of the expression 16 are shown for clarity. These aIR values and their corresponding errors appear in table 3. The residuals are computed as the average value of the log[L(Hαcorr)/(L(Hαobs) + aIR×L(IR))], where L(IR) could be 22 µm or TIR, after applying a 4 σ rejection. The hybrid tracers show dispersions lower than in the case of the single-band tracers. These hybrid tracers show the same pattern as the ones in Figure 14, the aIR coefficient decreases when galaxies hosting type-2 AGN are considered. These calibrations and the ones in Figure 14 show that applying an energy balance approximation is a good approach to obtain reliable SFR tracers for integrated measurements of nearby galaxies. Here, we use IFS data for the first time to achieve this goal. calibrations for FUV + 22 µm, FUV + TIR, Hα + 22 µm and Hα stant difference (∼ 9 %) in the coefficients of the infrared term, + TIR hybrid tracers that appear in Table 3. aIR, 3.55 (4.52) in the FUV + 22 µm and 0.33 (0.40) in FUV + TIR tracers with (and without) type-2 AGN host galaxies being The resulting hybrid-tracer calibrations obtained using FUV included. A ∼ 10 % difference also appears when we compare and Hα as observed luminosities are shown in Figures 14 and the aIR coefficients between the Hα + 22 µm, 0.015 (0.018), and 15. The dispersions found for FUV + 22 µm, FUV + TIR, Hα Hα + TIR, 0.0015 (0.0019), calibrators with (and without) type- + 22 µm and Hα + TIR tracers are 0.23 (0.19), 0.22 (0.19), 0.17 2 AGN, respectively. If we compare the ratio between these aIR (0.14) and 0.17 (0.15) dex when type-2 AGN are (not) included, coefficients for the combinations of 22 µm and TIR data with the respectively. The single best-fitting parameter in each of these luminosity ratio expected for infrared SEDs with different inter- plots is the median of the distribution of the coefficients that mul- stellar radiation fields, starlight intensities, dust chemical com- tiply the corresponding infrared luminosity in each galaxy (aIR) position, etc., we would predict the ratio of the energy absorbed to match the SFR based on the extinction-corrected Hα luminos- by dust at λ < or > 4000 (see Figure 2 of Cortese et al. 2008). ity. As explained before, aIR has been obtained using expressions The most optimal models for carrying out such study, those by 16 and 17 and could be found in these plots. The line shown in Draine & Li (2007), assume a specific and fixed shape for the these figures corresponds to the 1:1 relation in SFR. This line interstellar radiation field (the local one) so the effect of opti- corresponds to the 1:1 line also in luminosity in the case of the cal photons is hidden in the variation of the factor γ, which pa- Hα + IR tracers (Figure 15) but takes into account the different rameterize the fraction of dust heated by intense radiation fields. constant for FUV and Hα given in Equations 3 and 4, respec- The comparison of the aIR for 22 µm and aIR for TIR coefficients tively (Figure 14). An interesting result found is a nearly con- Article number, page 18 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample yields a factor of 0.1 between L(22) and L(TIR), which (accord- Regarding the FUV+IR hybrid tracers (four bottom panels ing to Figure 19 of Draine & Li 2007) corresponds to γ = 0.02, in Figure 17), we find that the median values for aIR are more quite independently of the fractional abundance of PAHs. similar between S0/a-Sab and Sb-Sbc galaxies. However, the In Figures 14 and 15 we also show the results of this analy- Sc-Sd-Sm-Irr galaxies still show the highest value for the aIR sis after including (left panels) or excluding (right panels) type-2 coefficient. Table 4 lists the resulting median values and their AGN from the sample in order to establish whether the behavior corresponding errors. of the hybrid calibrators changes in each case. The four hybrid These trends are likely the combination of multiple effects tracers show the same pattern, the aIR coefficient decreases when (especially given the large dispersion in the value of aIR within a galaxies hosting type-2 AGN are considered. As for the case of given subsample), namely: the single-band tracers (see Section 4.5.1) this decrease in the value of the aIR coefficient implies that we need to slightly re- (1) The contribution of obscured AGN to the IR luminosity (both duce the contribution of the infrared emission in type-2 AGN to at 22 µm and in the TIR). This partly explains the fact that match that measured in Hα. This implication means that either the average aIR decreases when type-2 AGN are included (1) galaxies hosting type-2 AGN are emitting more light in the in the sample. The fraction of type-2 AGN is larger within infrared that is not associated to the sites or processes that lead early-type galaxies, so part of the IR luminosity (without an to the Hα emission, both at 22 µm and TIR luminosities, than equivalent extinction-corrected Hα luminosity counterpart) normal star-forming galaxies or (2) the Balmer-corrected Hα lu- is arising from the (obscured) AGN itself. minosity underestimates the actual SFR in these galaxies. (2) A fraction of the SFR (that assumed to be in this case ac- The distribution of the aIR coefficient appears in the his- curately measured using an hybrid tracer with a nominal – tograms of Figure 16 where red dashed lines are referred to its large – value of the aIR coefficient) is missed when using median value. This coefficient has a large dispersion even when the extinction-corrected Hα luminosity. This happens, espe- only star-forming galaxies are studied. In Table 3 we give the re- cially, in early-type spirals so the aIR coefficient decreases in sulting median values of aIR and the corresponding dispersions these objects to compensate for the reduced amount of SFR (measured as the interval that includes 68% of the data points derived from Hα. When the Hα emission missed is exclu- around the median). These dispersions appear as red tick-marks sively due to an obscured AGN we will be in case (1) and at the top panels in Figure 16 while black tick-marks indicate the Hα would be a fair measure of the SFR. standard error of the median computed from the asymptotic vari- (3) There is a fraction of the infrared emission that is due to ance formula (which assumes that the underlying distribution is heating by optical photons. One would expect that this ef- Gaussian) using the previous 1 σ dispersions. These values are fect would be more notorious when the aIR coefficient refers in good agreement with the ones reported in the literature for in- to the TIR band, as optical photons are expected to heat the tegrated measurements of galaxies. Kennicutt et al. (2009) found dust at low temperatures, where the emission at 22 µm should 0.020 ± 0.005 and 0.0024 ± 0.0006 for L(Hα) + a × L(24 µm) IR be small. As discussed above, the value of aIR is smaller for and L(Hα) + aIR × L(TIR), respectively. For the case of the UV S0/a-Sab galaxies, which are galaxies that have older stel- luminosity, Hao et al. (2011) found 3.89 ± 0.15 and 0.46 ± 0.12 lar populations and optically bright bulges. Li et al. (2013) for L(FUV) + aIR × L(25 µm) and L(FUV) + aIR × L(TIR), re- also found that the coefficient that multiply the IR luminosity spectively. in the L(Hα) + a × L(70 µm) hybrid tracer is smaller when In the rest of this section we study the value of the aIR coef- larger apertures around star-forming regions are used. The ficient as a function of galaxy properties in order to get insights authors attribute this effect to the larger associated star for- on the origin of this spread. As we show below, the change in mation timescale and the consequent dust heated by old stel- the aIR coefficient with galaxy properties does not appear only lar populations. when studying nuclear activity but also galaxy morphology, stel- lar mass, color, axial ratio and attenuation. The fact that by removing type-2 AGN we reduce but not eliminate completely the morphological-type dependence of aIR indicates that while (1) appears to have some role, the other pos- a 4.5.3. Morphological-type dependence of IR in hybrid sibilities are also at play. Disentangling the contribution of dif- tracers ferent mechanisms listed above is not easy. In particular, the Given the large number of galaxies in our sample, we can now change observed in aIR when type-2 AGN are excluded from explore the origin of the differences between the various SFR the sample could be also due to a decrease in the number of tracers. In particular, we analyze the origin of the variation of red massive star-forming galaxies in each morphological-type bin, galaxies which are expected to suffer from mechanism (3) the aIR with different galaxy properties. Here we focus on the study of its dependence with galaxy morphology (see Walcher as well. The analysis of the variation of aIR with other prop- erties, mass, color, axial ratio and ionized-gas attenuation will et al. 2014). Figure 17 shows the distribution of the aIR coef- ficient in bins of morphological type. In the four top plots of help us understanding the relative contribution of these mecha- nisms and, therefore, the specific limitations of the different SFR this figure we can see a trend for the median value of the aIR co- efficient (vertical dashed lines) with the galaxy morphology for hybrid tracers. Hα + IR tracers. Star-forming galaxies of early-type, consid- ered here as S0/a, Sa and Sab, have lower median values for aIR 4.5.4. Stellar mass dependence of a in hybrid tracers (red dashed line) than intermediate-type spirals such as Sb and IR Sbc (grey dashed line). The last group of galaxies, Sc-Sd-Sm- Since morphology alone is not able to establish the origin of the Irr, shows the largest median value for the aIR coefficient (blue variation of aIR from galaxy to galaxy and within subsamples, dashed line). When type-2 AGN galaxies are excluded (right we now explore its dependence with stellar mass. We use the panels in Figure 17) the trend is less obvious, mainly because of total stellar masses for the CALIFA galaxies from Walcher et al. a drastic increase in the median aIR of early-type spirals. (2014), Section 6.3. (J. Walcher, priv. comm.). The masses are Article number, page 19 of 36 With Type 2 AGN Without Type 2 AGN 22 + Hα 22 + Hα 55 55 45 45 35 35 25 25 Number Number 15 15 5 5 0 0.02 0.04 0.06 0.08 0.10 0.12 0 0.02 0.04 0.06 0.08 0.10 0.12 aIR aIR TIR + Hα TIR + Hα 25 25 15 15 Number Number 5 5 0 0.002 0.004 0.006 0.008 0.010 0 0.002 0.004 0.006 0.008 0.010 aIR aIR 22 + FUV 22 + FUV 35 35 25 25 15 15 Number Number 5 5 -2 0 2 4 6 8 10 12 14 16 18 -2 0 2 4 6 8 10 12 14 16 18 aIR aIR TIR + FUV TIR + FUV 35 35 25 25 Number 15 Number 15 5 5 -0.4 -0.2 0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 1.6 -0.4 -0.2 0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 1.6 aIR aIR Fig. 16: Histograms showing the distribution of the aIR coefficient values obtained for the different hybrid tracers. aIR is computed using the expressions 16 and 17. The red dashed line corresponds to the median value of this coefficient. The red tick-marks shown at the top refers to the 1 σ dispersions measured as the interval that includes 68% of the data points around the median quoted in Table 3 while black tick-marks indicate the standard error of the median computed from the asymptotic variance formula using these 1 σ dispersions. Type-2 AGN galaxies are excluded from the histograms at right. Article number, page 20 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample With Type 2 AGN Without Type 2 AGN 22 + Hα 22 + Hα S0/a-Sab S0/a-Sab Sb-Sbc Sb-Sbc Sc-Sd-Sm-Irr Sc-Sd-Sm-Irr 20 20 15 15 Number Number 10 10 5 5 0 0.02 0.04 0.06 0.08 0.10 0.12 0 0.02 0.04 0.06 0.08 0.10 0.12 aIR aIR TIR + Hα TIR + Hα S0/a-Sab S0/a-Sab Sb-Sbc Sb-Sbc Sc-Sd-Sm-Irr Sc-Sd-Sm-Irr 15 15 10 10 Number Number 5 5 0 0.002 0.004 0.006 0.008 0.010 0 0.002 0.004 0.006 0.008 0.010 aIR aIR 22 + FUV 22 + FUV S0/a-Sab S0/a-Sab Sb-Sbc Sb-Sbc Sc-Sd-Sm-Irr Sc-Sd-Sm-Irr 15 15 10 10 Number Number 5 5 -2 0 2 4 6 8 10 12 14 16 18 -2 0 2 4 6 8 10 12 14 16 18 aIR aIR TIR + FUV TIR + FUV S0/a-Sab S0/a-Sab Sb-Sbc Sb-Sbc Sc-Sd-Sm-Irr Sc-Sd-Sm-Irr 15 15 10 10 Number Number 5 5 -0.2 0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 1.6 -0.2 0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 1.6 aIR aIR Fig. 17: Four left panels: Histograms showing the values of the coefficient that multiply the IR luminosity, aIR, in the hybrid tracers using bins of morphological types. Early-type star-forming galaxies, considered here as S0/a, Sa and Sab, are shown in red, intermediate-type spirals such as Sb and Sbc appear in grey and Sc-Sd-Sm-Irr galaxies are represented in blue. Vertical dashed lines correspond to the median value of each galaxy group. Black top marks show the median value for all the galaxies as in Figure 16. There is a clear trend with the morphological type, late-type galaxies need a higher value of the aIR coefficient than early-type galaxies. This trend could be explained in terms of the contribution of an obscured AGN, a missing fraction of the Hα extinction-corrected SFR or the heating by optical photons. Four right panels: Same histograms as before but removing the type-2 AGN. There is still a trend with the morphological type although less obvious that in the previous case. Article number, page 21 of 36 With Type 2 AGN Without Type 2 AGN 22 + Hα 22 + Hα log[M?] > 10.5 log[M?] > 10.5 10.0 < log[M?] < 10.5 10.0 < log[M?] < 10.5 25 log[M?] < 10.0 25 log[M?] < 10.0 20 20 15 15 Number Number 10 10 5 5 0 0.02 0.04 0.06 0.08 0.10 0.12 0 0.02 0.04 0.06 0.08 0.10 0.12 aIR aIR TIR + Hα TIR + Hα log[M?] > 10.5 log[M?] > 10.5 10.0 < log[M?] < 10.5 10.0 < log[M?] < 10.5 20 log[M?] < 10.0 20 log[M?] < 10.0 15 15 Number 10 Number 10 5 5 0 0.002 0.004 0.006 0.008 0.010 0 0.002 0.004 0.006 0.008 0.010 aIR aIR 22 + FUV 22 + FUV log[M?] > 10.5 log[M?] > 10.5 10.0 < log[M?] < 10.5 10.0 < log[M?] < 10.5 log[M?] < 10.0 log[M?] < 10.0 15 15 10 10 Number Number 5 5 -2 0 2 4 6 8 10 12 14 16 18 -2 0 2 4 6 8 10 12 14 16 18 aIR aIR TIR + FUV TIR + FUV log[M?] > 10.5 log[M?] > 10.5 10.0 < log[M?] < 10.5 10.0 < log[M?] < 10.5 log[M?] < 10.0 log[M?] < 10.0 15 15 10 10 Number Number 5 5 -0.2 0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 1.6 -0.2 0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 1.6 aIR aIR Fig. 18: Left four panels: Frequency histograms of aIR for different hybrid tracers as a function of stellar mass. Massive galaxies (log[M∗/M ]>10.5) appear in red, intermediate-mass galaxies (10 of Hα emission being completely absorbed either at the AGN or 2 SF in circumnuclear star formation or to correlations between nu- Type-2 AGN 0.04 < z < 0.1 Elbaz+07 clear activity and other properties, besides mass, such as mor- phological type or environment. The latter is related with the fact 1 that the presence of an AGN might impact the internal evolution )] of the galaxy quenching the SF by feedback mechanisms (for a corr complete review on this topic see Alexander & Hickox 2012). α The analysis of potential effects of the AGN on the current level 0 of star formation at fixed mass (e.g. preference of type-2 AGN for the Green Valley; see Kauffmann et al. 2003; Sánchez et al. 2004) is beyond the scope of this paper. log[SFR(H 1 − 4.5.5. Color dependence of aIR in hybrid tracers 2 − We address here the dependence of the aIR coefficient with the 8.0 8.5 9.0 9.5 10.0 10.5 11.0 11.5 color of the integrated stellar population (as traced by the global log[M?/M ] SDSS g − r color; see Walcher et al. 2014). Figure 20 shows the distribution of aIR coefficients as a function of the galaxies Fig. 19: Correlation between the extinction-corrected Hα SFR and the g−r color in the case of the Hα+22 µm (left panel) and Hα+TIR total stellar mass of the galaxies. Solid line represents the fit of Elbaz (right panel) hybrid tracers. In these plots, type-2 AGN galax- et al. (2007) with a slope of 0.77 for galaxies in the 0.04 Article number, page 23 of 36 SF 20 SF 20 Type-2 AGN Type-2 AGN 15 15 10 10 Galaxies Galaxies 5 5 # # 0 0 SF SF Type-2 AGN Type-2 AGN 0.08 0.008 0.006 m) 0.06 µ (TIR) (22 0.04 0.004 IR a IR a 0.02 0.002 0.00 0.000 0.2 0.3 0.4 0.5 0.6 0.7 0.8 0.9 1.0 0.2 0.3 0.4 0.5 0.6 0.7 0.8 0.9 1.0 g-r g-r Fig. 20: Left panel: Distribution of aIR coefficients as a function of the galaxies g − r SDSS color for the Hα+22 µm hybrid tracer. Orange points show type-2 AGN galaxies while blue points represent star-forming galaxies. Filled contours represent the 1 σ dispersion after applying a 5 σ rejection around the mean value expressed as a blue (orange) solid line for the SF (type-2 AGN) galaxies. The corresponding histogram with the distribution of the number of galaxies for each g − r SDSS color is plotted on the top for reference applying the same color-coding. Right panel: Same as the left panel showing the Hα+TIR hybrid tracer instead. There is a clear offset between the star-forming and type-2 AGN host galaxies with the aIR coefficient at any g − r SDSS color. 20 SF SF Type-2 AGN 15 Type-2 AGN 15 10 10 Galaxies 5 Galaxies 5 # # 0 0 SF SF Type-2 AGN Type-2 AGN 0.08 0.008 0.006 m) 0.06 µ (TIR) (22 0.04 0.004 IR a IR a 0.02 0.002 0.00 0.000 0.0 0.2 0.4 0.6 0.8 1.0 0.0 0.2 0.4 0.6 0.8 1.0 axial ratio axial ratio Fig. 21: Left panel: Distribution of aIR coefficients as a function of the galaxies axial ratio for the Hα+22 µm hybrid tracer. Orange points show type-2 AGN galaxies while blue points represent star-forming galaxies. Filled contours represent the 1 σ dispersion after applying a 5 σ rejection around the mean value expressed as a blue (orange) solid line for the SF (type-2 AGN) galaxies. The corresponding histogram with the distribution of the number of galaxies for each axial ratio is plotted on the top for reference applying the same color-coding. Right panel: Same as the left panel showing the Hα+TIR hybrid tracer instead. SF objects with low axial ratios (where highly-inclined disk galaxies would be located) show similar aIR values as the rest of the galaxies. Lower aIR would be expected if a fraction of the Hα emission is completely obscured. normal galaxies. They used the 24 µm flux as a tracer of the sis of obscuration effects is the subject of Section 4.5.7 below in obscured emission due to SF and the uncorrected Hα flux as a which efforts had been made to further explore this issue. tracer of the unobscured portion (the same way as we use our 22 µm + Hα hybrid tracer along this work, but we compute inte- grated measurements of galaxies). These authors conclude that 4.5.7. Attenuation dependence of aIR in hybrid tracers the fraction of highly obscured regions in normal, star forming In order to determine whether local obscuration effects in Hα disk galaxies is small on 500 pc scales. They are more luminous might be behind the decrease of aIR in galaxies of specific types, and tend to be closer to the center of the host galaxy. The analy- masses or colors (as this coefficient should be reduced to com- pensate by the SFR missed in Hα) we finally analyze its varia- Article number, page 24 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample log[M?] < 10.0 log[M?] < 10.0 0.12 10.0 < log[M?] < 10.5 0.12 10.0 < log[M?] < 10.5 log[M?] > 10.5 log[M?] > 10.5 0.10 0.10 0.08 m) m) 0.08 µ µ 0.06 (22 (22 0.06 IR IR a a 0.04 0.04 0.02 0.02 0.00 0.00 1 0 1 2 1 0 10− 10 10 10− 10− 10 A(Hα) (mag) A(FUV)IRX(Hao) (mag) 0.012 0.012 log[M?] < 10.0 log[M?] < 10.0 10.0 < log[M?] < 10.5 10.0 < log[M?] < 10.5 0.010 log[M?] > 10.5 0.010 log[M?] > 10.5 0.008 0.008 0.006 0.006 (TIR) (TIR) IR a IR 0.004 a 0.004 0.002 0.002 0.000 0.000 10 1 100 101 2 1 0 − 10− 10− 10 A(FUV) (mag) A(Hα) (mag) IRX(Hao) Fig. 23: Top panel: Variation of the a coefficient with the FUV atten- Fig. 22: Top panel: Variation of the a coefficient with the Hα at- IR IR uation derived using the IRX from Hao et al. (2011) for the Hα+22 µm tenuation derived using the Balmer decrement for the Hα+22 µm hy- hybrid tracer. Same symbol-coding as in Figure 22. Bottom panel: brid tracer. Black squares show galaxies with stellar masses lower than Same as previous panel this time for the a coefficient that corresponds log[M ] < 10.0, grey stars represent galaxies with stellar masses in the IR ? to the Hα+TIR hybrid tracer. range of 10.0 < log[M?] < 10.5 and, finally, open circles are for the most massive galaxies with log[M?] > 10.5. Bottom panel: Same as previous panel this time for the aIR coefficient that corresponds to the Hα+TIR hybrid tracer. FUV attenuation applying the IRX given by Hao et al. (2011) is shown in Figure 23 for comparison. Except for this tail at low attenuation, A(Hα)<0.2 mag, we find no correlation between the two parameters. Should a signif- tion as a function of ionized-gas attenuation. This attenuation is icant fraction of the SFR being missed when the extinction cor- derived from the Hα/Hβ Balmer decrement as described in Sec- rected Hα luminosity is used, one would expect to find a clear decrease in aIR as the ionized-gas attenuation gets larger. Only tion 3.1.4. Figure 22 shows the variation of the aIR coefficient with the attenuation measured in magnitudes in Hα (the use of when low-mass galaxies are analyzed separately they seem to log scale in the abscissa is justified by the large concentration show a decline in their aIR values above A(Hα)=0.4 mags, al- of points at low attenuations). The most remarkable feature in though with some discrepant points at A(Hα) > 1 mag. This this plot is that there is a number of galaxies with low global decline in aIR might be due to the fact that only in some of these naturally low-metallicity galaxies high attenuations are due ionized-gas attenuations that show very small values of aIR. We interpret this as consequence of dust emission that is caused by to the presence of active nuclear star formation events. De- the heating of photons different from those arising in sites of cur- spite that, the average and scatter obtained for the aIR coefficient rent star formation. In galaxies where the attenuation derived in in these galaxies are not very different from those obtained at Hα is compatible with no attenuation even a small amount of higher masses neither aIR reaches very low values. dust emission would lead to a null value for aIR, which results in the number of galaxies with low attenuations and low values 5. Summary and conclusions of aIR seen in Figure 22. These galaxies with very small val- ues of aIR indicate that, at this level, we are in the limit where In this work, we present the analysis of the Star Formation Rate A(Hα) can be properly derived, given the low global ionized-gas in a sample of 380 galaxies from the diameter-limited CALIFA attenuations found. The variation of the aIR coefficient using the survey. A total of 272 galaxies shows detected emission in both Article number, page 25 of 36 Hβ and Hα and are listed in Table 1 for reference. The avail- mass, color, and attenuation among the galaxies in the sam- ability of wide-field IFS for all the galaxies in the sample is a ple, there is no single physical property that can by itself major advantage over other techniques. Using IFS data we can explain the entire variation in aIR from galaxy to galaxy. recover the flux in galaxies with low equivalent widths and sep- (4) The analysis of the dependence of aIR with galaxy proper- arate Hα and the [NII] without assuming a [NII]/Hα ratio avoid- ties indicates that galaxies hosting type-2 AGN tend to re- ing problems associated with narrow-band imaging or long-slit duce the median value of aIR, likely due to the contribution spectroscopy. It also ensures a proper determination of the un- of obscured AGN to the infrared emission. The fact that aIR derlying stellar continuum and, consequently, of the extinction- does not show a particularly low value at high ionized-gas corrected Hα luminosity. attenuations nor low axial ratios, suggests that obscured star We have combined the aperture-corrected Hα measurements formation is, comparatively, playing a minor role. Part of from CALIFA with those measured in other bands that are also the dependence of the median value of aIR with the mor- used to estimate the SFR, including luminosity measurements in phological type disappears once the AGN contribution is re- the UV from GALEX (200 galaxies), 22 µm from WISE (265 moved although early spirals still show a somewhat lower aIR galaxies) and TIR luminosities from WISE+IRAS+AKARI than intermediate- and late-type spirals. This behavior, also SED fitting (221 galaxies). present when comparing massive with less massive systems, We first compare the extinction-corrected Hα SFR with mea- can be explained in part as due to the enhanced contribution surements from single-band (FUV, 22 µm and TIR) and hybrid- of optical photons to the heating of the dust in both early- tracers (Hα+22 µm, Hα+TIR, FUV+22 µm, FUV+TIR). In this type spirals and massive systems. part of the paper, we use recent compilations of SFR recipes by Calzetti (2013). The good correlation between the SFR sur- These conclusions will allow us to make use of the CALIFA face density obtained with extinction-corrected Hα tracer and IFS data to explore the distribution of the SFR with spatial res- FUV+22 µm hybrid tracer guarantees that potential linear cor- olution in a future work (see Catalán-Torrecilla et al. 2014). We relations between different SFR tracers (some of them not re- emphasize that the impact of potential differences in the selec- solved, such as those relying on TIR measurements) are not tion criteria should be addressed carefully when extrapolating driven by scaling effects and that global values of the SFR can these results to other samples of galaxies and, particularly, to be used reliably. Our results indicate that, overall, the extinction- other redshifts. corrected Hα luminosity (once underlying stellar absorption and Acknowledgements. This study makes uses of the data provided by the Calar dust-attenuation effects are properly accounted for) matches the Alto Legacy Integral Field Area (CALIFA) survey (http://califa.caha.es). CAL- IFA is the first legacy survey being performed at Calar Alto. The CALIFA col- SFR obtained from hybrid tracers combining the observed FUV laboration would like to thank the IAA-CSIC and MPIA-MPG as major part- or Hα and the IR (22 µm and TIR) luminosities with dispersions ners of the observatory, and CAHA itself, for the unique access to telescope around ∼0.20 dex being found. In the case of the comparison time and support in manpower and infrastructures. The CALIFA collabora- with single-band tracers we conclude (1) that the use of IR mea- tion thanks also the CAHA staff for the dedication to this project. C. C.-T. ∼ −1 thanks the support of the Spanish Ministerio de Educación, Cultura y Deporte by surements clearly underestimates the SFR below 1 M yr and means of the FPU fellowship program. The authors also thank the support from (2) the large uncertainty in the correction for attenuation when the Plan Nacional de Investigación y Desarrollo funding programs, AYA2012- only FUV−NUV color (similar to the UV slope, β) information 30717 and AyA2013-46724P, of Spanish Ministerio de Economía y Competi- is available. This factor introduces a very large scatter, particu- tividad (MINECO). P.G. P-G. acknowledges support from the AYA2012-30717 larly at SFR>5 M yr−1, where the β-corrected FUV luminosity and AYA2012-31277. J.I.P. acknowledges financial support from the Spanish MINECO under grant AYA2010-21887-C04-01 and from Junta de Andalucía also tend to underestimate the SFR. This prevents the use of the Excellence Project PEX2011-FQM7058. R.A. M. is funded by the Spanish pro- UV luminosity alone as a SFR tracer. gram of International Campus of Excellence Moncloa (CEI). MAPT acknowl- edges support from the Spanish MINECO through grant AYA2012-38491-C02- 02. AdO acknowledge financial support from the Spanish grant AYA2013- We also provide a new set of single-band calibrators anchored 42227-P. Support for LG is provided by the Ministry of Economy, Develop- to the extinction-corrected Hα luminosities. The values for these ment, and Tourism’s Millennium Science Initiative through grant IC120009, awarded to The Millennium Institute of Astrophysics, MAS. LG acknowledges coefficients appear in Table 2. In the case of the hybrid calibra- support by CONICYT through FONDECYT grant 3140566. JMG acknowl- tors we determine the best (median) fit for the coefficient that edges support from the Fundação para a Ciência e a Tecnologia (FCT) through weights the amount of IR luminosity reprocessed by dust, aIR. the Fellowship SFRH/BPD/66958/2009 from FCT (Portugal) and POPH/FSE We assume an energetic balance and calculate the aIR coeffi- (EC) by FEDER funding through the program Programa Operacional de Fac- cients for different combinations of observed (UV or Hα) and tores de Competitividade (COMPETE). JMG also acknowledges support by FCT under project FCOMP-01-0124-FEDER-029170 (Reference FCT PTDC/FIS- dust-reprocessed (22 µm or TIR) SFR contributions anchored to AST/3214/2012), funded by FCT-MEC (PIDDAC) and FEDER (COMPETE). the extinction-corrected Hα luminosities. These values appear This research has made use of the NASA/IPAC Extragalactic Database (NED) in Table 3 and are calculated with and without galaxies hosting which is operated by the Jet Propulsion Laboratory, California Institute of Tech- type-2 AGN being considered. nology, under contract with the National Aeronautics and Space Administration. This publication makes use of data products from the Wide-field Infrared Survey (1) This analysis allows us to provide, for the first time, a set Explorer, which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by of hybrid calibrations for different morphological types and the National Aeronautics and Space Administration. This research has made masses. These are particularly useful in case that the sample use of the NASA/ IPAC Infrared Science Archive, which is operated by the Jet to be analyzed shows a different bias in terms of morphology Propulsion Laboratory, California Institute of Technology, under contract with or, more commonly, luminosity or stellar mass (see Tables 4 the National Aeronautics and Space Administration. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participat- and 5). ing Institutions, the National Science Foundation, the U.S. Department of En- (2) We also study the dependence of this coefficient not only ergy, the National Aeronautics and Space Administration, the Japanese Monbuk- with morphological type and mass but also with color (SDSS agakusho, the Max Planck Society, and the Higher Education Funding Council g − r), axial ratio and ionized-gas attenuation. for England. The SDSS Web Site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The (3) The distributions of aIR values (for each of the hybrid trac- Participating Institutions are the American Museum of Natural History, Astro- ers) are quite wide in all cases. While part of the spread physical Institute Potsdam, University of Basel, University of Cambridge, Case can be attributed to changes in morphological type, stellar Western Reserve University, University of Chicago, Drexel University, Fermi- Article number, page 26 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample lab, the Institute for Advanced Study, the Japan Participation Group, Johns Hop- Gordon, K. D., Clayton, G. C., Witt, A. N., & Misselt, K. A. 2000, ApJ, 533, kins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute 236 for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chi- Hao, C.-N., Kennicutt, R. 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I., et al. 1995, ApJ, 439, 604 Yi, S., Demarque, P., & Oemler, Jr., A. 1997, ApJ, 486, 201 Article number, page 27 of 36 1 Departamento de Astrofísica y CC. de la Atmósfera, Universidad Complutense de Madrid, E-28040, Madrid, Spain e-mail: [email protected] 2 Instituto de Astrofísica de Andalucía-CSIC, Glorieta de la As- tronomía, 18008 Granada, Spain 3 Estación Experimental de Zonas Áridas (CSIC), Ctra. de Sacra- mento s/n, La Cañada, Almería, Spain 4 Instituto de Astronomía,Universidad Nacional Autonóma de Méx- ico, A.P. 70-264, 04510, México,D.F. 5 Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge, CB3 0HA, United Kingdom 6 CEI Campus Moncloa, UCM-UPM, Departamento de Astrofísica y CC. de la Atmósfera, Universidad Complutense de Madrid, E- 28040, Madrid, Spain 7 Leibniz-Institut für Astrophysik Potsdam (AIP), An der Sternwarte 16, D-14482 Potsdam, Germany 8 European Southern Observatory, Karl-Schwarzschild-Str. 2, 85748 Garching b. München, Germany 9 Instituto de Cosmologia, Relatividade e Astrofísica - ICRA, Cen- tro Brasileiro de Pesquisas Físicas, Rua Dr.Xavier Sigaud 150, CEP 22290-180, Rio de Janeiro, RJ, Brazil 10 European Southern Observatory, Casilla 19001, Santiago 19, Chile 11 Sydney Institute for Astronomy, School of Physics A28, University of Sydney, NSW2006, Australia 12 Astronomical Institute of the Ruhr-University Bochum Universi- taetsstr, 150, 44801 Bochum, Germany 13 Millennium Institute of Astrophysics, Universidad de Chile, Casilla 36-D, Santiago, Chile 14 Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile 15 Instituto de Astrofísica e Ciências do Espaço, Universidade do Porto, CAUP, Rua das Estrelas, PT4150-762 Porto, Portugal 16 Australian Astronomical Observatory, PO Box 915, North Ryde, NSW 1670, Australia 17 Department of Physics and Astronomy, Macquarie University, NSW 2109, Australia 18 GEPI, Observatoire de Paris, CNRS UMR8111, Université Paris Diderot, Place Jules Janssen, 92190 Meudon, France 19 Centro de Estudios de la Física del Cosmos de Aragón (CEFCA), 44001 Teruel, Spain 20 Departamento de Física Teórica, Universidad Autónoma de Madrid, Cantoblanco, E28049, Spain Article number, page 28 of 36 C. Catalán-Torrecilla et al.: Star Formation in the Local Universe from the CALIFA sample ) 1 − 0.94 1.55 0.19 0.33 0.38 0.32 0.18 0.13 0.69 0.38 0.34 0.45 0.56 0.43 0.05 0.41 0.27 0.56 0.61 0.24 0.15 0.09 0.12 0.37 0.35 0.31 0.84 0.14 0.37 0.27 0.18 0.66 0.22 0.17 0.29 0.34 m) ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± µ erg s 42 luminosity 10 α × )( 1 − 1.541.520.48 63.64 1.22 8.56 1.90 2.59 11.11 0.74 8.93 0.12 7.22 0.43 1.34 0.440.42 12.08 0.83 10.01 1.04 4.84 0.09 10.89 0.80 19.23 0.26 0.65 0.89 16.33 1.27 10.20 0.48 23.33 21.94 0.42 7.44 0.31 1.95 1.041.64 9.07 0.97 19.03 0.24 10.01 0.27 17.49 0.82 4.35 6.53 5.07 0.661.800.24 20.55 0.33 9.33 0.95 5.58 5.03 14.16 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 43 10 × flux uncertainty only, an additional ective temperature and therefore on ff α N L(TIR) L(22 − ) (AB mag) ( 10 % than what is inferred from Spitzer IRS and 1 )F − ∼ 0.200.960.73 0.770.21 0.410.36 0.24 45.41 0.04 13.16 0.310.14 2.19 0.580.12 0.48 11.22 0.22 0.520.57 8.03 0.560.05 0.620.02 . . . 10.02 0.26 0.55 10.06 0.13 1.050.67 . 0.34 . 1.59 . 12.49 0.25 0.46 15.40 0.520.06 0.76 0.56 3.57 19.27 0.17 0.670.04 9.37 20.59 0.48 0.58 23.13 0.12 0.490.18 0.830.10 0.430.15 1.86 1.1 0.480.12 8.42 . 0.87 . . 15.24 0.45 11.77 . . . 15.63 0.31 1.40 3.30 1.92 5.28 0.180.780.13 0.630.28 0.49 0.24 20.08 0.75 5.22 . . . 5.31 1.88 obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 42 10 × )( 1 − ) L(NUV 0.150.740.18 32.50 38.95 0.24 23.87 0.370.24 34.45 9.74 0.27 6.78 0.530.43 11.74 19.28 0.77 5.92 0.030.03 46.15 0.25 7.90 0.74 3.80 1.00 42.06 0.40 21.66 36.24 0.43 40.05 0.190.50 10.06 0.24 4.61 0.21 6.12 0.17 38.85 0.26 3.77 0.72 29.31 15.52 0.14 7.88 20.21 0.440.19 5.71 0.16 31.57 0.62 21.66 7.47 obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 42 10 × ] L(FUV )( 1 0.920.261.13 26.70 19.14 0.080.290.11 5.71 4.36 ...... 0.10 12.31 . . .0.36 2.85 0.530.45 1.41 1.37 0.23 18.84 0.51 6.96 5.21 ...... 6.13 2.04 15.99 0.48 25.89 0.050.382.77 7.27 0.27 11.51 0.35 3.34 0.320.04 27.81 0.63 3.00 0.45 2.78 .0.96 . . 27.53 0.82 13.41 0.04 21.64 0.09 21.61 0.14 5.90 0.39 . . . 2.93 . . . 26.14 . . .0.23 4.96 . . 15.19 . . . . 5.56 0.220.710.35 3.20 0.30 20.10 0.111.09 17.37 3.74 ...... 9.42 . 16.75 corr − α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± yr 1.1. Note that galaxies UGC 00987 and UGC 03973 are classified as type-1 AGN within m filter are systematically brighter by / µ ) (M 1 − ) SFR[H 1.840.93 5.27 8.67 1.49 0.34 3.70 0.56 2.93 1.60 0.44 1.82 1.17 0.61 0.37 3.43 0.43 1.54 1.54 3.20 15.17 0.73 0.90 0.25 2.11 0.06 1.87 0.79 0.23 0.50 3.73 1.39 2.69 1.04 5.71 1.67 4.84 0.39 0.38 0.66 0.55 0.46 0.78 0.94 0.56 0.66 2.31 0.38 2.06 0.65 3.20 0.65 2.63 4.51 1.29 0.29 1.80 1.35 1.15 12.01 0.881.87 1.17 0.75 4.10 0.40 2.07 0.66 1.79 6.26 0.61 3.92 0 that also appeared in this table). The errors listed here use the H ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs erg s = α ) 40 α 10 m (aperture photometry) luminosity. As explained in Section 3.3, it has been pointed out by several authors × luminosities, col. (8), and SFRs, col. (9) and in the attenuations used in the paper; (8) Observed H µ α SFR values; (10) Observed FUV luminosity; (11) Observed NUV luminosity; (12) UV color in AB system; (13) TIR ) L(H α 0.070.030.02 42.82 0.01 17.8 0.07 33.83 0.02 34.07 0.09 7.18 0.55 8.92 0.06 6.72 0.10 7.88 0.08 18.7 0.05 25.9 0.06 13.05 0.01 35.61 0.07 14.03 0.02 3.58 0.02 40.79 0.06 27.27 0.11 45.09 0.06 47.0 0.02 6.99 0.06 10.04 0.07 8.28 0.04 7.69 0.04 27.51 0.04 10.74 0.06 32.59 0.04 21.66 0.05 12.06 19.75 21.04 0.03 76.62 0.070.050.04 11.47 0.03 40.99 0.11 25.42 0.04 19.78 9.57 28.22 α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± m; (14) 22 µ 1000 22 0.87 1513 0.75 0.48 2033 0.0 0.0 1814 2.57 15 0.25 34 0.08 11 0.96 59 0.17 13 0.55 17 0.64 48 0.9 0.68 10 0.46 2011 0.86 15 0.72 12 0.55 0.16 1023 0.67 0.65 1310 0.54 22 0.16 1.01 − 5.9 0.46 2.49.6 0.11 0.94 9.2 0.44 6.04.4 0.0 3.1 0.0 3.4 0.58 0.3 7.29.3 1.35 0.63 7.1 0.43 140 1.14 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ) A(H ± cgs) (mag) ( obs 16 α − (H 10 f × flux using an elliptical aperture with a semi-major axis of 36 arcsec; (7) Attenuation obtained using the Balmer decrement. As explained in Section a / α ratio of 2.86 (although lower values are also physically possible in HII regions, depending on the electron density, e β H / α m luminosity values given in column (14) have to be multiplied by 1 µ (Mpc) ( This table is available in its entirety in an online version. A portion is shown here for reference regarding its form and content. flux have been taken into account for the errors in the H ects; (9) Balmer-corrected and aperture-corrected H β ff (1) ID CALIFA identifier; (2) Galaxy name; (3) Morphological type; (4) Distance in Mpc calculated from redshift; (5) Minor-to-major axis ratio as given by the RC3 catalog, i.e. measured m data. For that reason, the 22 µ ID name type dist b (1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12) (13) (14) 1011 NGC 003612 UGC 0013914 UGC 0014815 UGC 0031216 ESO539-G014 Sb19 Scd MCG-02-02-04020 Sc ESO540-G00321 Scd 87.09 Sd 56.45 Scd NGC 016022 NGC 0.62 0165 0.46 101.16 59.8823 Sb NGC 50.38 0169 62.1325 0.26 NGC 0.1 0171 0.2826 0.46 NGC 47.63 0180 Sa 2209.2 27 NGC 1933 0192 Sb28 0.4 NGC 0216 Sab 2379.8 30 255.4 5192 75.09 NGC 0214 Sb31 5977 84.19 NGC 0.57 0237 Sb33 66.2 NGC 0.83 0234 Sab38 56.14 1560 NGC 0.25 0257 Sd39 75.65 NGC 0.87 0447 Sbc 60.1740 NGC 0.81 1770.5 0444 Sc41 443 0.47 UGC 22.3 00809 Sc 64.6642 UGC 1160 00841 Sc43 0.34 0.74 NGC 2358 0477 59.5 Sa45 63.73 IC 2267 1683 Scd48 0.57 2533 75.17 Scd NGC 049649 1.0 79.94 Sbc NGC 0.71 0523 69.1852 5111 6034 60.04 UGC 0.98 00987 Sbc53 0.31 79.83 NGC 0551 0.15 4650 UGC 01057 Sb Scd 83.83 8269 0.2 5054 Sd 634.8 0.54 Sa 973.5 85.95 69.36 980.5 Sbc 0.54 69.53 Sc 0.46 725.9 66.53 0.29 74.37 1598 0.3 90.68 0.44 3039.2 1829.7 0.32 1629 1341 2132 1716 23 UGC5 00005 NGC7 7819 IC8 1528 UGC9 00036 Sbc NGC 0001 NGC 0023 Sc 104.31 0.51 Sab Sbc 71.47 Sbc 89.82 0.76 Sb 54.77 0.46 64.92 0.48 2677 0.76 65.26 2617.2 0.65 3146 840 4068 12029 1 IC 5376 Sb 71.97 0.17 720 24 our sample. As a consequence, they have been excluded from the SFR analysis. (Wright et al. 2010; Jarrett et al. 2013; Brown et al. 2014) that SF galaxies measured with the WISE 22 the chemical abundance. This leads7% to uncertainty a in number the of H corrected galaxies for for aperture which e we assumedluminosity A(H defined as the total infrared emission in the range 8 Table 1: in the D25 B-band isophote;3.1.4, we (6) have Observed-H assumed an intrinsic H Article number, page 29 of 36 Table 2: Values of the a, a0 and b coefficients for the calibration of single-band tracers: Linear and non-linear fits, as explained in Section 4.5.1. The numbers in brackets refer to the number of galaxies used in each case (same notation applied in Tables 3, 4 and 5). The luminosities are −1 −1 expressed in erg s and the values of the SFR are in M yr . Single-band Tracers Without type-2 AGN With type-2 AGN × +2.1 × −43 +1.6 × −43 SFR = a L(22 µm) a = 3.0−0.4 10 [164] a = 2.8−0.6 10 [263] × +1.7 × −44 +1.6 × −44 SFR = a L(TIR) a = 2.8−0.5 10 [135] a = 2.3−0.5 10 [218] SFR = a0 × [L(22 µm)/1043]b (log scale fit) a0 = 2.12 ± 0.18 ; b = 0.733 ± 0.053 a0 = 1.86 ± 0.11 ; b = 0.702 ± 0.039 SFR = a0 × [L(22 µm)/1043]b (curve_fit) a0 = 2.60 ± 0.12 ; b = 0.697 ± 0.051 a0 = 2.38 ± 0.12 ; b = 0.564 ± 0.043 SFR = a0 × [L(TIR)/1043]b (log scale fit) a0 = 0.403 ± 0.032 ; b = 0.720 ± 0.042 a0 = 0.359 ± 0.028 ; b = 0.719 ± 0.038 SFR = a0 × [L(TIR)/1043]b (curve_fit) a0 = 0.350 ± 0.063 ; b = 0.845 ± 0.064 a0 = 0.285 ± 0.054 ; b = 0.881 ± 0.061 Table 3: Values of the aIR coefficients for the calibration of hybrid tracers: GLOBAL VALUES (see Section 4.5.2 for a detailed explanation) The errors quoted here are the 1 σ dispersions measured as the interval that includes 68% of the data points around the median and correspond with the spread√ of the histograms in Figure 16. Note that the standard error of the median, computed from the asymptotic variance formula as 1.253×σ/ N, where σ is referred to the values listed here and N is the number of galaxies shown in brackets, decreases these errors considerably (black tick-marks shown at the top in Figure 16). The luminosities in these expressions are in erg s−1 and the values of the SFR are expressed in −1 M yr . Hybrid Tracers Without type-2 AGN With type-2 AGN × −42 × +0.018 +0.018 SFR = 5.5 10 [L(Hα)obs + aIR L(22 µm)] 0.018−0.006 [164] 0.015−0.006 [263] × −42 × +0.0015 +0.0016 SFR = 5.5 10 [L(Hα)obs + aIR L(TIR)] 0.0019−0.0005 [135] 0.0015−0.0006 [218] × −44 × +3.55 +3.38 SFR = 4.6 10 [L(FUV)obs + aIR L(22 µm)] 4.52−1.14 [113] 3.55−0.95 [187] × −44 × +0.33 +0.29 SFR = 4.6 10 [L(FUV)obs + aIR L(TIR)] 0.40−0.09 [94] 0.33−0.07 [156] Table 4: Values of the aIR coefficients for the calibration of hybrid tracers: BY MORPHOLOGICAL TYPE (see Sections 4.5.2 and 4.5.3 for a −1 detailed explanation). The recipes to compute the SFR in M yr shown on the left column are the same as in Table 3. The luminosities are in units of erg s−1. Without type-2 AGN With type-2 AGN Hybrid Tracers S0/a-Sab Sb-Sbc Sc-Irr S0/a-Sab Sb-Sbc Sc-Irr × +0.037 +0.016 +0.016 +0.026 +0.017 +0.015 L(Hα)obs + aIR L(22 µm) 0.010−0.005 [21] 0.022−0.007 [76] 0.021−0.009 [67] 0.006−0.006 [58] 0.017−0.005 [131] 0.019−0.008 [74] × +0.0029 +0.0013 +0.0012 +0.0022 +0.0012 +0.0013 L(Hα)obs + aIR L(TIR) 0.0014−0.0006 [16] 0.0020−0.0006 [68] 0.0022−0.0006 [51] 0.0008−0.0004 [44] 0.0016−0.0004 [116] 0.0018−0.0005 [58] × +2.76 +4.52 +3.97 +4.37 +2.90 +4.61 L(FUV)obs + aIR L(22 µm) 4.17−0.86 [14] 4.67−0.88 [51] 4.75−1.37 [48] 2.55−0.46 [39] 3.65−0.78 [93] 4.04−1.04 [55] × +0.24 +0.25 +0.32 +0.31 +0.21 +0.39 L(FUV)obs + aIR L(TIR) 0.37−0.15 [11] 0.42−0.06 [46] 0.43−0.12 [37] 0.27−0.07 [30] 0.34−0.07 [82] 0.36−0.08 [44] Table 5: Values of the aIR coefficients for the calibration of hybrid tracers: BY STELLAR MASS (see Sections 4.5.2 and 4.5.4 for a detailed −1 explanation). The stellar masses are in units of M?/M . Expressions to compute the SFR (M yr ) appeared on the left column. These recipes are the same as the ones in Table 3. The luminosities are in units of erg s−1. Without type-2 AGN With type-2 AGN Hybrid Tracers log[M?]>10.5 10.0 Article number, page 30 of 36 A&A–SFR_CALIFA, Online Material p 31 ) 1 − 0.18 0.66 0.22 0.17 0.29 0.34 0.94 1.55 0.19 0.33 0.38 0.32 0.18 0.13 0.69 0.38 0.34 0.45 0.56 0.43 0.05 0.41 0.27 0.56 0.61 0.24 0.15 0.09 0.12 0.37 0.35 0.31 0.84 0.14 0.37 0.27 0.07 0.08 0.83 0.43 0.60 m) ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± µ erg s 42 10 luminosity × α )( 1 − 0.66 20.55 1.800.240.33 9.33 0.95 5.58 1.54 5.03 1.52 14.16 63.64 8.56 0.481.221.90 2.59 11.11 8.93 0.740.12 7.22 1.34 0.430.440.42 12.08 0.83 10.01 1.04 4.84 0.09 10.89 0.80 19.23 0.26 0.65 0.89 16.33 1.27 10.20 23.33 21.94 0.48 7.44 0.42 1.95 0.31 9.07 1.04 19.03 1.640.970.24 10.01 0.27 17.49 0.82 4.35 6.53 0.17 5.07 1.67 2.86 41.14 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 43 10 × flux uncertainty only, an additional ective temperature and therefore on ff α N L(TIR) L(22 − ) (AB mag) ( 1 )F 10 % than what is inferred from Spitzer IRS and − 0.180.78 0.63 0.49 20.08 . . . 1.88 0.130.280.20 0.240.96 0.75 5.22 0.77 5.31 0.41 45.41 13.16 0.730.21 0.24 0.31 2.19 11.22 0.36 0.58 . . . 1.59 0.040.14 0.48 0.52 8.03 . . . 3.57 0.120.220.57 0.560.05 0.620.02 10.02 0.26 0.55 10.06 0.13 1.050.67 0.34 12.49 0.25 0.46 15.40 0.52 0.76 0.56 19.27 0.67 9.37 20.59 23.13 0.06 0.58 . . . 1.40 0.17 0.49 1.86 0.040.48 0.83 0.43 8.42 . . . 1.92 0.12 1.1 15.24 0.180.100.15 0.480.12 0.870.03 0.45 11.77 15.63 0.09 0.31 3.30 0.43 5.28 0.49 34.78 . . . 0.55 0.31 0.59 . . . 4.34 0.07 0.53 . . . 3.42 ∼ obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 42 10 × )( 1 − ) L(NUV 0.440.19 5.71 31.57 0.160.62 21.66 0.15 7.47 0.74 32.50 38.95 0.180.24 23.87 34.45 0.37 9.74 0.240.27 6.78 11.74 0.530.43 19.28 0.77 5.92 0.030.03 46.15 0.25 7.90 0.74 3.80 1.00 42.06 0.40 21.66 36.24 40.05 0.43 10.06 0.19 4.61 0.500.24 6.12 38.85 0.21 3.77 0.170.260.72 29.31 15.52 0.14 7.88 0.05 20.21 0.09 2.73 14.90 0.62 7.21 0.10 6.06 obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 42 10 × ] L(FUV )( 1 0.22 3.20 0.71 20.10 0.35 ...... 9.42 0.300.111.09 17.37 2.04 3.74 0.92 15.99 26.70 ...... 16.75 0.26 19.14 1.130.48 25.89 ...... 12.31 0.08 5.71 0.29 4.36 0.110.05 7.27 ...... 1.41 0.38 11.51 2.770.270.35 3.34 0.320.04 27.81 0.63 3.00 0.45 2.78 .0.96 . . 27.53 0.82 13.41 21.64 21.61 ...... 4.96 0.04 ...... 5.56 0.09 5.90 0.14 2.93 0.10 2.85 0.39 26.14 0.36 1.37 0.53 18.84 0.450.230.51 6.96 0.23 5.21 0.020.26 15.19 0.76 1.84 . . . 9.49 ...... 6.13 . 3.20 0.43 4.19 0.02 3.72 corr − α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± yr 1.1. Note that galaxies UGC 00987 and UGC 03973 are classified as type-1 AGN within m filter are systematically brighter by / µ ) (M 1 − ) SFR[H 0.88 1.17 1.87 4.10 0.75 2.07 0.400.66 1.79 6.26 0.61 1.15 3.92 1.84 12.01 5.27 0.93 1.49 8.670.34 3.70 2.93 0.56 0.44 1.60 1.17 1.820.61 0.37 0.43 3.43 1.54 1.543.20 15.17 0.73 0.90 0.25 2.11 0.06 1.87 0.79 0.23 0.50 3.73 1.39 2.69 1.04 5.71 4.84 1.67 0.38 0.39 0.55 0.66 0.78 0.46 0.56 0.94 2.31 0.66 2.06 0.38 3.20 0.650.65 2.63 4.51 1.29 0.29 1.80 0.01 1.35 1.61 0.14 0.55 0.80 4.49 0.56 2.11 0.22 0.32 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs 0 that also appeared in this table). The errors listed here use the H erg s α = 40 ) 10 α × m (aperture photometry) luminosity. As explained in Section 3.3, it has been pointed out by several authors luminosities, col. (8), and SFRs, col. (9) and in the attenuations used in the paper; (8) Observed H µ α ) L(H SFR values; (10) Observed FUV luminosity; (11) Observed NUV luminosity; (12) UV color in AB system; (13) TIR 0.07 11.47 0.05 40.99 0.04 25.42 0.030.110.04 19.78 0.03 9.57 0.07 28.22 76.62 42.82 0.03 17.8 0.020.01 33.83 34.07 0.07 7.18 0.02 8.92 0.090.55 6.72 7.88 0.06 18.7 0.100.080.05 25.9 0.06 13.05 0.01 35.61 0.07 14.03 0.02 3.58 0.02 40.79 0.06 27.27 45.09 47.0 0.11 6.99 0.06 10.04 0.02 8.28 0.06 7.69 0.07 27.51 0.04 10.74 0.04 32.59 0.040.060.04 21.66 0.05 12.06 0.04 19.75 0.04 21.04 0.08 2.60 5.66 36.34 0.05 5.6 0.14 5.89 α α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± m; (14) 22 µ 1000 7.13 0.43 5.91 0.46 2.36 0.11 9.65 0.94 9.19 0.44 5.97 0.0 4.42 0.0 3.06 0.58 3.37 0.3 7.21 1.35 9.26 0.63 6.31 2.09 6.49 0.0 − 10.36 0.67 22.90 0.65 12.7410.40 0.54 22.32 0.16 1.01 21.73 0.87 14.9413.19 0.75 0.48 20.2233.11 0.0 0.0 17.7414.08 2.57 14.99 0.25 34.39 0.08 10.96 0.96 58.51 0.17 12.55 0.55 16.74 0.64 48.01 0.9 0.68 10.14 0.46 20.5011.35 0.86 14.69 0.72 12.49 0.55 14.38 0.16 39.20 0.0 51.60 1.03 0.88 ) A(H ± ± ± ± ± ± ± ± ± ± ± ± ± 138.27 1.14 cgs) (mag) ( ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs ± 16 α − (H f 10 × flux using an elliptical aperture with a semi-major axis of 36 arcsec; (7) Attenuation obtained using the Balmer decrement. As explained in Section a / α ratio of 2.86 (although lower values are also physically possible in HII regions, depending on the electron density, e β H / α m luminosity values given in column (14) have to be multiplied by 1 µ (Mpc) ( flux have been taken into account for the errors in the H ects; (9) Balmer-corrected and aperture-corrected H β ff (1) ID CALIFA identifier; (2) Galaxy name; (3) Morphological type; (4) Distance in Mpc calculated from redshift; (5) Minor-to-major axis ratio as given by the RC3 catalog, i.e. measured m data. For that reason, the 22 µ ID name type dist b the chemical abundance. This leads to a number of galaxies for which we assumed A(H 24 our sample. As a consequence, they have been excluded from the SFR analysis. Table 1: in the D25 B-band isophote;3.1.4, we (6) have Observed-H assumed an intrinsic H (Wright et al. 2010; Jarrett et al. 2013; Brown et al. 2014) that SF galaxies measured with the WISE 22 7% uncertainty in the H corrected for aperture e luminosity defined as the total infrared emission in the range 8 (1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12) (13) (14) 1 IC 5376 Sb 71.97 0.17 719.75 23 UGC 00005 NGC 7819 Sbc Sc 104.31 0.51 71.47 2676.95 0.76 2617.25 578 IC 15289 UGC 0003610 NGC 000111 NGC 0023 NGC 0036 UGC 00139 Sab Sbc Sbc 89.82 Sb 54.77 Sb 0.46 Scd 64.92 0.48 0.76 65.26 87.09 56.45 3146.26 839.79 0.65 0.62 0.46 4068.14 12029.28 1933.46 2209.21 1214 UGC 0014815 UGC 00312 ESO539-G014 Sc Scd Sd 101.16 59.88 62.13 0.26 0.1 0.46 5192.47 5977.27 255.45 16 MCG-02-02-040 Scd 50.38 0.28 2379.79 1920 ESO540-G00321 NGC 0160 NGC 0165 Sb 47.63 Sa Sb 0.4 75.09 84.19 1559.81 0.57 0.83 443.08 1770.5 2223 NGC 016925 NGC 017126 NGC27 0180 NGC 019228 NGC 021630 Sab NGC31 0214 Sb NGC 023733 Sb NGC 023438 66.2 Sab NGC 0257 56.14 0.25 Sd NGC 0447 75.65 0.87 Sbc 60.17 0.81 Sc 1159.63 0.47 22.3 Sc 2358.23 64.66 Sc 2267.29 0.34 0.74 2532.58 Sa 59.5 63.73 0.57 5110.81 75.17 6034.21 1.0 79.94 0.71 4649.69 0.98 8268.97 5054.03 634.8 39 NGC 0444 Scd 69.18 0.31 973.51 40 UGC 00809 Scd 60.04 0.15 980.49 41 UGC 00841 Sbc 79.83 0.2 725.9 4243 NGC 0477 IC 1683 Sbc Sb 83.83 0.54 69.36 0.46 1597.8 1829.71 45 NGC 0496 Scd 85.95 0.54 3039.25 4849 NGC52 0523 UGC 0098753 NGC 055156 UGC 0105761 UGC 0111065 Sd NGC66 0681 Sa NGC 0716 Sbc UGC 01368 69.53 Sc 66.53 Sc 0.29 74.37 0.3 Sa 90.68 0.44 Sb 1628.81 40.23 0.32 Sab 1340.65 0.36 25.16 2131.63 114.74 65.66 1716.37 0.63 0.33 0.42 824.77 4494.5 5263.42 324.27 70 IC 1755 Sb 113.21 0.22 276.5 A&A–SFR_CALIFA, Online Material p 32 ) 1 − 0.45 0.34 0.08 0.04 0.27 0.06 0.22 0.72 0.45 0.08 0.58 0.58 0.21 0.06 0.15 0.30 0.30 0.12 0.47 0.12 0.59 2.25 0.48 0.22 0.69 0.20 0.25 2.25 0.15 0.20 0.23 0.02 0.11 0.54 0.72 0.81 0.38 0.73 0.07 0.32 0.27 0.20 0.03 0.01 0.40 0.33 0.62 0.09 0.18 m) ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± µ erg s 42 10 × )( 1 − 0.671.15 12.55 12.61 1.450.330.19 0.83 0.32 6.24 1.25 3.63 0.46 4.97 38.99 7.47 0.400.440.47 1.83 9.06 17.58 0.260.25 2.23 4.17 0.620.87 15.73 11.80 0.91 25.43 0.201.801.90 1.65 1.89 14.99 116.55 0.281.36 24.97 3.98 29.99 1.45 7.20 0.58 8.28 5.250.45 122.22 5.71 0.190.46 5.35 4.38 0.300.37 2.28 7.36 1.68 27.82 0.33 2.62 0.49 4.33 0.04 0.33 0.79 5.02 0.550.11 5.04 0.02 1.25 0.18 0.20 2.00 0.110.49 1.41 7.15 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 43 10 × N L(TIR) L(22 − ) (AB mag) ( 1 )F − 0.490.21 0.48 0.31 10.74 14.42 0.18 0.450.270.02 0.26 0.57 . . . 4.36 27.24 0.71 1.18 0.62 . . . 1.31 0.120.51 0.62 0.69 13.44 13.93 0.11 0.39 . . . 0.40 0.090.200.480.13 0.370.22 0.690.11 0.11 1.46 0.48 16.23 0.82 51.06 0.74 16.75 5.27 32.44 0.12 0.41 6.66 0.190.05 0.66 1.1 134.64 8.32 0.01 0.36 ...... 0.02 0.23 ...... 0.140.05 0.33 0.7 4.22 5.02 0.010.030.20 1.15 0.62 0.47 3.06 7.52 . . . 0.37 1.06 0.66 25.07 0.17 0.33 ...... 0.510.76 0.48 0.26 ...... 7.88 . . . 1.80 0.98 4.67 0.03 0.26 0.39 0.120.050.00 1.2 0.43 0.46 1.26 0.20 . . . 1.49 0.10 1.11 . . . 2.27 1.130.11 0.47 0.32 2.29 . . . 2.14 0.07 0.35 1.73 10.16 5.59 21.77 obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s ± 42 10 × )( 1 − ) L(NUV 0.711.67 40.06 34.38 0.24 2.30 1.260.11 21.80 3.91 0.10 114.38 0.98 15.72 0.100.27 19.85 27.45 0.22 8.62 0.630.810.64 13.81 0.13 6.63 0.31 77.17 0.59 21.24 8.97 17.96 0.24 18.92 0.230.03 15.55 7.83 0.03 1.55 0.02 3.18 0.310.29 22.62 8.69 0.010.030.20 0.36 5.67 33.13 0.95 17.26 0.25 4.54 0.850.98 20.61 3.28 3.14 19.54 0.08 5.64 0.480.050.03 9.74 7.55 0.65 0.27 2.79 1.210.24 18.40 17.85 0.15 11.51 obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 42 10 × ] L(FUV )( 1 0.74 25.74 0.55 25.84 0.01 1.52 0.290.350.140.450.51 17.16 2.31 ...... 4.85 . . 3.34 . 7.95 0.04 0.66 0.08 ...... 3.64 0.340.880.04 8.88 ...... 8.66 15.33 0.07 ...... 0.14 ...... 2.19 0.27 ...... 3.77 0.19 11.21 0.42 14.54 0.11 6.02 0.67 ...... 15.48 0.12 9.82 0.444.910.23 3.51 0.14 69.74 2.67 13.65 4.22 9.09 0.35 ...... 8.27 0.37 12.97 0.05 8.47 0.12 2.84 0.01 1.12 0.02 2.57 0.270.29 16.69 4.56 0.01 0.13 0.110.33 3.20 21.49 0.63 9.40 0.06 3.35 0.18 13.25 0.170.01 2.58 ...... 2.98 0.10 7.92 0.03 4.44 0.09 ...... 5.89 0.23 ...... 5.58 0.230.070.02 3.23 5.08 0.42 0.05 1.00 0.10 11.94 0.24 13.29 0.09 8.34 0.53 ...... 6.76 corr − α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± yr ) (M 1 − ) SFR[H 1.94 3.68 0.81 3.31 0.28 0.11 3.321.62 1.17 0.63 1.25 0.17 0.89 7.71 2.70 1.89 1.09 0.24 0.53 0.31 3.074.78 1.12 1.12 3.32 0.65 1.02 0.20 1.26 0.55 2.89 0.96 0.16 1.13 1.38 2.37 1.18 0.43 7.09 1.84 0.22 0.70 1.39 2.37 0.360.42 1.33 0.75 1.94 1.11 0.28 2.24 1.62 0.37 0.41 0.57 0.10 0.03 0.27 0.09 0.280.52 1.62 1.69 0.02 0.03 0.130.48 0.68 1.90 2.40 3.27 0.70 0.22 0.97 1.09 1.940.24 0.69 0.06 2.06 0.79 0.06 0.22 1.81 0.33 2.13 0.87 0.750.06 1.38 0.20 0.41 0.09 0.50 0.25 1.43 0.78 0.28 1.44 0.10 0.52 3.68 1.96 13.73 18.20 15.33 7.89 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs erg s ± ± α 40 10 × 0.1 0.47 ) L(H 0.09 47.04 0.04 32.37 0.10 2.01 0.050.020.07 17.9 0.05 7.34 0.07 16.12 10.52 34.38 0.09 4.42 0.10 2.66 0.140.110.12 12.2 26.17 11.89 0.02 3.43 0.04 6.59 0.06 12.90 0.03 9.01 0.06 36.28 0.03 6.67 0.05 22.12 0.04 10.67 0.090.170.09 18.62 75.27 0.100.05 23.21 8.32 52.74 0.06 7.25 0.02 23.91 0.20 6.65 0.13 9.32 0.08 0.4 0.03 1.43 0.020.06 19.26 17.2 0.020.06 6.41 23.85 0.11 24.48 0.03 3.56 0.12 19.9 0.100.06 10.51 1.01 0.19 14.36 0.02 4.06 0.10 5.95 0.04 10.33 0.100.010.19 25.13 5.36 1.00 0.26 4.62 0.09 14.18 0.04 16.76 0.03 7.23 0.04 17.4 α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 6.08 0.0 8.20 0.0 7.10 0.81 2.46 0.06 7.53 0.45 8.81 0.9 3.57 0.18 7.28 0.19 7.90 1.11 7.99 0.58 6.15 0.00 2.69 0.25 3.91 0.11 3.56 0.16 8.14 0.96 5.13 0.13 7.68 0.0 8.27 0.0 4.43 0.0 5.61 0.46 8.19 0.0 5.75 0.0 9.14 0.29 12.53 0.67 11.3723.11 0.19 11.59 1.23 13.47 0.0 23.12 1.67 0.0 12.2237.35 0.55 11.18 0.91 0.00 13.16 0.32 87.09 0.18 27.36 0.45 13.14 0.91 14.3832.52 0.05 75.78 0.53 1.09 12.1112.33 0.46 0.63 35.6654.72 0.71 0.4 11.3634.93 0.19 0.0 17.58 0.0 11.8013.79 0.0 15.85 0.36 0.57 11.34 0.48 63.11 0.78 ) A(H ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 104.59 0.38 196.92 1.61 124.48 0.11 cgs) (mag) ( ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs ± ± ± 16 α − (H f 10 × a / (Mpc) ( ID name type dist b (1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12) (13) (14) 71 NGC 0768 Sc 100.01 0.48 3116.02 7377 NGC 0776 NGC 0825 Sb Sa 70.29 48.52 0.98 0.39 4594.42 475.93 7888 NGC 0833100 UGC108 01938 NGC 1056115 NGC 1093116 UGC 02403131 UGC Sa 02405 Sbc NGC 1542 Sa Sbc 54.94 91.36 Sb 0.48 0.23 Sbc 22.14 75.85 0.48 Sab 58.73 110.12 0.65 1433.02 829.28 0.38 0.4 53.38 9129.11 1646.19 0.37 1712.66 2022.92 603.47 132133 UGC 03038134 UGC 03107141 NGC 1645 IC 2095 Sab Sb 115.71 S0a 0.23 121.24 Sc 0.28 69.65 0.45 647.4 1205.67 40.89 0.22 645.37 986.85 143 NGC 1677 Scd 39.57 0.34 2771.85 146147 UGC 03253 NGC 2253 Sb Sbc 59.0 51.24 0.63 0.74 2236.96 2418.9 149150 NGC 2347 UGC 03899 Sbc Sd 63.67 55.53 0.71 0.27 5451.42 1473.73 151152 NGC 2410 UGC 03944 Sb Sbc 66.82 55.82 0.29 0.45 2201.04 1854.12 153154 UGC 03969155 UGC 03973156 UGC 03995165 NGC 2449186 UGC 04132 Sb IC 2247 Sbc Sb 117.84 94.29 Sab 0.2 Sbc 68.28 1.0 70.01 0.45 74.46 Sab 908.47 5742.33 0.48 0.3 1535.73 61.8 1278.65 4701.41 0.16 1008.76 208 UGC 04461 Sbc 71.31 0.35 2548.35 213 NGC 2623 Scd 79.17 0.42 606.17 219226 NGC 2639 UGC 04659 Sa Sdm 25.97 46.52 0.47 0.6 2892.4 341.16 231 UGC 04722 Sdm 24.66 0.33 1556.99 232273 NGC 2730274 IC 2487 IC 0540 Scd Sc 54.95 Sab 0.74 62.25 29.6 4530.91 0.33 0.27 1731.12 393.61 275277 NGC 2906278 NGC 2916 UGC 05108 Sbc Sbc Sb 30.59 53.31 0.59 116.19 0.68 0.65 5416.94 3630.93 1008.07 306 UGC 05358 Sd 41.59 0.62 1331.74 307 UGC 05359 Sb 121.27 0.35 686.14 309310 UGC 05396 UGC 05391 Sbc Sm 77.51 22.41 0.45 0.90 1171.19 1159.09 311312 NGC 3106 NGC 3057 Sab Sdm 88.75 21.79 1.0 0.59 1005.35 3416.93 314 UGC 05498NED01 Sa 90.12 0.2 345.63 326 UGC 05598 Sb 80.39 0.25 1013.28 340353 NGC 3303361 NGC 3381364 UGC 05990 UGC 06036 S0a Sd Sc 85.71 Sa 0.71 23.21 23.13 0.91 93.57 0.25 1056.24 0.21 6641.87 1187.42 315.62 381 IC 0674 Sab 107.44 0.44 561.39 388414 NGC 3614 NGC 3687 Sbc Sb 33.15 0.61 35.75 1.0 2513.2 3223.0 436 NGC 3811 Sbc 43.92 0.76 5634.04 A&A–SFR_CALIFA, Online Material p 33 ) 1 − 0.14 0.22 0.21 0.21 0.05 0.27 0.08 0.31 0.15 0.44 0.08 0.08 0.17 0.17 0.57 0.06 0.05 0.27 0.14 0.35 0.10 0.37 0.29 0.08 0.06 0.37 0.33 0.02 0.05 0.05 0.16 0.41 0.07 0.88 0.19 0.57 1.09 0.09 0.36 0.06 0.07 0.25 0.05 0.02 0.24 1.03 0.87 0.15 0.77 0.11 0.08 0.14 0.11 m) ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± µ erg s 42 10 × )( 1 − 0.140.460.46 3.43 8.96 11.72 0.520.87 4.91 0.320.50 13.34 0.21 3.09 4.12 2.05 1.24 15.74 0.250.25 2.81 1.810.17 4.38 18.79 2.02 0.120.60 0.64 11.33 1.540.23 12.03 1.01 3.34 12.22 0.18 2.26 0.080.560.46 0.65 21.27 4.85 0.08 1.06 0.070.150.63 0.44 2.96 12.81 0.620.700.14 36.41 0.35 9.72 0.18 4.80 12.10 2.59 0.371.22 2.75 11.38 0.120.64 0.62 3.611.32 7.78 0.20 37.42 4.53 52.50 0.24 2.82 0.12 36.30 0.54 5.62 2.80 3.13 0.09 1.35 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 43 10 × N L(TIR) L(22 − ) (AB mag) ( 1 )F − 0.210.240.17 0.5 0.24 0.51 3.73 7.83 8.82 0.94 0.51 ...... 0.060.030.07 0.38 1.26 0.41 0.79 14.94 4.43 0.01 0.43 . . . 0.10 0.05 0.35 0.93 0.160.100.16 0.24 0.33 1.07 0.44 2.80 13.95 0.011.96 0.72 0.66 ...... 0.67 4.72 0.030.180.21 0.410.42 0.830.04 0.39 21.54 0.41 9.99 0.52 . . . 11.44 2.77 1.11 0.030.07 0.38 0.85 4.04 13.00 0.070.211.78 0.40.240.16 0.880.08 0.55 0.73 0.420.10 9.63 0.46 40.92 0.07 0.45 31.86 2.65 0.42 37.44 0.5 3.01 4.67 0.30 0.24 11.70 obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 42 10 × )( 1 − ) L(NUV 0.330.290.15 11.39 39.51 13.45 0.36 48.57 2.54 17.00 0.090.190.76 4.50 4.29 11.97 0.03 1.13 0.05 3.83 0.060.110.18 8.53 16.38 6.42 0.053.07 1.60 26.57 0.220.550.22 5.00 7.48 0.64 33.53 0.03 33.91 6.08 0.080.41 2.45 10.77 0.180.18 5.66 2.051.12 11.29 0.15 26.33 0.16 19.86 12.70 0.10 12.75 0.21 16.06 11.78 obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 42 10 × ] L(FUV )( 1 0.18 7.19 0.411.07 31.67 8.41 0.03 ...... 4.80 1.000.040.290.22 ...... 15.45 . . . 4.50 4.00 2.73 0.170.13 ...... 11.06 0.87 0.05 ...... 0.55 0.390.271.270.15 ...... 3.10 . . . . 3.72 22.91 2.10 0.120.08 ...... 0.67 2.95 0.41 ...... 10.04 0.13 ...... 2.47 0.17 ...... 13.75 0.350.060.48 38.94 ...... 3.86 . . . 3.42 0.02 ...... 1.89 0.05 10.63 0.03 3.17 0.310.08 1.35 8.21 0.01 0.76 0.06 2.78 0.05 6.83 0.240.32 12.09 2.40 0.02 0.82 0.25 14.47 0.03 3.43 0.470.360.36 3.48 0.88 23.41 0.26 23.24 3.77 ...... 4.15 0.15 1.72 0.31 4.92 0.04 3.92 0.111.371.46 5.02 0.16 15.87 1.44 13.49 0.48 8.31 0.25 8.43 0.22 10.91 7.43 ...... 5.33 0.01 . . .0.14 ...... 0.09 1.27 corr − α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± yr ) (M 1 − ) SFR[H 0.42 1.08 0.465.58 2.48 4.38 1.19 0.36 7.360.51 3.48 3.53 0.14 2.15 0.85 0.74 1.84 0.65 0.78 0.51 0.84 0.23 3.68 1.39 2.48 1.07 1.82 0.58 2.38 0.52 1.26 0.35 3.07 1.57 1.29 0.51 1.97 0.65 3.041.18 1.21 0.67 0.35 2.89 0.53 0.10 1.38 0.82 0.07 0.17 1.640.43 0.97 0.63 0.03 0.05 0.11 0.38 0.18 0.31 0.301.57 1.42 0.69 0.06 0.11 2.00 1.18 0.18 0.22 0.520.51 2.76 2.62 2.14 7.29 1.30 2.07 3.56 0.92 0.26 0.85 0.44 1.82 0.08 0.24 0.316.29 0.64 5.31 0.310.45 0.95 3.18 8.25 0.28 1.78 0.58 1.52 1.25 0.05 0.03 1.92 0.43 16.72 3.71 10.95 4.25 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs erg s ± ± α 40 10 × ) L(H 0.03 12.12 0.010.02 36.46 31.51 0.12 6.5 0.030.060.08 31.30 0.03 2.44 11.73 8.99 0.02 9.02 0.14 4.24 0.02 4.12 0.010.04 16.05 0.070.02 12.78 56.23 9.28 0.07 9.43 0.02 6.28 0.09 15.2 0.05 6.61 0.09 10.15 0.050.150.03 12.59 6.44 43.88 0.11 1.84 0.24 14.9 0.08 3.04 0.070.07 6.03 11.47 0.04 0.9 0.04 4.15 0.03 5.25 0.020.37 16.42 11.3 0.07 1.75 0.14 16.82 0.16 3.94 0.040.050.01 14.05 0.05 28.21 0.02 11.87 40.11 9.01 0.05 5.76 0.04 16.79 0.03 3.37 0.090.060.020.03 8.6 32.62 36.38 0.060.02 11.03 0.01 21.9 0.09 14.89 18.14 13.25 0.17 0.57 0.04 6.86 α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 7.42 0.0 6.44 0.3 2.79 0.85 5.77 0.0 5.80 0.0 8.91 0.0 4.07 0.36 6.58 0.17 2.54 0.0 8.05 0.0 9.62 0.0 2.96 0.0 5.69 0.57 4.95 0.07 5.80 0.14 9.59 0.0 4.50 1.08 3.96 0.26 8.32 0.0 7.59 0.15 22.02 0.52 21.5497.94 0.23 1.01 28.1313.11 0.76 11.78 0.05 14.31 0.3 0.44 14.94 0.5 13.1527.48 0.45 13.10 0.2 0.13 13.82 0.69 36.6018.63 0.6 29.45 0.0 0.19 14.4120.30 1.17 0.0 14.8378.03 0.49 0.12 14.79 0.26 15.3919.18 1.38 18.44 0.35 24.56 0.74 30.18 0.52 0.67 14.97 0.73 11.7218.79 0.33 20.42 1.18 13.80 0.82 16.32 0.49 30.71 2.09 10.41 0.84 16.59 0.45 0.58 ) A(H ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± cgs) (mag) ( ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs 16 α − (H f 10 × a / (Mpc) ( ID name type dist b (1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12) (13) (14) 437 NGC 3815 Sbc 53.08 0.54 3007.07 475476 NGC 3991479 NGC 3994 NGC 4003 Sm Sbc S0a 46.46 44.16 0.28 93.86 0.56 12687.14 0.59 11022.62 496.7 486489 UGC 07012502 NGC 4047515 NGC 4149518 NGC 4185520 NGC 4210 Scd NGC 4211NED02 Sbc 43.78 Sa S0a Sbc 0.54 49.2 Sb 95.62 44.02 0.85 55.52 0.29 3001.7 0.2 0.74 38.85 7567.21 0.78 1811.94 311.56 792.2 3446.38 528 IC 0776 Sdm 34.89 0.66 2147.53 548569 NGC 4470580 NGC 4644593 NGC 4711603 UGC 08107606 NGC 4961 Sc UGC 08231 Sb Sbc Sa 33.71 70.52 0.72 Scd 58.25 Sd 118.66 0.31 0.57 0.43 8400.82 36.64 35.38 0.69 1054.53 2460.2 1929.57 0.5 4494.24 3228.46 608609 NGC 5000 UGC 08250 Sbc Sc 79.99 75.41 0.85 0.22 1520.88 757.06 610 UGC 08267 Sb 103.64 0.23 618.61 611613 NGC 5016614 UGC 08322615 NGC 5056 IC 4215 Sbc S0a Sc 37.37 108.76 0.74 Sb 0.2 79.41 5183.36 0.57 371.21 55.7 4232.77 0.15 299.29 624 NGC 5157 Sab 104.68 0.72 576.79 630 NGC 5205 Sbc 25.31 0.57 1614.02 634651 NGC 5218652 NGC 5267 UGC 08662 Sab Sab Scd 41.41 84.69 0.69 28.9 0.37 1846.77 0.3 1051.8 700.48 653 NGC 5289 Sab 36.18 0.32 1336.39 657 UGC 08733 Sdm 33.42 0.65 1737.45 659663 NGC 5320664 IC 0944 UGC 08778 Sbc Sb Sab 37.49 0.51 100.11 46.38 0.34 0.21 3216.34 625.43 596.48 665676 UGC 08781 NGC 5378 Sb Sb 108.38 0.5 42.89 0.83 655.29 704.07 680684 NGC 5394707 NGC 5406714 NGC 5480715 UGC 09067718 NGC 5520 Sbc UGC 09113 Sb Scd 48.73 Sbc 0.56 77.16 Sbc 27.05 112.16 Sb 0.74 0.66 0.49 3805.28 26.83 45.59 2406.29 0.55 9886.46 2150.23 0.22 7436.51 985.09 720725 NGC 5522 IC 0995 Sb Scd 65.58 0.21 45.2 0.24 1594.71 1122.24 737740 NGC 5602741 NGC 5614742 NGC 5616743 NGC 5610747 NGC 5622748 S0a UGC 09262749 Sa NGC 5633750 Sc 31.78 NGC 5630751 Sb NGC 0.53 5635 55.59 Sbc 122.17 UGC 09291 Sab 0.83 72.39 0.2 Sbc 55.54 122.33 0.33 Sdm 323.8 1490.19 0.57 0.14 Sa 1434.33 33.24 Scd 37.98 3307.52 0.62 2352.76 0.36 992.55 41.31 61.7 8369.48 0.59 0.49 6765.18 1392.19 2057.09 A&A–SFR_CALIFA, Online Material p 34 ) 1 − 0.19 0.17 0.16 0.20 0.10 0.03 0.06 0.92 0.02 0.05 0.04 0.08 0.12 0.10 0.55 0.92 0.18 0.23 0.17 0.21 0.08 0.29 0.11 0.07 1.57 0.15 0.33 0.76 0.10 0.43 0.27 0.12 0.15 0.04 0.07 0.19 3.94 0.06 0.03 0.06 0.38 0.08 0.35 0.04 0.19 0.05 0.15 0.15 0.37 0.25 0.11 0.23 0.08 0.04 m) ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± µ erg s 42 10 × )( 1 − 0.34 9.01 0.66 8.71 0.640.180.43 3.65 12.22 3.78 0.10 0.64 0.13 1.90 0.04 0.32 0.23 1.71 0.25 0.95 0.17 2.27 0.34 3.27 0.330.621.13 1.94 3.86 18.17 0.370.200.85 4.70 5.12 5.99 0.30 2.68 0.201.280.07 3.45 7.22 0.96 0.131.59 2.54 9.10 0.980.472.70 5.31 17.77 35.95 0.661.07 33.34 10.21 0.05 0.83 0.20 16.87 25.20.13 534.17 1.05 0.12 1.32 0.330.90 4.74 20.22 0.16 1.57 0.40 6.96 0.220.540.22 1.41 0.90 7.63 0.55 4.72 0.12 15.40 0.61 6.39 1.27 5.23 0.13 1.71 0.09 0.47 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 43 10 × N L(TIR) L(22 − ) (AB mag) ( 1 )F − 0.20 0.55 9.75 0.06 0.47 7.04 0.240.080.08 0.8 0.65 0.94 5.66 8.03 4.99 0.02 0.55 0.68 0.04 0.54 2.19 0.18 0.47 13.82 0.070.120.44 0.42 0.43 0.85 4.20 4.70 7.61 0.19 0.42 2.41 0.022.340.06 0.53 0.31 0.66 3.74 11.57 0.68 0.060.14 0.85 0.61 3.22 7.98 0.530.130.30 0.49 0.92 0.51 8.28 17.86 36.23 0.060.090.03 0.34 1.01 0.93 15.53 12.95 . . . 0.51 0.170.03 0.31 0.48 0.90 . . . 3.84 0.020.07 0.41 1.15 10.03 . . . 0.41 0.260.68 1.34 0.34 407.08 1.23 0.05 0.41 . . . 0.34 0.09 0.53 1.19 0.100.150.12 0.62 1.2 0.59 5.03 18.35 . . . 1.19 0.05 0.34 1.94 0.24 0.46 7.81 0.050.060.090.17 0.490.15 0.460.05 0.870.44 1.79 0.49 8.48 1.17 3.10 16.11 0.4 0.83 8.30 1.08 7.09 0.09 0.29 1.52 0.03 0.63 0.52 obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 42 10 × )( 1 − ) L(NUV 0.09 16.27 0.05 9.20 0.140.070.03 6.48 12.89 6.80 0.02 3.64 0.12 3.50 1.05 29.16 0.200.590.39 10.85 19.14 10.17 0.32 10.37 0.074.930.07 2.97 47.50 3.41 0.071.07 1.94 22.58 0.830.181.27 14.22 5.09 24.47 0.060.020.30 4.55 4.64 4.23 0.120.03 5.68 5.02 0.070.10 2.89 5.47 0.310.54 8.39 10.06 0.18 4.00 0.11 2.83 0.560.250.40 8.21 4.04 18.69 0.08 1.46 0.47 19.42 0.050.161.06 8.17 1.29 2.27 0.67 7.53 0.15 27.45 0.26 7.95 7.96 11.91 0.75 7.59 0.01 2.36 obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 42 10 × ] L(FUV )( 1 0.31 9.81 0.57 5.97 0.16 3.10 1.360.12 7.08 2.86 0.06 2.19 0.20 2.13 0.20 ...... 5.15 0.02 ...... 0.41 0.11 ...... 1.99 0.13 ...... 0.91 0.23 ...... 2.53 0.35 ...... 3.73 0.05 ...... 2.59 0.230.30 18.92 ...... 4.97 0.24 7.37 0.200.07 12.88 4.65 0.07 7.04 0.17 1.82 0.330.02 35.70 1.86 0.05 0.89 0.25 12.88 0.24 9.05 0.542.96 2.18 15.30 0.03 3.33 0.371.87 1.83 1.80 0.12 4.27 0.12 3.23 0.160.03 1.98 ...... 2.76 0.24 1.90 0.20 2.44 0.09 7.36 0.01 2.74 0.07 1.73 0.08 4.64 0.020.75 1.34 10.85 0.06 1.07 0.18 12.71 0.15 5.20 0.290.141.14 1.49 0.07 3.38 0.08 17.48 0.24 2.71 5.51 5.55 0.23 5.81 0.06 1.32 corr − α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± yr ) (M 1 − 2.4 2.41 ) SFR[H 0.44 1.83 4.17 1.77 0.79 0.86 0.36 0.68 0.04 0.34 0.29 1.15 3.20 0.96 0.16 0.08 0.74 0.45 1.88 0.45 1.92 0.94 2.73 1.28 2.08 0.47 2.70 0.83 0.40 1.44 0.470.81 1.14 0.58 0.35 0.51 0.22 1.01 5.140.18 1.47 0.15 0.22 0.27 4.02 1.24 1.04 1.34 0.87 3.06 0.06 0.17 0.374.05 2.21 3.19 0.46 0.67 0.12 0.73 1.61 0.59 0.32 0.13 0.18 1.45 0.46 1.13 0.73 0.46 0.08 0.12 0.23 0.42 0.35 0.47 0.340.92 0.26 4.31 0.15 0.34 0.37 1.06 0.17 0.87 0.410.39 1.64 0.75 0.79 0.57 6.73 0.14 0.57 0.84 0.49 1.33 0.29 1.35 0.12 0.34 ± 10.43 3.89 13.32 9.70 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs erg s ± ± α 40 10 × ) L(H 0.03 20.49 0.03 14.98 0.05 9.72 0.010.11 33.94 10.25 0.02 4.22 0.03 7.35 0.11 17.37 0.04 0.74 0.02 3.82 0.01 7.72 0.02 10.26 0.03 12.84 0.13 8.58 0.040.16 12.02 43.89 0.04 13.05 0.070.12 16.67 10.62 0.07 9.34 0.02 7.99 0.100.15 26.71 2.69 0.22 3.03 0.10 22.46 0.05 17.15 0.060.04 17.93 52.15 0.17 2.77 0.020.45 11.2 14.17 0.05 8.08 0.07 9.19 0.040.05 7.5 1.76 0.02 9.12 0.18 7.73 0.03 7.18 0.11 2.25 0.05 4.16 0.10 8.63 0.160.04 4.67 36.04 0.06 2.15 0.05 16.14 0.01 10.34 0.050.170.10 8.33 0.22 8.79 0.04 37.21 0.04 10.35 6.83 13.63 0.04 8.77 0.02 4.26 α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 9.47 0.83 5.47 0.51 4.40 0.4 4.72 1.13 1.63 0.76 5.54 0.83 8.35 0.06 6.56 0.55 8.84 0.65 9.81 0.0 6.60 0.75 7.38 0.0 7.19 0.9 8.72 0.51 8.74 0.38 5.58 0.15 3.18 0.45 4.29 0.38 1.98 0.33 9.93 1.06 6.45 0.0 6.61 0.67 7.29 0.0 2.72 1.14 9.97 0.19 8.56 0.47 4.77 1.12 7.48 0.41 75.32 0.53 20.2722.43 0.8 0.20 10.36 0.0 12.1018.49 0.25 0.0 28.9943.32 0.24 0.0 35.5418.06 0.0 0.0 24.54 0.0 29.7012.74 1.23 1.32 13.33 1.39 16.33 0.41 30.83 1.15 15.06 0.16 10.1693.74 0.0 0.84 26.1614.81 1.38 35.04 0.53 12.28 1.29 11.62 0.0 18.29 0.29 0.62 ) A(H ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 192.02 1.53 cgs) (mag) ( ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs ± 16 α − (H f 10 × a / (Mpc) ( ID name type dist b (1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12) (13) (14) 753 NGC 5656 Sb 45.13 0.78 6985.15 754 NGC 5657 Sbc 56.06 0.41 2452.82 755 NGC 5659 Sb 64.36 0.27 1074.51 756757 NGC 5665758 NGC 5675 NGC 5682 Sc Sa Scd 31.75 56.62 0.71 32.47 0.35 0.34 15687.1 1185.3 2441.9 763 NGC 5714 Sb 32.28 0.12 1105.81 764765 NGC 5720 UGC 09448 Sbc Sb 111.3 31.61 0.66 0.11 940.15 460.05 766 NGC 5730 Scd 36.15 0.2 1920.86 767 NGC 5731 Sd 36.07 0.32 3391.82 768 NGC 5732 Sbc 53.88 0.54 2352.9 769 UGC 09476 Sbc 46.77 0.65 3659.45 770 NGC 5739 S0a 79.11 0.91 785.48 771774 NGC 5735775 UGC 09537 UGC 09542 Sbc Sb Sc 53.61 126.22 0.79 0.22 78.41 0.29 2522.88 484.34 1130.64 777778 NGC 5772779 NGC 5784 UGC 09598 Sab S0 Sbc 69.77 0.62 77.68 79.89 0.98 0.38 1801.35 1305.07 964.16 783 UGC 09665 Sb 36.56 0.27 2797.73 784787 NGC 5829788 NGC 5876 UGC 09759 Sc S0a Sab 81.11 46.7 49.29 0.87 0.49 0.35 2667.82 472.87 541.31 789790 NGC 5888 UGC 09777 Sb Sbc 124.62 67.16 0.6 0.59 985.88 1769.6 791792 NGC 5908793 UGC 09842 UGC 09849 Sa Sbc Sd 127.36 47.86 0.36 0.38 49.64 0.22 1775.46 2489.7 585.91 795796 NGC 5930797 NGC 5934 UGC 09873 Sab Sb Sb 37.5 0.53 79.8 80.52 0.47 0.25 4311.22 1274.0 805.85 798799 UGC 09892800 NGC 5951 UGC 09901 Sbc Scd Sc 81.13 0.23 25.42 45.44 0.43 0.15 706.62 3010.59 573.54 801802 NGC 5953 ARP 220 Sa Sd 28.29 0.83 77.74 0.81 6257.23 905.5 803804 NGC 5957 NGC 5971 Sb Sb 26.19 0.93 48.3 0.35 3272.14 473.6 805 UGC 09919 Sc 45.53 0.28 1237.27 807 IC 4566 Sb 79.79 0.65 974.58 809810 NGC 5987811 NGC 5980 UGC 10043 Sa Sbc Sab 42.86 58.41 0.32 31.0 0.35 0.33 464.21 6441.01 490.3 813 NGC 6004 Sbc 54.8 0.93 3553.6 817 IC 1151 Scd 31.1 0.41 4574.39 818820 UGC 10123821 NGC 6032822 NGC 6060823 UGC 10205824 NGC 6063 Sab825 IC 1199 Sbc UGC 10257 53.85 Sb S0a 0.23 62.1 Sbc 64.56 93.98 0.45 1924.86 Sbc 0.53 Sb 0.59 40.86 1030.17 0.59 54.57 3361.37 67.71 705.13 0.22 2146.08 0.34 1880.03 1180.6 827 UGC 10297 Sc 32.89 0.29 1910.39 A&A–SFR_CALIFA, Online Material p 35 ) 1 − 0.3 0.24 0.50 0.22 0.27 0.38 0.08 0.40 0.09 0.17 0.06 0.39 0.08 1.69 0.03 0.39 0.65 0.14 0.60 0.41 0.63 0.09 0.24 0.37 0.95 0.35 0.31 1.17 0.87 0.42 0.38 0.14 0.13 0.05 0.43 0.52 0.36 0.38 0.17 0.56 0.04 0.47 0.53 0.35 0.92 0.07 0.16 0.20 0.27 0.12 0.20 0.29 0.40 ± m) ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± µ erg s 42 10 × )( 1 − 0.530.600.80 11.82 5.79 5.33 1.080.221.32 6.58 0.19 4.13 0.59 16.56 3.25 8.33 1.060.111.24 8.79 1.41 17.46 0.080.30 0.52 0.36 21.04 2.520.29 3.83 0.95 25.01 0.09 2.47 26.19 1.30 0.89 8.84 0.820.46 9.28 5.82 0.750.693.85 50.78 7.45 21.55 0.27 3.01 0.600.880.95 8.83 14.06 16.20 0.19 2.48 0.46 7.74 2.283.02 18.74 0.631.40 16.07 2.71 10.08 13.24 58.82 2.47 3.90 0.29 3.69 0.270.560.18 2.54 0.41 7.79 0.78 8.00 1.07 3.18 8.98 19.90 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 43 10 × N L(TIR) L(22 − ) (AB mag) ( 1 )F − 0.080.650.10 0.35 0.38 0.61 7.30 6.26 7.14 0.15 0.54 19.88 0.050.090.050.45 0.39 0.71 0.53 0.39 10.05 1.18 15.27 . . . 1.11 0.040.120.100.47 0.320.07 0.51 0.46 0.48 0.55 22.83 0.09 0.69 3.45 19.63 0.4 . . . 1.02 5.22 0.290.38 0.47 0.37 8.44 . . . 1.55 0.25 0.39 7.95 0.67 0.35 . . . 1.18 3.340.11 1.44 0.65 35.91 7.13 0.05 0.41 ...... 1.02 1.52 2.46 0.120.100.350.12 0.410.92 0.68 0.55 0.68 9.58 0.62 15.69 16.69 . . . 0.12 . . . 2.54 0.10 0.37 2.44 0.31 0.76 7.99 0.240.12 0.630.131.10 22.82 0.38 0.78 12.29 0.4 13.74 45.53 0.84 1.380.04 17.36 0.55 4.16 0.040.180.050.05 0.660.18 0.410.39 0.89 2.07 0.53 6.94 0.52 6.29 0.97 2.75 10.10 14.00 10.17 0.95 18.33 obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s ± 42 10 × )( 1 − ) L(NUV 0.091.250.71 13.48 17.51 16.89 0.37 8.36 0.050.750.12 7.51 2.35 14.31 4.35 72.95 0.050.290.10 6.76 1.58 6.29 0.12 15.78 18.95 11.86 0.22 3.81 0.500.93 9.35 20.34 0.53 41.38 1.78 12.11 0.700.28 28.59 18.50 0.03 4.28 0.10 2.66 0.110.571.87 2.41 0.18 16.64 1.89 56.29 18.77 16.54 0.53 16.13 0.27 8.37 1.31 19.61 0.250.093.14 19.62 21.09 22.43 0.351.20 15.08 70.23 0.11 6.48 0.010.300.06 2.18 0.10 9.54 0.17 3.72 0.56 8.59 28.96 4.51 obs ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 42 10 × ] L(FUV )( 1 0.46 9.77 0.340.24 12.34 9.63 0.10 ...... 2.89 0.13 ...... 5.77 0.450.690.220.44 5.08 ...... 4.11 . 3.86 7.67 0.22 5.24 0.730.060.87 7.44 2.67 50.94 0.04 5.04 1.660.090.25 3.93 0.24 10.33 0.15 11.42 1.11 6.28 0.07 2.63 ...... 2.03 29.47 0.04 6.06 0.35 14.47 0.06 28.89 0.31 ...... 5.02 0.07 8.78 0.32 7.59 0.340.12 10.17 ...... 18.30 0.06 2.93 0.29 ...... 4.04 0.34 0.66 0.01 ...... 0.67 0.05 1.65 0.560.590.51 8.89 0.17 33.92 10.03 9.34 0.13 11.47 0.35 4.16 0.08 ...... 0.17 0.49 10.98 0.801.022.53 13.82 10.28 15.52 0.34 4.23 0.13 3.90 1.80 29.28 0.09 1.19 0.350.480.15 6.54 0.43 1.64 1.38 5.27 17.94 1.84 corr − α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± yr ) (M 1 − ) SFR[H 0.38 2.80 1.930.37 1.26 1.46 2.17 0.64 2.71 0.54 3.270.87 1.55 1.24 4.03 3.46 0.88 1.50 2.66 0.65 1.250.25 3.94 1.50 0.31 4.81 0.04 0.21 2.620.37 6.28 0.78 0.88 0.34 1.47 2.70 1.36 8.72 0.52 0.20 3.34 0.39 0.36 0.30 1.03 1.97 0.07 0.33 4.59 0.98 0.73 0.43 5.54 1.65 0.332.88 2.04 0.66 0.09 0.35 2.45 1.11 1.92 1.17 0.17 0.05 0.24 0.25 1.340.89 2.06 0.49 3.51 1.10 3.00 0.88 0.21 0.80 1.80 1.90 1.23 0.24 1.12 2.85 6.140.58 2.67 4.24 6.15 13.68 3.54 1.37 0.12 0.78 5.92 10.22 0.14 0.51 0.682.15 2.04 0.14 1.86 0.31 0.93 2.38 2.59 6.51 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs erg s α 40 10 × ) L(H 0.01 24.61 0.170.03 21.18 18.0 0.12 11.72 0.14 9.86 0.020.020.02 13.65 0.04 22.98 6.94 14.46 0.01 9.2 0.060.100.11 27.25 5.55 48.57 0.05 3.8 0.050.180.05 13.15 0.05 16.04 0.06 17.98 0.14 18.8 0.07 9.41 37.49 3.87 0.06 5.48 0.05 24.03 0.12 4.38 0.11 17.47 0.10 7.89 0.09 30.06 0.070.10 23.88 11.95 0.01 4.61 0.06 11.97 0.06 7.9 0.14 0.84 0.02 3.06 0.180.050.05 19.4 0.10 50.61 29.08 12.04 0.06 12.01 0.05 22.99 0.02 3.96 0.06 32.26 0.040.030.03 24.59 52.61 52.69 0.05 19.77 0.04 8.46 0.03 102.89 0.06 4.3 0.030.010.05 18.66 0.02 10.7 0.08 10.13 30.52 18.33 α ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 5.27 0.79 5.66 0.0 2.79 0.0 5.30 0.26 7.69 0.03 3.34 0.0 2.58 0.35 8.11 0.0 8.08 0.0 3.55 0.34 8.73 0.57 7.33 1.07 1.48 0.0 4.42 0.41 9.86 0.21 9.41 0.44 3.71 0.11 9.62 0.25 5.96 0.56 5.55 0.84 51.7518.34 0.09 0.42 12.2222.74 0.79 11.17 1.26 13.45 0.91 0.69 21.5926.78 1.05 35.09 0.0 0.64 46.8112.25 2.35 12.20 0.0 14.07 0.43 11.10 0.3 0.0 22.75 0.65 35.86 0.43 23.82 0.02 22.4614.24 0.48 0.0 29.3949.73 0.25 11.54 0.68 0.31 16.84 0.51 15.1819.06 0.74 24.33 0.82 1.69 27.00 0.64 22.3224.44 0.75 10.78 1.25 15.50 0.55 20.76 0.47 2.03 ) A(H ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 123.99 0.52 212.32 0.71 cgs) (mag) ( ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs ± ± 16 α − (H f 10 × a / (Mpc) ( ID name type dist b (1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12) (13) (14) 828 UGC 10331 Sc 65.1 0.29 3409.89 830831 UGC 10337833 NGC 6132 NGC 6154 Sb Sbc Sab 125.4 71.21 0.3 0.32 85.3 0.93 881.95 2417.63 1076.34 834 UGC 10380 Sb 125.22 0.16 326.25 836837 NGC 6155841 UGC 10384842 NGC 6168843 NGC 6186 UGC 10650 Sc Sb Sc 34.59 Sb 70.75 Scd 0.69 36.84 0.2 42.12 42.31 0.35 6598.14 0.76 0.24 2961.0 3470.4 4712.9 2549.51 847848 UGC 10710849 NGC 6310852 NGC 6301 UGC 10796 Sb Sb 119.82 Sbc Scd 0.24 48.74 119.74 43.82 0.25 0.62 1017.53 0.76 1680.26 969.33 1348.1 853854 NGC 6361856 UGC 10811857 IC 1256861 NGC 6394862 UGC 10972863 Sab NGC 6478 Sb866 NGC 6497 54.69 UGC 11262 124.67 Sb Sbc 0.29 Sbc 0.29 121.81 Sc 68.16 2842.15 Sab 0.27 66.5 Sc 591.91 0.71 0.27 97.15 44.99 931.74 0.37 79.81 0.46 2185.2 1063.16 0.38 2381.38 1204.58 499.54 868869 MCG-02-51-004 NGC 6941 Sb 80.7 Sb 0.35 88.94 2679.21 0.71 220.02 871872 NGC 6978 UGC 11649 Sb Sab 85.44 54.24 0.47 0.74 1293.92 1009.84 873 UGC 11680NED01 Sb 111.39 0.65 978.61 876877 NGC 7047878 UGC 11717 MCG-01-54-016 Scd Sbc Sab 42.09 82.85 90.72 0.14 0.55 0.36 2744.54 1451.01 836.95 879 UGC 11740 Sbc 92.19 0.32 895.94 880 UGC 11792 Sbc 68.78 0.21 957.31 884 UGC 11982 Scd 69.6 0.14 111.21 885 UGC 12054 Sc 29.92 0.24 1738.99 886887 NGC 7311889 NGC 7321890 NGC 7364891 UGC 12185 UGC 12224 Sa Sbc Sab Sb 64.46 101.84 Sc 0.48 0.66 69.41 95.05 3342.64 0.6 50.42 2977.12 0.54 0.87 3883.28 705.74 2270.75 892 VV488NED02 Sb 70.49 0.53 1255.21 895 UGC 12308 Scd 32.37 0.27 1738.61 896897 NGC 7466 UGC 12348899901 NGC 7536904 NGC 7549906 Sbc NGC 7591 Sb IC 5309 107.4 107.14 Sc 0.3 0.29 Sbc Sbc 67.15 1590.85 67.42 1168.95 0.39 Sc 70.82 0.46 0.41 3062.51 60.03 3865.33 3622.21 0.45 1694.21 898 NGC 7489907 Sbc NGC 7608 89.12 0.54 Sbc 4720.99 50.31 0.33 1169.96 909913 UGC 12519914 NGC 7625915 NGC 7631918 NGC 7653 UGC 12633 Sc Sa Sb 62.6 Sb Sab 23.23 0.37 53.67 0.91 61.04 59.65 3204.47 0.4 12721.06 0.89 0.68 1942.46 5614.42 2474.82 A&A–SFR_CALIFA, Online Material p 36 ) 1 − 0.27 0.16 0.20 0.14 0.34 1.20 0.75 0.24 1.15 0.30 0.07 0.13 0.07 0.22 m) ± ± ± ± ± ± ± ± ± ± ± ± ± ± µ erg s 42 10 × )( 1 − 0.390.93 3.88 0.10 5.89 1.46 0.772.861.21 6.28 68.51 23.34 0.500.92 3.01 12.53 1.21 9.56 0.34 1.36 0.160.091.18 3.17 1.02 7.84 ± ± ± ± ± ± ± ± ± ± ± ± ± erg s 43 10 × N L(TIR) L(22 − ) (AB mag) ( 1 )F − 0.160.190.100.08 0.38 0.52 0.28 3.23 0.31 5.90 1.69 . . . 1.05 0.510.27 0.13 0.55 6.25 50.58 0.14 0.32 3.48 0.10 0.66 10.49 0.02 0.43 2.16 0.94 0.31 3.57 obs ± ± ± ± ± ± ± ± ± ± erg s 42 10 × )( 1 − ) L(NUV 0.160.131.02 25.37 0.17 7.55 15.88 12.19 0.390.08 3.94 14.60 0.47 22.83 0.04 7.73 0.02 3.09 0.79 19.08 obs ± ± ± ± ± ± ± ± ± ± erg s 42 10 × ] L(FUV )( 1 0.24 17.88 0.040.380.11 4.67 12.27 9.17 0.17 3.50 1.440.93 8.80 ...... 19.29 0.24 17.01 1.47 ...... 12.96 0.18 4.21 0.08 2.08 0.19 14.34 0.250.28 ...... 0.94 9.39 corr − α ± ± ± ± ± ± ± ± ± ± ± ± ± ± yr ) (M 1 − ) SFR[H 0.34 1.42 0.180.74 0.17 0.10 2.27 0.65 0.92 0.64 1.258.04 8.49 3.62 0.41 1.42 9.22 4.01 0.28 1.07 0.05 0.47 1.32 1.06 2.852.59 0.74 1.40 ± ± ± ± ± ± ± ± ± ± ± ± ± ± obs erg s α 40 10 × ) L(H 0.15 19.33 0.510.040.05 2.88 25.39 9.61 0.11 4.68 0.110.05 41.51 41.76 0.04 19.56 0.12 29.81 0.08 12.24 0.04 4.44 0.05 17.1 0.010.05 9.69 22.31 α ± ± ± ± ± ± ± ± ± ± ± ± ± ± 5.71 0.99 9.78 0.3 2.94 0.5 6.56 0.71 4.25 0.13 24.82 0.31 17.3011.65 0.06 13.45 0.53 0.23 12.1526.06 1.43 0.49 79.35 0.97 14.5210.09 0.35 0.14 ) A(H ± ± ± ± ± cgs) (mag) ( ± ± ± ± ± ± ± ± ± obs 16 α − (H f 10 × a / (Mpc) ( ID name type dist b (1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12) (13) (14) 920 NGC 7691 Sbc 57.66 0.76 2563.25 921922 UGC 12653924 UGC 12688926 NGC 7716 UGC 12723 Sc Scd Sb 57.48 Sc 74.68 0.31 0.33 36.63 78.11 0.83 456.04 2457.96 0.33 3426.16 536.26 927929 NGC 7738930 UGC 12810 UGC 12816 Sb Sbc Sc 115.57 97.92 0.76 0.4 76.21 0.55 1757.99 2053.33 1902.54 931933 NGC 7782 NGC 7787 Sb Sab 76.87 95.22 0.54 0.26 2350.61 553.26 934 UGC 12857 Sbc 35.15 0.21 1920.31 935 UGC 12864 Sc 66.83 0.6 1591.1 937938 NGC 7800 NGC 5947 Ir Sbc 24.88 84.84 0.69 0.98 7629.71 2426.3