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A Spectroscopically Confirmed Excess of 24 micron Sources in a Super Galaxy Group at z=0.37: Enhanced Dusty Star Formation Relative to the Cluster and Field Environment

To trace how dust-obscured star formation varies with environment, we compare the fraction of 24 micron sources in a super galaxy group to the field and a rich galaxy cluster at z~0.35. We draw on multi-wavelength observations that combine Hubble, Chandra, and Spitzer imaging with extensive optical spectroscopy (>1800 redshifts) to isolate galaxies in each environment and thus ensure a uniform analysis. We focus on the four galaxy groups in supergroup 1120-12 that will merge to form a galaxy cluster comparable in mass to Coma. We find that 1) the fraction of supergroup galaxies with SFR(IR)>3 Msun/yr is four times higher than in the cluster (32% vs. 7%); 2) the supergroup's infrared luminosity function confirms that it has a higher density of IR members compared to the cluster and includes bright IR sources not found in galaxy clusters at z<0.35; and 3) there is a strong trend of decreasing IR fraction with increasing galaxy density, i.e. an IR-density relation, not observed in the cluster. These dramatic differences are surprising because the early-type fraction in the supergroup is already as high as in clusters, i.e. the timescales for morphological transformation cannot be strongly coupled to when the star formation is completely quenched. The supergroup has a significant fraction (~17%) of luminous, low-mass, IR members that are outside the group cores (R>0.5 Mpc); once their star formation is quenched, most will evolve into faint red galaxies. Our analysis indicates that the supergroup's 24 micron population also differs from that in the field: 1) despite the supergroup having twice the fraction of E/S0s as the field, the fraction of IR galaxies is comparable in both environments, and 2) the supergroup's IR luminosity function has a higher L(IR)* than that previously measured for the field.


INTRODUCTION
Galaxies in the field environment span a wide range in morphology, color, and ongoing star formation (e.g. Marzke et al. 1998). In contrast, the significantly more crowded environment of galaxy clusters is dominated by passive, red, earlytype galaxies that formed the bulk of their stars at z > 2 (Gre-gory & Thompson 1978;Dressler 1980;Bower et al. 1992;van Dokkum et al. 1998a). Using spectroscopically defined samples, several studies show that galaxies in clusters differ from their field counterparts even up to z ∼ 1.4 (e.g. Holden et al. 2007;Mei et al. 2009;Lidman et al. 2008).
However, we know that the galaxy populations in clusters have evolved since at least z ∼ 1. Observational examples of how the galaxy mix in clusters evolves with increasing redshift include: the increasing fraction of blue/starforming members (Butcher & Oemler 1978;Ellingson et al. 2001); the increasing fraction of spectroscopically confirmed 24µm sources (Saintonge et al. 2008, Paper I); the increasing fraction of massive post-starburst members (Tran et al. 2003;Poggianti et al. 2004); the increasing fraction of active galactic nuclei (Eastman et al. 2007;Kocevski et al. 2009); the increasing fraction of star-forming galaxies with increasing galaxy density at z ∼ 1, a reversal of what is observed at z ∼ 0 (Elbaz et al. 2007;Cooper et al. 2008); the decreasing fraction of S0 galaxies (Postman et al. 2005;Moran et al. 2007); and the decreasing fraction of faint red galaxies (z > 0.5; Tanaka et al. 2007;Stott et al. 2007;De Lucia et al. 2007, cf. Crawford et al. 2009).
The problem is that we have yet to identify clearly the physical mechanisms responsible for the dramatically different galaxy populations in clusters versus the field, nor the timescales needed for these mechanisms to operate. Although a plethora of physical processes have been invoked to quench star formation and transform galaxies into spheroidal systems, e.g. starvation (Bekki et al. 2002), ram-pressure stripping (Abadi et al. 1999), and galaxy harassment (Moore et al. 1998), stringent observational tests of when star formation is quenched and whether quenching is coupled to morphological transformation are needed to assess the relative importance of the physics at work. Simulations are sufficiently advanced that new insight can be obtained by, e.g. comparing star formation and gas-loss rates as a function of local density to the observations (Tonnesen et al. 2007).
Also, instead of focusing on massive clusters, the key to understanding the interplay between galaxy evolution and environment is to study galaxy groups because: 1) most galaxies in the local universe are in groups (e.g. Geller & Huchra 1983); and 2) hierarchical structure formation predicts that galaxy clusters assemble from the merger and accretion of smaller structures, e.g. groups (Peebles 1970). Simulations show that the physical mechanisms normally associated with galaxy clusters are also effective in groups (Hester 2006;Kawata & Mulchaey 2008;McCarthy et al. 2008), but that they operate in groups at lower redshifts (Romeo et al. 2008).
In fact, the galaxy groups in the local universe do have more in common with galaxy clusters than with the field population, i.e. higher early-type fractions and lower mean star formation rates than the field (Zabludoff & Mulchaey 1998;Hashimoto et al. 1998;Tran et al. 2001;Blanton & Berlind 2007;Rasmussen et al. 2008). With the advent of large spectroscopic studies such as CNOC (Yee et al. 1996) and SDSS (Abazajian et al. 2003), the galaxy populations in groups can now be studied for statistically large samples (e.g. Yang et al. 2007), and at intermediate redshifts (e.g. Poggianti et al. 2008;Gal et al. 2008;Knobel et al. 2009). Although still nascent, spectroscopic studies of galaxy groups at z > 0.3 find that the groups already have high early-type fractions (Jeltema et al. 2007;Wilman et al. 2008). However, whether star formation in z > 0.3 groups is enhanced or simply quenched relative to the field is debated Balogh et al. 2009).
The question then is whether the evolution of galaxies in clusters is driven primarily on group or on cluster scales. Our discovery of a supergroup of galaxies at z = 0.37 allows us to uniquely answer this question. The supergroup (hereafter SG1120) is composed of multiple galaxy groups that we have shown will merge into a cluster comparable in mass to Coma by z ∼ 0 , unlike the majority of clusters studied at z > 0.3 that are too massive to be Coma progenitors. Because we know the galaxies in SG1120 will evolve into a cluster population, we can test whether the group galaxies are already like those in clusters. First results from our multiwavelength study of SG1120 show that the group galaxies are in transition: SG1120 has a high fraction of early-type members ), yet several of the most massive group galaxies are growing by dissipationless merging at z < 0.4 ).
Here we focus on the dust-obscured star formation in the supergroup as measured with MIPS (Rieke et al. 2004) 24µm observations. Studies find a surprising number of midinfrared sources at cluster and group densities at z > 0.3 (Elbaz et al. 2007;Bai et al. 2007;Koyama et al. 2008;Dressler et al. 2009), but this may be due to the general increase in the fraction of mid-infrared galaxies with redshift (Le Floc'h et al. 2005). To determine if there is an excess of IR sources in the galaxy groups making up SG1120, we compare the 24µm members in SG1120 to their counterparts in both the field and cluster environment at the same redshift.
For the cluster environment, we use CL 1358+62, a massive, dynamically evolved, X-ray luminous galaxy cluster at z = 0.328 with a line-of-sight velocity dispersion of 1027 +51 −45 km s −1 (Fisher et al. 1998, hereafter F98). Most of the 232 spectroscopically confirmed members are passive, earlytype galaxies (van Dokkum et al. 1998b). In addition to the spectroscopy, we have HST/WFPC2 mosaics taken in F606W and F814W covering ∼ 502 ′ , and MIPS 24µm imaging.
The field sample is drawn from extensive observations of two higher redshift clusters: MS 2053-04 (z = 0.59; Tran et al. 2005) and MS 1054-03 (z = 0.83; Tran et al. 2007). In the combined area of approximately 502 ′ , we have measured spectroscopic redshifts for nearly 300 field galaxies at 0.09 < z < 1.36. These fields were also imaged with HST/WFPC2 in F606W and F814W, and with MIPS at 24µm. The depth and uniformity of our spectroscopic and photometric observations in these fields makes for a unique dataset that enables robust comparison across environment at z ∼ 0.35.
Our study of how 24µm galaxies vary across environment is Paper II in our SMIRCS (Spitzer/MIPS Infra-Red Cluster Survey) series and complements Paper I (Saintonge et al. 2008) where we explored how the 24µm fraction increases with increasing redshift in massive galaxy clusters at (0 < z < 0.83; see also Finn et al., in prep). Throughout the paper, we use H 0 = 70 km s −1 Mpc −1 , Ω M = 0.3, and Ω Λ = 0.7; at z = 0.37, this corresponds to a scale of 5.12 kpc per arcsec and a lookback time of 4 Gyr. All restframe magnitudes are in the Vega system.

OBSERVATIONS
At these redshifts (z ∼ 0.35), our large number of spectroscopically confirmed group (174) and cluster members (232) combined with multi-wavelength imaging that includes MIPS observations is unique among existing surveys. Unlike many spectroscopic cluster surveys at intermediate redshifts, we do not select by optical color which can bias a sample towards members that are already on the red sequence, e.g. against blue, star-forming members. The depth of our redshift surveys in the cluster fields also enables us to identify a sample of field galaxies (0.25 ≤ z ≤ 0.45;z = 0.35) that have been observed and analyzed in the same manner as the group and cluster galaxies. The uniformity of our observations allows us to compare directly galaxy populations across a range of environments at z ∼ 0.35.
Line-matched photometric catalogs were generated using the VIMOS R mosaic as the master detection image (SExtractor v2.5.0; Bertin & Arnouts 1996). While several close galaxy pairs (separation< 1 ′′ ) are considered single objects in the R catalog, this is appropriate for our analysis given that the same close pairs are also blended sources in the MIPS catalog (see §2.4). We use k-correct v4.1 (Blanton & Roweis 2007) to determine rest-frame absolute magnitudes (Vega) and Kcorrections. As input, we use the MAG − AUTO photometry from the g ′ BV r ′ R imaging and assumed minimum photometric uncertainties in each bandpass of 0.05 mag. The photometry has been corrected for foreground Galactic extinction using the Schlegel et al. (1998) dust maps and the O'Donnell (1994) Milky Way extinction curve, assuming R V = 3.1.
For consistency and to thus ensure that our comparisons are robust across the supergroup, cluster, and field samples, we calculate stellar masses in the same manner by following the prescription in Bell et al. (2003). Here mass-to-light ratios (M * /L) B are calculated from (B − V ) colors using: and a diet Salpeter initial mass function (IMF) is assumed. We use the diet Salpeter IMF defined in Bell & de Jong (2001) as having x = 0 below 0.6M ⊙ and so the stellar mass using a diet Salpeter IMF is 70% of that for a regular Salpeter IMF (Salpeter 1955). Using an absolute magnitude for the sun of The galaxy populations in CL 1358+62 (z = 0.328; F98) have been studied extensively using optical imaging and spectroscopy. For the cluster galaxies, we use the optical photometry measured by Holden et al. (2007, hereafter H07) from the HST/WFPC2 mosaics (total area of ∼ 502 ′ ). To summarize, a Sérsic profile (1 ≤ n ≤ 4) was fit to the surface brightness distribution in the HST/WFPC2 imaging of each spectroscopically confirmed member; over 85% of the cluster members have n ≥ 2. Galaxy colors were determined from fluxes measured within a half-light radius; the half-light radii were determined using the F814W imaging. Note that at z ∼ 0.33, the redshifted B and V filters are well-matched to F606W and F814W. As in SG1120, we convert the observed fluxes (corrected for Galactic extinction) to rest-frame BV magnitudes using k − correct, and estimate stellar masses using Eq. 1. For more details about the photometry and testing the robustness of the stellar mass determination for the cluster galaxies, we direct the reader to the extensive dicussion in H07.

Field Galaxies (z = 0.35)
Our field sample is drawn from a larger program that focused on galaxies in X-ray luminous clusters at intermediate redshifts. To select field galaxies in the same redshift range as SG1120 and CL1358, we use observations of galaxy clusters MS 2053-04 (z = 0.59; Tran et al. 2005) and MS1054-03 (z = 0.83; Tran et al. 2007). Both galaxy clusters were imaged by HST/WFPC2 in the F606W and F814W filters; each image mosaic is composed of six overlapping pointings and each mosaic covers an area of ∼ 252 ′ . The image reduction and photometry are detailed for MS2053 and MS1054 in Hoekstra et al. (2002) and van Dokkum et al. (2000), respectively. 12 http://www.ucolick.org/∼cnaw/sun.html Photometric catalogs were generated using SExtractor (see Tran et al. 2004) and we used k − correct to convert observed fluxes (measured with MAG_AUTO and corrected for Galactic extinction) to rest-frame BV magnitudes. As in the cases of the supergroup and cluster galaxies, stellar masses for the field galaxies are estimated using Eq. 1.
Spectra from all of the observing runs were reduced using a combination of IRAF 13 routines and custom software provided by D. Kelson (Kelson et al. 2000); see Tran et al. (2005) for further details on the spectral reductions. Redshifts were determined using IRAF cross-correlation routines, and each assigned redshift was visually compared to the 1D spectrum. Each redshift was then given a quality flag where Q = 3, 2, & 1 corresponded to definite, probable, and maybe (single emission line). The spectral range for most of the supergroup members covers [OII]λ3727 to [OIII]λ5007.
The spectroscopic completeness in the HST/ACS footprint is shown in Fig. 2. Due to the supergroup's elongated structure (see Fig. 1), spectra for a few of the bright (m ACS 814 < 18.5) galaxies have not been obtained; however, these are foreground galaxies. The brightest group galaxy has m ACS 814 = 17.5 mag, and the survey remains > 50% complete to m ACS 814 = 20.5 mag. For red supergroup members, the adopted magnitude limit used in our analysis of M V =−20.5 mag (see §2.2.2) corresponds to m ACS 814 = 20.4 mag. In the larger 20 ′ × 20 ′ region centered on the HST/ACS mosaic, we have obtained spectra of 603 unique objects. Guided by breaks in the redshift distribution, we define group members to be at 0.34 ≤ z ≤ 0.38 (Fig. 3). Considering only galaxies with redshift quality flag of Q = 3 gives 174 supergroup members. Four of the five X-ray luminous regions correspond to galaxy groups at 0.35 < z < 0.37 while the fifth is a galaxy cluster at z = 0.48 . The coordinates, mean redshifts, and line-of-sight velocity dispersions of the individual groups are listed in Table 1, and Fig. 1 shows the spatial distribution of members on the HST/ACS mosaic.
A complete description of the spectroscopic survey in CL1358 including target selection, spectral reduction, wavelength calibration, sky subtraction, etc., is presented by F98. To summarize, WHT and MMT spectroscopy targeted objects with R ≤ 21 mag over a 10 ′ × 11 ′ region; at this magnitude limit, the spectroscopic survey is > 80% complete and not dependent on color (see F98, Fig. 2). The magnitude limit corresponds approximately to M V =−20.5, and we use this limit for our luminosity-selected samples. For reference, the Coma cluster has M * V = −20.6 14 From nearly 400 redshifts, cluster membership was confirmed for 232 galaxies; in our analysis, we consider only the 171 members that fall on the HST/WFPC2 mosaic (total area of ∼ 502 ′ ) that were studied by H07.

Field Galaxies
As part of our program on galaxy clusters at intermediate redshift, we also obtained redshifts for a large sample of field galaxies. Spectroscopic targets were selected using a magnitude cut of m W FPC2 814 ≤ 23 and m W FPC2 814 ≤ 23.5 in the MS2053 and MS1054 fields respectively. These magnitude-limited spectroscopic surveys were completed with Keck/LRIS (Oke et al. 1995) and resulted in a total of over 800 redshifts in the two fields; excluding the cluster members and considering only redshifts with Q = 3 provides 295 field galaxies at 0.09 < z < 1.36. Further observational details for each field are in Tran et al. (2004). Notably, the spectroscopic completeness in both cluster fields is > 80% at m W FPC2 814 < 21 (see Figs. 2 in Tran et al. 2005. To ensure that we are observing the field galaxies at the same epoch as the group and cluster galaxies, we use only the field galaxies at 0.25 ≤ z ≤ 0.45 (z = 0.35; Fig. 4). In this redshift range, we have 87 field galaxies; applying a magnitude (M V <−20.5) or mass (M * ≥4 × 10 10 M ⊙ ) selection decreases the field sample to 28 and 21 galaxies respectively (see Table 2). Note that both galaxy clusters MS2053 (z = 0.59; Tran et al. 2005) and MS1054 (z = 0.83; Tran et al. 2007) are at higher redshift, thus our field sample is not contaminated with cluster galaxies.
In SG1120, we use the T-types assigned by Kautsch et al. (2008) to 142 of the 143 group galaxies that fall on the HST/ACS mosaic. The galaxy groups in SG1120 have velocity dispersions that are significantly lower than in massive clusters such as CL1358 (303 − 580 km s −1 vs. 1027 km s −1 ; see Table 1), yet the groups are already dominated by earlytype members: SG1120's early-type fraction of ∼ 70% is already comparable to that of galaxy clusters at intermediate redshifts .
For the cluster (CL1358) and field galaxies, we have Ttypes assigned by D. Fabricant, M. Franx, and P. van Dokkum using HST/WFPC2 imaging . This team classified all galaxies in the cluster fields brighter than m W FPC2 814 = 22; these classifications have been published in vD98, van Dokkum et al. (2000), and Tran et al. (2005). From this database, 161 of the 171 CL1358 galaxies in H07's sample and 67 of the 87 field galaxies (0.25 ≤ z ≤ 0.45) have visual classifications.
14 Here we use m * V = 14.5 (Abell 1977) and a distance modulus of (m−M) = 35.11 (Baum et al. 1997) for the Coma cluster.

MIPS 24µm Imaging
Deep wide-field 24µm imaging of all the fields presented in our study was taken with MIPS (Rieke et al. 2004). We briefly summarize here the procedure for retrieving, reducing, and analyzing the MIPS imaging; further details are in Paper I. For SG1120, we retrieved the MIPS 24µm data sets from the Spitzer archive and corrected the individual frames with the scan mirror position-dependent flats before combining the frames with the MOPEX software to a pixel size of 1.2 ′′ . The integration time per pixel was 1200 seconds and the background level 35.5 MJy/sr.
With the large SG1120 MIPS mosaic (22 ′ × 56 ′ ), we were able to determine a good point spread function (PSF) and thus measure 24µm fluxes via profile fitting. As a check, we compared the fluxes measured via profile fitting to aperture fluxes and found the values to be consistent; for the latter, we used an aperture diameter of 2 ′′ as a compromise between maximizing the flux and minimizing contamination from close neighbors, and applied corrections based on fluxes derived from modeled PSFs (see Paper I). We matched the centroid position of the MIPS sources to the master R-band catalog. We estimated the completeness of the SG1120 24µm catalog by adding 50 sources modeled on the empirical PSF to the mosaic and repeating this process 20 times.
To convert the 24µm fluxes to star formation rates, we determined the total infrared luminosity (F 8−1000µm ) for each galaxy using a family of infrared spectral energy distributions (SEDs) from Dale & Helou (2002). Using the range of SEDs that are representative of galaxies in the Spitzer Infrared Nearby Galaxies Survey (Dale et al. 2007), we adopt the median conversion factor from F 24µm to F 8−1000µm at z ∼ 0.37 where the SEDs give essentially the same values and the error due to the adopted conversion factor is only ∼ 10 − 20%. Combining this conversion with the completeness simulations, we estimate that the SG1120 24µm catalog is 80% complete to log(L IR )[erg s −1 ]=43.8; this corresponds to a 24µm flux of approximately 105µJy and an IR star formation rate of SFR IR ≥ 3M ⊙ yr −1 ).
For the smaller cluster fields, we followed essentially the same procedure except that we used APEX to measure fluxes within a 2 ′′ diameter aperture and corrected the measured fluxes using the PSF from the SG1120 mosaic. The total integration times and background levels in these mosaics vary, but the 24µm imaging is essentially confusion-limited in these fields (see Paper I). We estimated the completeness of the 24µm catalogs by adding 30 sources into each mosaic and repeated the process 20 times for each mosaic. The 24µm catalogs are deeper than in the SG1120 field, e.g. the CL1358 24µm catalog is 80% complete to log(L IR )[erg s −1 ]=43.5.

RESULTS
In the following analysis, we consider only the 143 supergroup galaxies in SG1120 that fall on the HST/ACS mosaic, the 171 cluster galaxies in CL1358 with photometry measured by H07 from the HST/WFPC2 mosaics, and the 87 field galaxies at 0.25 ≤ z ≤ 0.45, all of which also have HST/WFPC2 imaging. We are thus assured of a uniformly selected sample and can directly compare results across the three environments.
Our 24µm imaging identifies all galaxies with obscured star-formation rates of 3M ⊙ yr −1 or greater, regardless of galaxy mass. However, actively star-forming galaxies tend to have lower mass-to-light ratios than galaxies dominated by 24µm Sources in a Super Galaxy Group at z = 0.37 older stars, e.g. galaxies on the red sequence, thus an optical luminosity-selected sample is likely to differ from a massselected sample. For this reason, we use both luminosity and mass-selected samples in our analysis to check the robustness of our results. Note that in fitting the infrared luminosity functions, we select based on total IR luminosity as determined with the 24µm fluxes.

Fraction of 24µm Sources
We first apply a luminosity limit set by the spectroscopic completeness in the CL1358 field (see §2.2) and consider only galaxies with absolute V -band (Vega) magnitude brighter than −20.5; due to the mixed galaxy population in our z ∼ 0.35 samples, we do not correct for passive evolution. The colormagnitude (CM) diagram for the galaxies in all three environments is shown in Fig. 5. The slope of the CM relation each panel is from vD98 who measured the CM relation in CL1358 using the early-type members; the CM relation is normalized to the red sequence in CL1358. The color deviation from the CM relation is denoted as ∆(B − V ) where blue galaxies are classically defined has having ∆(B − V ) < −0.2 (Butcher & Oemler 1978).
In the luminosity-limited sample, the fraction of 24µm sources in the cluster is significantly lower than in the field (7±2% 15 vs. 36±9%; Table 2). However, we find that the fraction of 24µm sources in the supergroup (32 ± 5%) is comparable to the field and four times greater than in the cluster. Figure 6 shows HST/ACS images of the supergroup galaxies with SFR IR ≥ 3M ⊙ yr −1 and M V <−20.5.
Because ongoing star formation can increase a galaxy's total optical luminosity and thus scatter lower-mass systems into the luminosity-selected sample, we compare our results to a mass-selected sample ( Table 2). As in Paper I, we consider only galaxies with stellar masses greater than M * = 4 × 10 10 M ⊙ . Again, the fraction of 24µm sources in the supergroup is higher than in the cluster (19±5% vs. 5±2%; Table 2); however, the fraction in the supergroup is now only about half that of the field.
In applying a mass-cut, we discover that SG1120 has a significant number of members (17; see Fig. 6) that are bright (M V <−20.5), mostly late-type galaxies with stellar masses of (10.0 < log(M * )[M ⊙ ] < 10.6). Once star formation is quenched in these systems, they will fade and redden, and most will have L < L * , i.e. they will populate the faint end of the red sequence We have assumed that the 24µm emission is due to star formation but as many authors have noted (e.g. Donley et al. 2008), there is a likely contribution from dust-enshrouded active galactic nuclei (AGN). However, we stress that AGN contamination does not impact our conclusions because the relative fraction in each environment is small. In a 70 ksec Chandra/ACIS image of the SG1120 field , only 4 of the 143 group galaxies are detected as Xray point sources. As for the cluster galaxies (z = 0.33), Martini et al. (2007) estimate that the AGN fraction in two z ∼ 0.3 clusters is less than 3%. The possible number of field AGN is equally low: using Donley et al. (2008)'s survey of IR-detected AGN, we estimate that only one of the field IR sources can be an AGN. Note that if we account for these estimates of the AGN fraction, the difference in the IR star-forming fraction between the supergroup and cluster en-vironment only increases (∼ 28% vs. ∼ 4% in the luminosityselected samples).

Infrared Luminosity Function
To better quantify how the 24µm sources in the supergroup differ from their counterparts in the cluster and in the field, we compare the infrared luminosity functions (IR LFs) of the supergroup (SG1120) and cluster (CL1358) with published results from the field in Fig. 7. We correct the observations in both SG1120 and CL1358 for spectroscopic and 24µm incompleteness; the 80% completeness limit for the 24µm sources is log(L IR )[erg s −1 ]= 43.8 and 43.5 in the supergroup and galaxy cluster, respectively.
To fit the IR distribution, we follow Bai et al. (2006) and first use a Schechter function (Schechter 1976): where we fix the faint-end slope α and adopt the best-fit chisquare minimization method to determine L * and φ * ( Table 3). Because studies show the IR LF in general has a relatively large number of bright sources and is better described by a double-exponential function (Le Floc'h et al. 2005), we adopt their approach and also fit a double-exponential function: (3) where we fix the constants α and σ to the values measured for the field IR LF, and minimize with chi-square again. Note that to determine the faint-end slope α in either function, deeper IR observations are required (Bai et al. 2006).
The IR LFs in both the supergroup and galaxy cluster are well-fit by both a Schechter and a double-exponential function (Fig. 7) using different values for L * and φ * (see Table 3). However, the density of IR sources in the supergroup is dramatically higher than in the cluster, especially at log(L IR )[erg s −1 ]> 45.
Perhaps the large difference is due to CL1358 being unusually deficient in IR sources. We test this by taking the IR LF determined from 24µm observations of galaxy clusters at z ∼ 0 (Bai et al. 2006) and evolve the IR LF to z ∼ 0.35 using the observed evolution in the field IR LF (Le Floc'h et al. 2005). Bai et al. (2007Bai et al. ( , 2008 find that L * IR and φ * IR , as derived from 24µm observations, evolve in approximately the same manner in galaxy clusters and the field to z ∼ 0.8, although see Muzzin et al. (2008) for an alternative result. The IR LF for CL1358 (z = 0.33) is consistent with the evolved cluster IR LF ( Fig. 7; long-dashed curve). However, the density of IR sources in the supergroup is ∼ 10 times higher than the number predicted from the evolved IR LF at log(L IR )[erg s −1 ]> 45.
The IR LF in the supergroup also differs from the IR LF for the field measured by Le Floc'h et al. (2005). The best-fit double-exponential function to the group IR LF has a measurably larger L * IR compared to the value for field galaxies at (0.3 < z < 0.45): 43.71 +0.19 −0.19 vs. 43.28 +0.09 −0.03 . In comparison, Bai et al. (2008) find a similar L * IR value for both local cluster and field galaxies. (Table 3). SG1120's higher L * IR relative to even that measured for the field suggests that star formation is enhanced in the group environment.
To summarize, the number of IR sources in the galaxy groups that make up SG1120 is significantly higher than in CL1358, a rich galaxy cluster at z = 0.33, and includes a population of very bright IR sources (log(L IR )[erg s −1 ]> 45) that are not found in CL1358 nor in lower redshift clusters. The higher value of log(L * IR ) in the supergroup compared to that measured for the field at (0.3 < z < 0.45) also indicates that the IR sources in the supergroup differ from their counterparts in the field, i.e. that star formation likely is enhanced in the group environment.

Local Environment
Having established that the 24µm population in the supergroup (SG1120) is different from that in the cluster (CL1358) and likely also the field environment, we examine how the galaxy populations for the luminosity-selected samples (M V <−20.5) depend on local environment, i.e. how star formation rate relates to galaxy density (Balogh et al. 1998;Gómez et al. 2003). In addition to the IR-bright population, we separate galaxies into optically-defined absorptionline ([OII]λ3727 < 5Å) and emission-line ([OII]λ3727 ≥ 5Å) systems. In the supergroup and cluster, we define the local galaxy density Σ using the distance to the 10 th nearest spectroscopically confirmed neighbor; we note that the following results do not change if we use instead the 5 th nearest neighbor. Figure 8 (top left) shows how the fraction of SFR IR ≥ 3M ⊙ yr −1 members in the supergroup steadily increases with decreasing local density. The increasing fraction of emissionline members with decreasing Σ mirrors the trend for IR members, and the absorption-line fraction changes accordingly. The trend of an increasing IR fraction with decreasing local density remains even if we apply higher IR star formation rate threshold of 5 M ⊙ yr −1 . In contrast, the IR population in the cluster (top right) shows essentially no trend with local environment: the absorption-line population dominates throughout the range of local densities explored here (5 < Σ < 70 gal Mpc −2 ) 16 . These results are in line with Paper I where we also find an increase in 24µm members outside the cores of massive clusters (R pro j > 700h −1 kpc). Our results argue for a physical mechanism that quenches star formation before the members reach the group cores.
At the lowest galaxy densities, the fraction of IR members in the supergroup is higher than even in the field: considering all members with Σ < 20 gal Mpc −2 , the IR fraction increases to 49% (26/53) compared to the field value of 38% (Table 2). At these low galaxy densities, the higher fraction of IR members in the supergroup relative to the field, while statistically not significant, is consistent with the higher L * IR measured in the group environment (see Table 3).
If we now examine how the mass-selected (M * >4 × 10 10 M ⊙ ) samples depend on local environment (Fig. 8, bottom panels), the trend of increasing 24µm fraction with decreasing galaxy density in the supergroup is weaker: the absorption-line population dominates in both the group and the cluster environment at Σ > 10 gal Mpc −2 . These results are consistent with H07 who find that evolution in the earlytype fraction in massive clusters (0 < z < 0.8) is weaker when considering only galaxies with M * >4 × 10 10 M ⊙ versus a luminosity-selected sample, thus galaxies with lower masses play a significant role in the observed evolution of the cluster galaxy population. 16 While the galaxy density in the cluster environment extends to Σ > 100 gal Mpc −2 , we consider only the Σ range that overlaps with the galaxy groups.
We find that the supergroup has a population of luminous, 24µm detected late-type members with stellar masses of (10.0 < log(M * )[M ⊙ ] < 10.6) that are mostly outside the groups' cores (see Figures 1 & 6), i.e. at lower galaxy densities. It is these galaxies that cause the strong observed trend of decreasing 24µm fraction with increasing galaxy density in the luminosity-selected sample. As noted in §3.1, most of these will fade and redden and can populate the faint end of the red sequence once their star formation is quenched.

DISCUSSION
To study how galaxies evolve, galaxies are usually separated into active/emission-line and passive/absorption-line systems with the goal of isolating the physical process that connects the two phases, e.g. removal of a galaxy's gas halts its star formation and the galaxy evolves from an active system into a passive one (Abadi et al. 1999;Kawata & Mulchaey 2008). The high fraction of 24µm galaxies in the supergroup and the field means that obscured star formation (IR phase) is important for at least 30% of optically-selected galaxies in both these environments; the IR phase is likely to be as important in clusters given that clusters grow via the accretion of field and group galaxies (Peebles 1970). In the following, we examine the physical properties of the 24µm galaxies to better understand where the IR population fits into our current picture of galaxy evolution.

Morphologies of 24µm Galaxies
To determine what the typical morphology of a 24µm galaxy is across environment at z ∼ 0.35, we separate the samples into late-type (T > 0; disk-dominated) and early-type (T ≤ 0; bulge-dominated) systems. In both the field and supergroup, most (> 60%) of the late-type galaxies are IR-detected; the supergroup even has a few bulge-dominated systems that are IR-detected 17 . In contrast, only ∼ 30% of the late-type galaxies in the cluster are IR-detected, and none of the cluster's bulge-dominated members are IR-detected. These differences are true in both the luminosity and mass-selected samples ( Table 2).
Comparing the 24µm galaxies in the supergroup to the field again strongly suggests a difference between the two environments. In the luminosity-selected samples, the supergroup and field have similar 24µm fractions. However, the supergroup has a much higher fraction of E/S0 members: the E/S0 fraction in the supergroup is ∼ 60% but it is only ∼ 30% in the field 18 ( Table 2). Several of the 24µm galaxies in the supergroup are in merging, disk-dominated systems (see Fig. 6), but only one of the massive dissipationless merging pairs ) is detected at 24µm.
Despite having double the fraction of early-type galaxies compared to the field, the supergroup has a high 24µm fraction due to a population of luminous (M V <−20.5), lowmass (M * <4 × 10 10 M ⊙ ) late-type members with SFR IR ≥ 3M ⊙ yr −1 (see §3.1, §3.3, & Table 2). Note that in the massselected sample, the 24µm fraction in the supergroup drops from ∼ 32% to ∼ 19% while the field fraction remains high (∼ 35%). Our results show that the timescales for morpho-logical evolution and quenching of star formation must differ (see also Finn et al., in prep).

Star Formation on the Red Sequence
Across all three environments, there are 24µm sources that are also on the optical red sequence (see Fig. 5); here we use the classical definition of the red sequence as galaxies with ∆(B − V ) > −0.2 (Butcher & Oemler 1984). The fraction of red 24µm sources depends on environment: it is highest in the field (21%; 6/29), decreases in the supergroup (7%; 7/98), and is is lowest in the cluster (3%; 3/105). These results do not depend on whether we use the luminosity or mass-selected sample.
Our results appear to conflict with Gallazzi et al. (2008) who find that the fraction of red 24µm sources in the Abell 901/902 supercluster (z = 0.165; Gray et al. 2002) peaks at intermediate densities typical of cluster outskirts and galaxy groups 19 . However, the authors estimate the field contamination in their magnitude range can be as high as 20%, and our study shows that the fraction of red 24µm galaxies in the field is ∼ 3 times higher than in the groups. By using a spectroscopically selected sample, we circumvent possible problems due to field contamination.
In A901/902, Wolf et al. (2009) find that the dusty red starforming members are primarily spiral galaxies, and that this population mostly overlaps with the "optically passive" spirals (as defined by color). We find similar results: all of the red 24µm galaxies in the field and cluster are disk-dominated systems (T > 0), and most of the red 24µm galaxies (∼ 60%) in the groups are spirals as well (see Fig. 6).
We note that neither optical colors nor optical spectroscopy reliably identifies dusty red [∆(B − V ) > −0.2] star-forming spirals: summing across environment, only (10/15) of the red 24µm galaxies have [OII]λ3727 > 5Å, i.e. one third of 24µm members on the red sequence show no significant [OII] emission. This result is in line with earlier studies, e.g. Moustakas et al. (2006), that show optical spectroscopy can severely underestimate the level of activity. On a related note, many 24µm galaxies can be strongly extincted with E(B-V) values as high as 0.6 (Cowie & Barger 2008); once corrected for extinction, many of the 24µm galaxies would not lie on the red sequence.

Progenitors of Faint Red Galaxies
As the groups in SG1120 merge to form a galaxy cluster, how do the 24µm members impact the overall galaxy population? In Fig. 9, we plot specific star formation rates (defined as SFR IR ≥ 3M ⊙ yr −1 divided by stellar mass) versus stellar mass for the 24µm galaxies in the supergroup, field, and cluster. Assuming the 24µm members maintain their current star formation rates, perhaps only five out of the 72 massive (M * >4 × 10 10 M ⊙ ) group galaxies will double their stellar masses, i.e. virtually all of the massive galaxies that will end up in the cluster are already in place.
In our analysis, we have identified a considerable number of luminous (M V <−20.5) galaxies in the supergroup that have SFR IR ≥ 3M ⊙ yr −1 and stellar masses below 4 × 10 10 M ⊙ (17/98; Table 2); it is this population that contributes the most to the difference between the 24µm population in 19 Because Gallazzi et al. (2008) estimate local galaxy density differently, we cannot compare their values directly to Fig. 8. These authors use a spectroscopic sample supplemented with members selected with photometric redshifts. the supergroup and in the cluster. Fig. 9 shows that even if these group galaxies can maintain their current star formation rates, most (15/17; ∼ 90%) will still have stellar masses of M * < 10 11 M ⊙ at z ∼ 0; the current average stellar mass for all 17 galaxies is log(M * )[M ⊙ ] = 10.4. Note that most of these members are at R pro j > 0.5 Mpc from their respective group cores (see Fig. 1).
We test our hypothesis that these galaxies can evolve into (L < L * ) red galaxies by comparing their stellar masses to the faint red galaxies in CL1358, our massive galaxy cluster. Following De Lucia et al. (2007), we define faint red galaxies as having luminosities of (0.1 − 0.4)L * 20 and ∆(B −V ) > −0.2. The average stellar mass of the faint red galaxies in the cluster is log(M * )[M ⊙ ] = 10.3; this is comparable to the average stellar mass of the luminous, low-mass, SFR IR ≥ 3M ⊙ yr −1 supergroup galaxies. Assuming their star formation is quenched by z ∼ 0, these supergroup galaxies will fade and redden to lie on the CM relation in less than a Gyr (see models by Bruzual & Charlot 2003). Their younger luminosity-weighted ages relative to the more massive galaxies will be consistent with the observed age spread in the Coma cluster (Poggianti et al. 2001). We stress that the luminous, low-mass, SFR IR ≥ 3M ⊙ yr −1 supergroup galaxies are likely to be only one of multiple progenitors of faint red galaxies.

SUMMARY
To quantify how dust-obscurred star formation varies with environment, we compare galaxies in a super galaxy group to those in the field and in a massive cluster at z ∼ 0.35 using a rich multi-wavelength dataset that includes imaging from Hubble (optical), Chandra (X-ray), and Spitzer (24µm). The strength of our work relies on extensive optical spectroscopy in our fields: the magnitude-limited spectroscopic surveys yielded a total of over 1800 unique redshifts and enable us to securely identify field, supergroup, and cluster members. We focus on the four X-ray luminous galaxy groups at z ∼ 0.37 (SG1120-12) that will merge to form a galaxy cluster comparable in mass to Coma ; the groups have line-of-sight velocity dispersions of 303 − 580 km s −1 . To ensure robust comparison, we consider only field galaxies at 0.25 ≤ z ≤ 0.45 (z = 0.35) and confirmed members of the massive galaxy cluster CL1358+62 (z = 0.33, Fisher et al. 1998).
We find that the supergroup has a significantly higher fraction of dusty star-forming members than the massive galaxy cluster: in the luminosity-selected (M V <−20.5) samples, 32% of the supergroup members have SFR IR ≥ 3M ⊙ yr −1 compared to only 7% of the cluster members. The supergroup's infrared luminosity function confirms that the density of IR sources is dramatically higher in the groups compared to the cluster. The supergroup members also include bright IR sources (log(L IR )[erg s −1 ]> 45) not found in galaxy clusters at z 0.35.
When selected by luminosity, the supergroup members show a strong trend of decreasing 24µm fraction with increasing local galaxy density, i.e. an infrared-density relationship. This mirrors the trend in the optically active members (as defined by [OII] emission). In contrast, the fraction of 24µm sources in the massive cluster stays essentially zero at all densities.
Comparison to the mass-selected (M * >4 × 10 10 M ⊙ ) sam-20 Using M * V = −20.6, the corresponding Vega V -band magnitudes are (−20.1 < M V < 19.6). ples reveals that the higher 24µm fraction and the IR-density relation in the supergroup is due primarily to a population of luminous (M V <−20.5), lower-mass (10.0 < log(M * )[M ⊙ ] < 10.6), late-type members with SFR IR ≥ 3M ⊙ yr −1 (∼ 17%). Most of these members are outside of the group cores (R pro j ≥ 0.5 Mpc). Assuming their star formation is quenched in the next ∼ 3−4 Gyr, these members will fade and redden by z ∼ 0, and most will become fainter (L < L * ) galaxies on the colormagnitude relation. The physical mechanism that quenches their star formation must be effective outside the group cores, i.e. in lower density environments.
In the supergroup, the excess of 24µm sources, the number of very bright 24µm members, and the infrared-density relationship is surprising because the E/S0 fraction is already as high as in the cluster (> 60% for luminosity-selected sample; Kautsch et al. 2008). No further morphological evolution is required to bring the morphological distribution of the groups in line with the high early-type fractions observed in local galaxy clusters. In other words, the timescale for morphological transformation must not be strongly coupled to when star formation is completely quenched.
Our analysis indicates that the 24µm population in the supergroup differs even from the field: 1) the supergroup's IR luminosity function has a measurably higher L * IR than the field; and 2) the E/S0 fraction in the supergroup is twice that of the field, yet the 24µm fraction in both environments are comparable. If dusty star formation is enhanced in the supergroup relative to the field, our IR-density analysis suggests that it occurs at densities of Σ < 20 gal Mpc −2 . A larger field sample selected with the same criteria as in the supergroup and cluster is needed to answer this question securely.
Our study highlights the importance of understanding galaxy evolution on group scales. A significant fraction ( 30%) of optically selected galaxies in both the supergroup and field at z ∼ 0.35 have dust-obscured star formation; the IRphase must be as important in clusters because clusters grow by accreting galaxy groups and field galaxies. As demonstrated in recent simulations of galaxy groups (e.g. Hester 2006;Romeo et al. 2008;McCarthy et al. 2008;Kawata & Mulchaey 2008), the physical mechanisms that affect star formation and induce morphological evolution are already well underway in the galaxy groups that make up SG1120. We will continue dissecting how these galaxies are transformed by using recently obtained IFU observations to map the kinematics and star formation of the 24µm members.

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a For the Schechter profile, we set α IR = 1.414 (Bai et al. 2006(Bai et al. , 2007. b IR LF measured in z ∼ 0 galaxy clusters (Bai et al. 2006) evolved to z ∼ 0.35 using the evolution measured in the field IR LF (Le Floc'h et al. 2005). c For the double-exponential profile, we set α IR = 1.23 and σ IR = 0.72 (Le Floc'h et al. 2005) d Field L * IR measured by Le Floc'h et al. (2005) for galaxies at (0.3 < z < 0.45); we do not include φ * IR because it is normalized differently in the field and in clusters.
FIG. 1.-Spatial distribution of the spectroscopically confirmed supergroup galaxies that fall on the HST/ACS mosaic (small dots; 0.34 ≤ z ≤ 0.38; 143 of 174 total members); north is up and east to the left. Left: Open circles denote all group galaxies brighter than M V =−20.5, and large filled circles denote members brighter than M V =−20.5 with SFR IR ≥ 3M ⊙ yr −1 . The large dashed circles (R = 0.5h −1 Mpc) are centered on the brightest group galaxies listed in Table 1; these positions are well-matched to the extended X-ray emission detected in each group. Right: The same but for the mass-selected (M * > 4 × 10 10 M ⊙ ) supergroup members. Note the number of luminous, 24µm members outside the group cores that are low-mass systems ( FIG. 4.-The redshift distribution of the field galaxies at 0.1 < z < 0.6. These field galaxies are drawn from extensive spectroscopic surveys of two galaxy clusters, both of which are at z > 0.55. The redshift bins are ∆z = 0.002 and are 5 − 10 times larger than the typical error in the redshift; any apparent overdensities disappear with smaller bin sizes. The vertical dashed lines denote the redshift range of our comparison field sample (0.25 ≤ z ≤ 0.45).
FIG. 5.-Rest-frame color-magnitude (CM) diagram of galaxies brighter than M V =−20.5 (dashed vertical line) in the field (top), galaxy groups (middle), and galaxy cluster (bottom). The galaxies with SFR IR ≥ 3M ⊙ yr −1 are shown as filled circles; the corresponding fractions are listed in each panel. The slope of the CM relation (solid line) is from vD98 and normalized to the red sequence in CL1358; the same CM relation is shown in each panel. The 24µm fraction in the galaxy groups is nearly as high as in the field and is four times higher than in the cluster.
FIG. 8.-Relative fraction of absorption-line (circles), emission-line (triangles), and 24µm (squares) members as a function of local galaxy surface density in the groups (SG1120, left panels, filled symbols) and cluster (CL1358, right panels, open symbols); the points are offset slightly in log Σ for clarity. The top panels show the luminosity-selected (M V <−20.5) members and the bottom panels the mass-selected (M * >4 × 10 10 M ⊙ ) members. The long-dashed, dotted, and shortdashed horizontal lines show respectively the absorption-line, emission-line, and 24µm fractions in the field. Only in the supergroup with a luminosity-selected sample does the 24µm fraction increase steadily with decreasing local density, i.e. an infrared star formation-density relation.
FIG. 9.-The 24µm-derived specific star formation rate (SSFR) in units of Gyr for the galaxy groups (filled squares), cluster (open squares), and the field sample (stars); the dotted vertical line denotes M * =4 × 10 10 M ⊙ and we show only members with SFR IR ≥ 3M ⊙ yr −1 . If we assume constant SF rates, the galaxies above the horizontal line will more than double their stellar mass from z ∼ 0.37 to now. At most, perhaps four of the 72 massive galaxies in SG1120 can double their stellar mass. However, there are a number of lower-mass (10.0 < log(M * )[M ⊙ ] < 10.6) group galaxies that can grow substantially in stellar mass (upper left quadrant); these are luminous late-type members with SFR IR ≥ 3M ⊙ yr −1 .