EDP Sciences
Free Access
Issue
A&A
Volume 540, April 2012
Article Number A96
Number of page(s) 17
Section Extragalactic astronomy
DOI https://doi.org/10.1051/0004-6361/201117281
Published online 04 April 2012

© ESO, 2012

1. Introduction

Hickson compact groups (HCGs) (Hickson 1982) are dense and relatively isolated groups of four to eight galaxies in the nearby universe. The combination of high galaxy density (Hickson 1982) and low-density environment coupled with low systemic velocity dispersions (⟨σ⟩  = 200 km s-1, Hickson et al. 1992) make HCGs especially interesting systems for studying how gas content and star formation activity in galaxies are influenced by the environment.

The most remarkable effect of multiple and strong interactions between galaxies in HCGs involves an atomic gas (HI) deficiency. VLA measures of individual spiral galaxies in HCGs show they only have 24% of the atomic hydrogen (HI) mass, MHI, expected from their optical luminosities and morphological types (Verdes-Montenegro et al. 2001). The inferred deficiency becomes even greater if one assumes that many of the lenticular galaxies, which are overrepresented in HCGs, are stripped spirals. Verdes-Montenegro et al. (2001) propose an evolutionary sequence for HCGs in which the HI is continuously removed from the galaxies, finally leading to groups where most of the HI is located outside of the galaxies. However, not only are the individual galaxies in HCGs HI deficient, but also HCGs are so as a whole (Verdes-Montenegro et al. 2001). This leads to the still open question of where the missing HI has gone and by which mechanism it was removed. To investigate the role played by a hot intragroup medium (IGM), Rasmussen et al. (2008) performed Chandra and XMM-Newton observations to study eight of the most HI deficient HCGs. They find bright X-ray emission in only four groups suggesting that galaxy-IGM interactions are not the dominant mechanism driving cold gas out of the galaxies. Borthakur et al. (2010) find with new single-dish Green Bank Telescope (GBT) observations of HGCs an important diffuse, low-column density intragroup HI component, missed by VLA observations. Taking these components into account reduced, but did not completely eliminate, the HI-deficiency of the groups.

The effect of an extreme environment on molecular gas properties is controversial. An enhancement of molecular gas content with respect to isolated galaxies has been reported for strongly interacting systems (Casasola et al. 2004, and references therein), defined in that work as galaxies appearing to be clearly interacting with nearby objects, presenting tidal tails or bridges, and merging systems and galaxies with disturbed structures. With respect to galaxies in clusters, no deficiency in the molecular gas content has been found in any global studies of the Virgo cluster (Kenney & Young 1986; Boselli et al. 2002) and the Coma Supercluster (Casoli et al. 1991; Boselli et al. 1997) in spite of the large HI deficiencies that some galaxies presented. The spatially resolved study of Fumagalli et al. (2009) found, however, that a significant number (~40%) of HI-deficient spiral galaxies were also depleted in molecular gas, if the HI was removed from within the optical disk. Scott et al. (in prep.) find a trend for spirals in Abell 1367 in more evolved evolutionary states to be more depleted in MH2 than those in less evolved evolutionary states. Thus, there are indications that the cluster environment does affect the molecular gas content.

Observations of the molecular gas in galaxies in HCGs have given until now contradictory results. Leon et al. (1998) found the MH2/LB ratio of galaxies in HCGs to be enhanced with respect to a sample of field and interacting galaxies. In contrast, Verdes-Montenegro et al. (1998) found no evidence for an enhancement of the molecular gas mass (MH2) in HCG galaxies relative to a sample of isolated galaxies. Studying the relation between atomic and molecular gas for a sample of 32 spiral galaxies, Verdes-Montenegro et al. (2001) have found tentative evidence of a depressed molecular gas content in HI deficient galaxies in HCGs.

The level of star formation (SF) in HCGs has also been subject to considerable debate with original claims of a far-infrared (FIR) excess (Hickson et al. 1989), which was subsequently challenged (Sulentic & de Mello Rabaca 1993). From the enhanced SF observed in galaxy pairs (Xu & Sulentic 1991), an increase in SF in HCGs might be expected as a consequence of the continuous encounters and tidal interactions that take place within such groups. Nevertheless, the star formation rate (SFR) in HCGs, obtained from FIR (Verdes-Montenegro et al. 1998), mid-infrared (Bitsakis et al. 2010), a combination of mid-infrared and ultraviolet (Tzanavaris et al. 2010), and Hα luminosities (Iglesias-Páramo & Vílchez 1999) has been found to be similar to those of the control samples.

There have been a few attempts to study the relation of MH2 and LFIR with the HI properties of the HCG galaxies. Based on CO, FIR, and HI single-dish data, together with VLA maps for eight groups, Verdes-Montenegro et al. (2007) found that the MH2 and LFIR are lower than expected for HI deficient galaxies, when compared to a well-defined sample of isolated galaxies (AMIGA project, Analysis of the interstellar Medium of Isolated GAlaxies, http://amiga.iaa.es; Verdes-Montenegro et al. 2005). A possible explanation for this trend is that, since HI is needed to replenish the molecular clouds and molecular gas is necessary to fuel SF, an HI deficiency can ultimately lead to a decrease in the SFR. However, the result of Verdes-Montenegro et al. (2007) was based on a small sample of galaxies that does not cover the wide range of properties of HCGs and was therefore not statistically significant. On the other hand, while previous works studying MH2 and SFR of galaxies in HCGs (Verdes-Montenegro et al. 1998; Leon et al. 1998) were based on larger samples, those did not have the HI mass of the individual galaxies to compare with MH2 and LFIR. Thus, to date, no study of the relation between MH2, MHI  and the SFR for a statistically significant sample has been carried out.

To shed light on the relations between MH2 and SFR with MHI properties of the HCGs, we present here a systematic study of galaxies in a sample of 20 HCGs for which we have HI measurements for the entire groups, as well as for a large fraction of the individual galaxies. This enables us to consider MHI of the galaxies as an additional parameter, as well as the evolutionary phase of the group according to Verdes-Montenegro et al. (2001) (see Sect. 2). We compare the properties of galaxies in HCGs with those of isolated galaxies in the AMIGA sample (Verdes-Montenegro et al. 2005). Our goal is to determine whether deviations in the HI content with respect to isolated galaxies translate into anomalies in the MH2 and the SFR.

The outline of this paper is as follows. We present the sample in Sect. 2. CO(1–0) and CO(2–1) data coming either from our observations or from the literature, together with reprocessed IRAS FIR data, are presented in Sect. 3. In Sect. 4 we compare the MH2, LFIR (as a tracer of the SFR), and MHI of the galaxies, studying their deficiencies and their relation with the HI content and evolutionary phase of the group. A discussion of a possible evolutionary sequence for the molecular gas content in the HCGs is presented in Sect. 5. Finally, the conclusions of our work are summarized in Sect. 6.

2. The samples

2.1. Galaxies in HCGs

Our sample was selected from the revision of the original Hickson (1982) catalog performed by Hickson et al. (1992). From the groups included in that work, we studied 86 galaxies belonging to 20 different HCGs: 7, 10, 15, 16, 23, 25, 30, 31, 37, 40, 44, 58, 67, 68, 79, 88, 92, 93, 97, and 100. The groups, which cover all evolutionary stages and a wide range of HI deficiencies, satisfy the following criteria:

  • have at least four members, so triplets are excluded, according tothe original Hickson (1982) criterion. We alsoexclude false groups, where a single knotty irregular galaxy hasbeen confused with separated galaxies (Verdes-Montenegroet al. 2001);

  • contain at least one spiral galaxy, since we are mainly interested in studying the relation between the SF process and MH2, which are most clearly linked for spiral galaxies;

  • are at a distance D ≤ 100 Mpc (assuming H0 = 75 km s-1 Mpc-1), in order to guarantee a good sensitivity and to minimize possible source confusion within the telescope beam. At 100 Mpc, the 30 m beam would have a size of 10.7 kpc and the VLA beam (considering a size of 50′′ × 50′′) 24.2 kpc.

The HCGs in our sample cover the full range of HI contents. Their deviation from normalcy is measured with respect to that of isolated galaxies, as given by Haynes & Giovanelli (1984).

Table 1

Basic parameters of the galaxies in the HCG sample.

This deviation is usually referred to as deficiency and is defined as the decimal logarithm of the ratio between the sum of the expected HI masses for all the galaxies in the group based on their optical luminosity and morphology and of the HI mass of the entire group as derived from the single-dish observations in Borthakur et al. (2010; see Verdes-Montenegro et al. 2001, and also Sect. 4.2). As a function of their total HI deficiency, the HCGs in our sample can be classified as:

  • HCGs with a normal HI content (at least 2/3 of its expected value):HCG 23, 25, 68, and 79;

  • HCGs with a slight HI deficiency (between 2/3 and 1/3 of the expected value): HCG 7, 10, 15, 16, 31, 37, 40, 58, 88, 92, 97, and 100;

  • HCGs with a large HI deficiency (under 1/3 of the expected value): HCG 30, 44, 67, and 93.

Verdes-Montenegro et al. (2001) propose an evolutionary sequence model where the HI is continuously stripped from the galaxies. According to this model, HCGs can be classified into three phases as follows. In Phase 1 the HI is mainly found in the disks of galaxies. In Phase 2, 30% to 60% of the HI has been removed from the disks by tidal interaction. Finally, in Phase 3, almost all the HI is found outside of the galactic disks, either forming tidal tails of stripped gas (Phase 3a) or, in a few cases, in a large HI cloud with a single velocity gradient in which the galaxies are embedded (Phase 3b).

According to the evolutionary phases defined in Verdes-Montenegro et al. (2001), the HCGs in our sample were classified by Borthakur et al. (2010) as

  • Phase 1: HCG 7, 23, 67, 79, and 88;

  • Phase 2: HCG 10, 16, 25, 31, 401, 58, and 100;

  • Phase 3: HCG 15, 30, 37, 44, 68, 92, 93, and 97.

The evolutionary state is an indicator of the evolution of the cold ISM of the group, but it does not necessarily give the age of the group; e.g. HCG79 consists of three early-type galaxies and one intruding spiral galaxy. Stellar halo data indicates that it is an old group (Durbala et al. 2008). However, since the main part of the HI is located within the disk of the intruder galaxy, it is classified in evolutionary phase 1.

We revised the velocities of the individual galaxies in the HCGs of our sample. Two galaxies not considered in Hickson et al. (1992) have been added: HCG100d, which had no velocity data in that work, and HCG31g, added to the catalog of HCGs by Rubin et al. (1990).

The basic properties of the galaxies in our sample are detailed in Table 1. The columns are

  • 1.

    Galaxy: galaxy designation, following the notation ofHickson (1982).

  • 2.

    V: heliocentric radial velocity in km s-1 (weighted average of optical measurements taken from the LEDA2 database) converted from the optical to the radio definition for comparison with the CO spectra.

  • 3.

    σV: velocity dispersion of the galaxies in the group.

  • 4.

    D: distance to the corresponding HCG in Mpc, derived from the mean heliocentric velocity of the group as D = V/H0, assuming a value of H0 = 75 km s-1 Mpc-1. The mean velocity of the group is calculated averaging the velocity of the individual galaxies (Col. 2).

  • 5.

    T: morphological type taken from LEDA, following the RC3 classification (de Vaucouleurs et al. 1991).

  • 6.

    D25: optical major diameter in arcmin at the 25 mag arcsec-2 isophot taken from LEDA.

  • 7.

    : apparent blue magnitude taken from LEDA, corrected for Galactic dust extinction, internal extinction, and K-correction.

  • 8.

    log(LB): decimal logarithm of the blue luminosity, derived from as (1)This definition provides an estimate of the blue luminosity (νLν) at 4400 Å.

  • 9.

    log(LK): decimal logarithm of the luminosity in the K-band in units of the solar luminosity in the KS-band (LK, ⊙  = 5.0735 × 1032 erg s-1), calculated from the extrapolated magnitude in the KS (2.17 μm) band from the 2MASS Extended Source Catalog (Jarrett et al. 2000). We calculated the KS luminosity, LK, from the total (extrapolated) KS flux, fK, as LK = νfK(ν) (where ν is the frequency of the K-band, 1.38  ×  1014 Hz) .

  • 10.

    log(MHI): logarithm of the mass of the atomic hydrogen, in solar masses, for 66 of the galaxies in our sample observed with the VLA, using different combinations of the C and D configurations with beam sizes ranging from 16′′ × 14′′ to 72′′ × 59′′ (Verdes-Montenegro et al. 2001; and Verdes-Montenegro, priv. comm.).

2.2. Reference sample: isolated galaxies

We chose the AMIGA sample of isolated galaxies (Verdes-Montenegro et al. 2005), which is based on the CIG catalog (Karachentseva 1973), as a reference for the FIR and molecular gas properties. The FIR properties of an optically complete subsample of this catalog have been studied in Lisenfeld et al. (2007). The data that we are using for this subsample are slightly different from Lisenfeld et al. (2007) because we have taken a recent update of some basic properties into account that affects the blue magnitude, distance, morphological type, and isolation degree (detailed information provided in Fernández-Lorenzo 2012). We also used the CO data of a velocity restricted subsample (1500 < V < 5000 km s-1) of 173 AMIGA galaxies (Lisenfeld et al. 2011) to compare to our galaxies. The data for LK is also taken from Lisenfeld et al. (2011). When we quote mean values of the AMIGA galaxies we derived them for the subsample of galaxies with MH2 data. For the analysis of the HI properties we used the work of Haynes & Giovanelli (1984), which presents the observations and analysis of CIG galaxies, as a reference.

thumbnail Fig. 1

Morphological types, T(RC3), for the AMIGA isolated (black line) and HCGs (red filled bars) galaxies.

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There are two intrinsic differences between the AMIGA and the HCG samples that must be taken into account when performing a comparison: (i) the HCG sample has a higher rate of early-type galaxies (~45%), whereas 87% of the AMIGA galaxies are spirals (Fig. 1); and (ii) the sample of AMIGA galaxies with MH2 data is restricted to a velocity range of 1500 < V < 5000 km s-1, while the range of the HCGs extends to higher velocities. Thus, the isolated galaxies are, on average, at a shorter distance (47 Mpc average distance versus 68 Mpc for the HCG sample), which can explain their lower average luminosities (see Table 6). However, the values of the deficiencies, ratios or correlations, that we are going to discuss in the following are not expected to be affected by the difference in distance.

3. The data

3.1. CO data

We obtained CO data either by our observations or from the literature, for 86 galaxies in the selected 20 HCGs. CO data are missing for only two galaxies in these 20 groups, HCG 67d and HCG 92f. The CO(1–0) line was detected for 45 galaxies.

3.1.1. IRAM CO(1–0) and CO(2–1) observations and data reduction

We observed 47 galaxies belonging to 14 different HCGs. The observations of the CO rotational transition lines J = 1 → 0 and J = 2 → 1 (at 115.271 and 230.538 GHz, respectively) were carried out with the IRAM 30 m radio telescope at Pico Veleta3 during June, October, and December 2006. We performed single-pointing observations using the wobbler switch mode, with a switch frequency of 0.5 Hz and a throw of 200′′. We checked for all the objects that the off-position did not coincide with a neighboring galaxy.

The dual polarization receivers A100 and B100 were used to observe the CO(1–0) and CO(2–1) lines in parallel. The median system temperature was 231 K for the CO(1–0) observations, with  ~80% of the galaxies observed with system temperatures between 150 and 350 K. In the case of CO(2–1), the median system temperature was 400 K, with a temperature range between 230 and 800 K for 85% of the galaxies. For CO(1–0) the 1 MHz filterbank was used, and for CO(2–1) the 4 MHz filterbank. The corresponding velocity resolutions were 2.6 km s-1 for CO(1–0) and 5.3 km s-1 CO(2–1), respectively. The total bandwidth was 1 GHz. The half power beam width (HPBW) is 22′′ and 11′′ for 115 and 230 GHz, respectively. All CO spectra and intensities are presented on the main beam temperature scale (Tmb) which is defined as . The IRAM forward efficiency, Feff, was 0.95 at 155 GHz and 0.91 at 230 GHz and the beam efficiency, Beff, was 0.75 and 0.54, respectively.

Table 2

Observed and derived molecular gas properties.

The data reduction and analysis was performed using CLASS, while further analysis used GREG, both part of the GILDAS4 package developed by IRAM. First we visually inspected the spectra and discarded bad scans. Then, spikes were removed and a constant or linear baseline was subtracted from each spectrum. The scans were then averaged to achieve a single spectrum for each galaxy and transition. These spectra were smoothed to a velocity resolution of 21 to 27 km s-1, depending on the rms. A total of 24 galaxies were detected in CO(1–0) (2 of them marginal), 22 in CO(2–1) (4 of them marginal) and 18 in both transitions. The spectra are shown in Appendix A, where Fig. A.1 displays the spectra detected in CO(1–0) and Fig. A.2 those detected in CO(2–1).

For each spectrum, we integrated the intensity along the velocity interval where emission was detected. For nondetections we set an upper limit as (2)where δV is the channel width and ΔV the total line width. We used a value of ΔV = 300 km s-1 for nondetections in both CO(1–0) and CO(2–1). When the source was detected in only one transition, this line width was used to calculate the upper limit in the other transition.

The results of our CO(1–0) and CO(2–1) observations are displayed in Table 2. The columns are

  • 1.

    Galaxy: galaxy designation;

  • 2.

    ICO(1−0): velocity integrated intensity of the CO(1–0) emission in K km s-1;

  • 3.

    rms: root-mean-square noise of the CO(1–0) spectrum (if available) in mK;

  • 4.

    Ref.: reference of the CO(1–0) data, detailing whether data come from our observations or from the literature (see Sect. 3.1.2);

  • 5.

    Beam: HPBW of the telescope in arcsecond;

  • 6.

    ΔVCO(1−0): line width of the CO(1–0) emission (if detected) in km s-1;

  • 7.

    ICO(2−1): velocity integrated intensity of the CO(2–1) emission (if observed) in K km s-1;

  • 8.

    rms: rms of the CO(2–1) spectrum (if observed) in mK;

  • 9.

    ΔVCO(2−1): line width of the CO(2–1) emission (if detected) in km s-1;

  • 10.

    log(MH2 obs): logarithm of the H2 mass (in solar masses) calculated from the observed central ICO (see Sect. 3.1.3);

  • 11.

    log(MH2): logarithm of the H2 mass (in solar masses) extrapolated to the emission from the total disk (see Sect. 3.1.3).

3.1.2. CO(1–0) data from the literature

We have searched in the literature for available CO(1–0) data for the 20 HCGs of our sample and have compiled data for the velocity-integrated CO(1–0) intensities and line widths (also listed in Table 2) from the following sources:

  • Verdes-Montenegro et al. (1998): 24 galaxies from nine different HCGs. Twenty of them were observed with the NRAO 12 m telescope at Kitt Peak with a beam size of 55′′. The data from the four other galaxies are from Boselli et al. (1996), observed with the SEST 15 m telescope, with a 43′′ beam. Two of these galaxies (68d and 88c) were also observed by us, but we chose the Verdes-Montenegro et al. (1998) data because of their better quality;

  • Leon et al. (1998): 17 galaxies corresponding to ten different HCGs, observed with the IRAM 30 m telescope with a similar setting as in our observations (see Sect. 3.1.1).

There are 16 galaxies that were observed both by us and by Verdes-Montenegro et al. (1998) or Leon et al. (1998). Furthermore, 14 were observed by both Verdes-Montenegro et al. (1998) and Leon et al. (1998). To choose between the different existing spectra (either from our observations or from the literature), we first checked that they were consistent and then applied the following criteria: if available, we chose the spectrum with detected emission. If more than one detected spectrum existed, we chose the one with the lower rms or – in case of comparable rms- the spectrum observed with a larger beam, in order to probe a larger fraction of the disk. Except for the two galaxies mentioned above, we always selected our data due to their better quality in case of duplication. In total, we have CO(1–0) spectra for 86 galaxies (45 from our own observations and 41 from the literature) for our statistical analysis.

3.1.3. Molecular gas mass

We calculate the molecular gas mass, MH2 using the following equation: (3)

where Ω is the area covered by the observations in arcsec2 (i.e. Ω = 1.13 θ2 for a single pointing with a Gaussian beam where θ is the HPBW). This equation assumes a CO-to-H2 conversion factor X = NH2/ICO = 2 × 1020   cm-2 (K km s (e.g. Dickman et al. 1986). No correction factor for the fraction of helium and other heavy metals is included. The molecular gas masses of the AMIGA galaxies are calculated in the same way.

In both the observations that we carried out and the data from the literature, a single position at the center of the galaxy was observed. Because of this and the different beams used by us and others we need to correct for possible emission outside the beam. To extrapolate the observed CO intensities to the total value within the disk we need to know the distribution and extension of the CO emission. Different authors (Nishiyama et al. 2001; Regan et al. 2001; Leroy et al. 2008) have found that the velocity-integrated CO intensity in spiral galaxies follows an exponential distribution as a function of radius with a scale length re: (4)We adopt a scale length of re = 0.2 × r25, where r25 is the major optical 25 mag arcsec-2 isophotal radius, following Lisenfeld et al. (2011), who derive this scale length from studies of the mentioned authors and from their own CO data. We used this distribution to calculate the expected CO emission from the entire disk, taking the galaxy inclination into account (see Lisenfeld et al. 2011, for more details). This approach assumes that the distribution of the molecular gas in galaxies in HCGs is the same as in field spiral galaxies. The implications of this approach are discussed below.

thumbnail Fig. 2

Distribution of the aperture correction factor for MH2.

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The resulting aperture correction factor for MH2 (defined as the ratio between MH2 observed in the central pointing and MH2 extrapolated to the entire disk) is shown in Fig. 2. The ratio between the extrapolated and central intensities is below two for most galaxies (66 out of 86, or 77%), with an average value of 1.78. To check the consistency of the extrapolation, we also performed the analysis presented in this paper for a sample restricted to galaxies with a small (less than a factor 1.6) aperture correction (n = 45), finding no significant differences with respect to the full sample. Thus, we conclude that the aperture correction does not introduce any bias into the results.

The values for the molecular gas mass in the central pointing and the extrapolated molecular gas mass are listed in Table 2. Here, and in the following, we always use the extrapolated molecular gas mass and denote it as MH2 for simplicity. The MH2 distribution is shown in Fig. 3. The average value for spiral galaxies (T ≥ 1) is listed in Table 6. The distribution and average values of MH2, as well as the statistical distributions and average values throughout this work, have been calculated using the Kaplan-Meier estimator implemented in ASURV5, to take the upper limits in the data into account.

thumbnail Fig. 3

MH2 and LFIR distribution of the HCG galaxies. The red filled bins show the distribution of the detected galaxies, the dashed line gives the distribution of upper limits and the full line shows the distribution calculated with ASURV.

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3.2. Far-infrared data

FIR fluxes were obtained from ADDSCAN/SCANPI, a utility provided by the Infrared Processing and Analysis Center (IPAC)6. This is a one-dimensional tool that coadds calibrated survey data of the Infrared Astronomical Satellite (IRAS). It makes use of all scans that passed over a specific position and produces a scan profile along the average scan direction. It is three to five times more sensitive than the IRAS Point Source Catalog (PSC) since it combines all survey data, so it is more suitable for detection of the total flux from slightly extended objects.

We have compiled the FIR data (also derived using ADDSCAN/SCANPI) for 63 galaxies in our sample from Verdes-Montenegro et al. (1998). In the case of the remaining 23 galaxies, we derived FIR fluxes directly using ADDSCAN/SCANPI. To choose the best flux estimator we have followed the guidelines given in the IPAC website7, which are explained in Lisenfeld et al. (2007). As a consistency check, we also applied this procedure to 14 galaxies in the list of Verdes-Montenegro et al. (1998). We found no significant differences, with an average difference of 15% between our reprocessed fluxes and those in Verdes-Montenegro et al. (1998).

From the fluxes at 60 and 100 μm the FIR luminosity, LFIR, is computed as (5)

where FIR is defined as (Helou et al. 1988): (6)The computed LFIR, together with the 60 and 100 μm fluxes compiled from ADDSCAN/SCANPI, are detailed in Table 3. The distribution of LFIR is shown in Fig. 3. The average value of LFIR for spiral galaxies is given in Table 6.

For galaxies HCG 31a and HCG 31c, the FIR fluxes could not be separated. Therefore, we use the sum of both. When comparing LFIR to other magnitudes (LB, MH2, or MHI), we also use the sum of both galaxies.

Table 3

FIR, SFR, SFE, and sSFR.

In order to check the accuracy of the low resolution IRAS data, we compared them to 24 μm data from Spitzer for the 12 groups for which Spitzer data is available. We compared the SFR derived from LFIR (calculated from Eq. (10)) to that derived from the Spitzer 24 μm luminosity, L24   μm, (from Bitsakis et al. 2010, 2011) using the equation SFR(M yr-1) = 8.10 × 10-37(L24   μm(erg   s-1))0.848 (from Calzetti et al. 2010). For most of the objects the agreement was satisfactory: the values of the SFR derived in both ways agreed to better than a factor 2.5 or, in case of IRAS upper limits, the resulting upper limits for the SFR were above those derived from L24   μm. There were only three galaxies in two groups with a larger discrepancy. For HCG 79b and for HCG 37b we obtained a value of the SFR derived from IRAS that was a factor of six higher than the SFR from the 24 μm data and for HCG 37a the difference was a factor of ten. After checking the Spitzer images and IRAS data, we found that in the case of HCG 79 the reason for the discrepancy was the blending of HCG 79a and 79b in the IRAS beam. We thus assumed that the value of LFIR given for HCG 79b in Verdes-Montenegro et al. (1998) arises from both galaxies, so we assigned to each object a fraction of the IRAS fluxes and LFIR such that SFR(IRAS) = SFR(24 μm). A similar situation occurred in the case of HCG 37, where three objects (HCG 37a, HCG 37b and HCG 37c) are blended in the IRAS beam. Here, we assumed that the value of LFIR given for HCG 37b in Verdes-Montenegro et al. (1998) was emitted from all three galaxies and corrected in the same way as for HCG 79.

4. Results

In this section, we aim to study the relation between MH2 and the SFR in HCG galaxies and compare them to isolated galaxies. Furthermore, we search for relations with the atomic gas deficiency of the galaxies and the groups and with the evolutionary phase of the groups. We furthermore investigate the ratio between the two CO transitions, CO(1–0) and CO(2–1).

To search for differences to isolated galaxies, we used two methods: (i) we normalized MH2 and LFIR to the blue luminosity, LB, or the luminosity in the K-band, LK, and compared the ratios to those of isolated galaxies, and (ii) we calculated the deficiency parameters of MH2, LFIR and MHI of the galaxies (see Sect. 4.2). In general, we obtained very consistent results for LB and LK.

We carry out this analysis separately for early-type galaxies and spirals for the following reasons. (i) The morphological distribution is very different for both samples with a much higher fraction of early-type galaxies among HCG galaxies. (ii) The number of early-type galaxies in the AMIGA reference sample is very small so that no statistically significant comparison sample is available. In particular, no deficiency parameter can be derived. (iii) Early-type galaxies tend to have a significantly lower molecular gas content than late-type galaxies, and their FIR emission is not as clearly related to their SFR as it is in late-type galaxies, as a result of the lack of strong SF. Therefore, the use of LFIR as an SF tracer is more questionable.

4.1. Relation between MH2, LFIR, MHI, and LB

Table 4

Correlation analysis of MH2 vs. LB, LFIR vs. LB, and MH2 vs. LFIR.

Figure 4 shows MH2 (top) and LFIR (bottom) versus LB for spirals galaxies (left) and early-type galaxies (right). For spiral galaxies good correlations exist between both MH2, LFIR and LB. A linear fit to the total sample of HGCs is plotted, along with the corresponding fit to the AMIGA sample. The coefficients are listed in Table 4. The fits on LFIR vs. LB for the AMIGA sample are slightly different from the values in Lisenfeld et al. (2007) because we have considered a recent update of the basic properties of the galaxies (e.g. distance and morphological type; see Fernández-Lorenzo 2012, for more details). A slight shift towards higher values in MH2 seems to be present in comparison to the best-fit line of isolated galaxies. The linear regressions between LFIR and LB, MH2 and LB, or MH2 and LFIR (Table 4) show no significative differences between HCGs and isolated galaxies. No clear correlation is visible for early-type galaxies and for log(LB) ≳ 10, the values of both MH2 and LFIR are below those of spiral galaxies.

thumbnail Fig. 4

MH2 vs. LB for a) spiral galaxies (T ≥ 1) and b) elliptical (circles) and S0 galaxies (triangles). LFIR vs. LB for c) spiral galaxies (T ≥ 1) and d) elliptical (circles) and S0 galaxies (triangles). The full green line corresponds to the bisector fit found for HCG galaxies (fit parameters are given in Table 4), while the blue dashed-dotted line corresponds to the bisector fit found for the AMIGA isolated galaxies. Both fits are done for the entire range of morphological types. The dashed black lines are offset by the standard deviation of the correlation for the isolated galaxies, which is  ± 0.35 for the MH2 and  ± 0.4 for LFIR. Black symbols denote detections and red symbols upper limits.

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We note that, in contrast to MH2 and LFIR, MHI shows no correlation with LB (Fig. 5), reflecting that HI is very strongly affected by the interactions and in many galaxies of our evolved groups largely removed from the galaxies.

thumbnail Fig. 5

MHI vs. LB for late-type (T ≥ 1, squares) and early-type (crosses) galaxies. Black symbols denote detections and red symbols upper limits.

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Previous surveys (see e.g. Young & Scoville 1991) have found a linear correlation between MH2 and LFIR. A linear correlation can also be seen in our sample (Fig. 6). We include in this figure the lines for constant LFIR/MH2 values equal to 1, 10, and 100 L/M. Practically all of our galaxies lie in the range of LFIR/MH2 = 1 − 10   L/M, which is typical of normal, quiescent galaxies (Young & Scoville 1991).

thumbnail Fig. 6

MH2 vs. LFIR for late-type (T ≥ 1, squares) and early-type (E+S0, crosses) galaxies. The green line corresponds to the bisector fit found for HCGs galaxies, while the blue dashed-dotted line corresponds to the bisector fit found for the AMIGA isolated galaxies from Lisenfeld et al. (2011). The fits are detailed in Table 4. The dotted black lines correspond to the LFIR/MH2 ratios 1 (left), 10 (middle), and 100 (right) L/M. Black symbols denote detections and red symbols upper limits.

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Finally, we have directly compared E and S0 galaxies in HCGs to galaxies of the same types in the AMIGA sample. In the case of lenticular galaxies we limited the sample of galaxies in HCGs to the same distance range as the AMIGA sample (40–70 Mpc) since for the largest distances the rate of upper limits is very high in HCGs and does not provide any further information. In Fig. 7 (top) we show the relation between MH2 and LB for the S0s in HCGs and from the AMIGA sample. Even though the number of data points is low, a trend seems to be present for S0s in isolated galaxies to have a higher MH2 for the same LB. A similar result is found for LFIR (not shown here), where most lenticular isolated galaxies present higher values than expected for their optical luminosity, while most of the objects in HCGs show upper limits excluding any excess. If S0 galaxies in these dense environments originate from stripping of spirals, this might indicate that molecular gas has also been removed in the process. Although this interpretation is speculative owing to the low-quality statistics, it provides hints for further research in future works.

Concerning the elliptical galaxies, none of the isolated galaxies is detected in CO, while among the four detections in HCGs two have a mass similar to the expected for spiral isolated galaxies (HCG 15d and HCG 79b), and the other two show significantly lower masses (HCG 37a and HCG 93a), pointing to an external origin (Fig. 7, bottom). The FIR luminosity of the Es in HCGs (not shown here) is similar to what is expected for spiral galaxies. It is also noticeable that, while the range of LB values for the S0s in HCGs covers about the same range as for isolated galaxies, Es in HCGs are up to half an order of magnitude more luminous than isolated Es.

thumbnail Fig. 7

MH2 vs. LB for early-type galaxies in HCGs (full symbols) and from the AMIGA sample of isolated galaxies (open symbols) with distances between 20 and 70 Mpc. The lines are the same as in Fig. 4. Black symbols denote detections and red symbols upper limits. Top: S0 galaxies (triangles); bottom: elliptical galaxies (circles).

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4.2. Deficiencies

We have calculated the MH2, LFIR, and MHI deficiencies following the definition of Haynes & Giovanelli (1984) as (7)where we calculated the predicted value of the variable X from LB. Following this definition, a negative deficiency implies an excess with respect to the predicted value.

The expected MH2 for each galaxy is calculated from its LB using the fit to the AMIGA sample in Lisenfeld et al. (2011). The fit, which is given in Table 4, was calculated without distinguishing morphological types. The dominance of spiral galaxies in the AMIGA sample means that the fit is only adequate for spiral galaxies. Because of the low number of early-type galaxies in the AMIGA sample, it is not possible to derive a meaningful deficiency parameter for them. In addition, we calculated the deficiency derived from the relation between MH2 and LK of the AMIGA sample (Lisenfeld et al. 2011), log (MH2) =  −2.27 + 1.05    ×    log  (LK). In a similar way, the expected LFIR is calculated from the fit between LFIR and LB obtained for the AMIGA isolated galaxies (Table 4) for the sample presented in Lisenfeld et al. (2007).

The correlations between MH2 (LFIR) and LB, or LK, have a considerable scatter with standard deviations of 0.35 dex for MH2 and 0.4 dex for LFIR. These standard deviations are much higher than the observational measurement errors. In this case, the error of the mean values are completely dominated by the statistical errors, so we neglect the observational errors in our calculations. The high standard deviation means that individual galaxies with deficiencies up to about these values represent normal deviations from the mean. However, the much smaller error of the mean deficiency allows samples of galaxies (here: galaxies in HCGs and isolated galaxies) to be compared with a much higher precision.

The HI deficiency of the galaxies is calculated following the morphology-dependent fit between MHI and LB in Haynes & Giovanelli (1984). We have considered h = H0/100 = 0.75. We adapted their results, which were based on magzw to our use of with the relation found by Verdes-Montenegro et al. (2005) (magzw = +0.136). Furthermore, considering that we express LB as a function of the solar bolometric luminosity (mag = 4.75), we introduce the correction (8)to express LB in the terms we assume (Sect. 2) to calculate the expected content of HI. The deficiencies in MH2, LFIR, and MHI derived from LB are listed in Table 5.

Table 5

Deficiencies of MH2, LFIR, and MHI derived from LB.

4.2.1. MH2 and LFIR deficiencies

Table 6

Mean values and their errors for spiral galaxies (T ≥ 1) in HCGs and from the AMIGA sample.

The mean MH2 and LFIR deficiencies for spiral galaxies in HCGs are similar, as can be seen in Table 6. Galaxies showing an excess in MH2 or LFIR have values spanning the full range of LB, as can be seen in Fig. 4. Thus, the excess in MH2 or LFIR is not associated with the brightest objects per se. We checked in detail the properties of the nine galaxies showing the largest MH2 excess (HCG 10c, HCG 16a, HCG 16c, HCG 16d, HCG 23b, HCG 23d, HCG 40c, HCG 58a, HCG 88c), and we find that half of them present strong signs of distortion (tidal tails in the optical and/or HI, kinematical perturbations, etc.).

thumbnail Fig. 8

MH2 deficiency vs. LFIR deficiency for late-type (T ≥ 1) galaxies in HCGs (left) and from the AMIGA sample (right). Red symbols represent upper limits in either MH2 or LFIR, and black symbols detections. The y = x line is plotted as reference and does not represent a fit to the data.

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Figure 8 (left) shows Def(MH2) (from LB) vs. Def(LFIR) for each galaxy. Both are strongly correlated, which can be understood coming from the causal relation between the molecular gas and SFR, leading to a lower SFR if the molecular gas as the fuel for SF decreases. For comparison, Fig. 8 (right) displays Def(MH2) of the isolated galaxies versus their Def(LFIR). The behavior of the isolated galaxies does not show a significant difference compared to galaxies in HCGs with a very similar range covered by both samples. However, for the isolated galaxies, Def(MH2) extends to slightly lower values for a given Def(LFIR). This is also reflected in the mean values of Def(MH2) and Def(LFIR) of AMIGA and HCG galaxies (Table 6): (i) values of Def(LFIR) for spiral galaxies are almost the same for both samples, and (ii) Def(MH2) in spirals is higher by 0.15–0.20 for HCG than for AMIGA galaxies (corresponding to a 40–60% higher MH2 than expected for isolated galaxies).

The histograms shown in Fig. 9 underline these findings: whereas the distribution of Def(MH2) for spiral galaxies in HCGs is shifted to negative deficiencies (i.e. an excess) compared to AMIGA galaxies, the distribution of Def(LFIR) is very similar for spiral galaxies in HCGs and in the AMIGA sample. Two sample tests (Gehan’s Generalized Wilcoxon Test and Logrank Test) confirm that the distributions of Def(MH2) are different with a probability of  >96%, whereas the distributions Def(LFIR) are identical with a non-negligible probability.

As an additional test, we compared the ratios MH2/LB and MH2/LK of HCG galaxies to those of isolated galaxies (listed in Table 6). In the case of MH2/LB we derived the ratios both for the entire luminosity range and for low (LB  ≤  1010L) and high (LB  >  1010 L) luminosity galaxies in order not to be affected by the nonlinearity of the MH2-LB relation. In all cases we found a lower ratio (by  ~ 0.2–0.3 dex) for the isolated galaxies, confirming our findings from the deficiency parameter.

thumbnail Fig. 9

Def(MH2) (top) and Def(LFIR) (bottom) distribution of spiral galaxies in AMIGA (black line) and in HCGs (red filled bars), calculated with ASURV in order to take the upper limits into account.

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The higher MH2 for a given LB found for spiral galaxies in HGCs could be explained in three ways: a) a real excess of the total molecular gas mass (and will be further discussed as such in the following section); b) a higher concentration towards the center of the molecular gas in HCG galaxies compared to isolated galaxies, so that the extrapolation of the flux based on a similar extent (see Sect. 3.1.3) would lead to an overestimate of MH2, or c) a systematic difference in the CO-to-H2 conversion factor between the AMIGA and HCG sample. Although we cannot exclude this possibility, we do not consider it very likely. The CO-to-H2 conversion factor is known to depend on a number of galactic properties such as the metallicity, gas temperture, gas density, and velocity dispersion (e.g. Maloney & Black 1988; Narayanan et al. 2011). These properties are likely similar in both samples because of the similar ranges in LB and LFIR (tracing SFR) that they cover. The first two effects (a and b) could both be at work at the same time. In fact, as indicated in e.g. Leon et al. (2008), galaxies in the AMIGA sample are dominated by disk SF while surveys of compact groups (Menon 1995) show that most radio detections involve compact nuclear emission. This can be explained since nuclear emission is thought to be enhanced by interactions that produce a loss of angular momentum of the molecular gas, which subsequently falls towards the center of the galaxy. These dissipative effects are likely near minimum in isolated galaxies. This result was also proposed by Verdes-Montenegro et al. (1998), where the enhanced I25/I100 ratio in HCGs is said to be caused by local starbursts, presumably in the nuclear region. This result is still compatible with the conclusion of a normal level of FIR emission among the HCG galaxies that we find here, if the activity responsible for enhanced 24 μm emission and enhanced/more concentrated molecular gas is localized compared to the overall distribution of gas and dust in the galaxies.

4.2.2. Comparison to the MHI deficiency

thumbnail Fig. 10

MHI vs. MH2 deficiencies (left) and MHI vs. LFIR deficiencies (right) for spiral galaxies (T ≥ 1). The red lines show Def(MHI) = 0, Def(MH2) = 0 and Def(LFIR) = 0, and the dashed lines give Def(MH2) = 0.75, separating low and highly HI-deficient galaxies in our analysis. Red symbols denote upper limits in MH2 or LFIR.

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In Fig. 10 we display Def(MHI) vs. Def(MH2) (left) and Def(LFIR) (right). The mean value of Def(MHI) of the galaxies with available HI data is 0.93 ± 0.13 (12% of the expected value) for spiral galaxies and 1.31 ± 0.11 (5% of the expected value) for all morphological types, which is one order of magnitude more than Def(MH2) and Def(LFIR). We stress here that the samples used in the present paper and in Verdes-Montenegro et al. (2001) are not the same. This earlier study concentrated on the set of data available at that time, which was biased towards HI bright groups. Later, more groups with higher HI deficiencies have been observed with the VLA (Verdes-Montenegro et al. 2007) and are part of the present sample. Therefore, the mean HI deficiency of the galaxies in Verdes-Montenegro et al. (2001) (25% of the expected value for spiral galaxies) is less than the mean HI deficiency of the present sample. We have checked that the HI deficiencies calculated in this paper are consistent with the values for the groups in common with Verdes-Montenegro et al. (2001).

Table 7

Mean values and their errors of deficiencies and ratios of MH2 and LFIR for different samples, considering only spiral galaxies (T ≥ 1).

Most noticeable in Fig. 10 is that even very HI-deficient galaxies have a fairly normal MH2 or LFIR. There is no clear correlation between Def(MHI) and Def(MH2) or Def(LFIR). There might be a weak trend in the sense that a higher MHI deficiency leads to higher MH2 and LFIR deficiencies. This trend is also seen when calculating the mean deficiencies and ratios separately for low and highly MHI deficient galaxies, here chosen as galaxies with Def(MHI)  <  0.75 and Def(MHI)  >  0.75 in order to obtain two groups of roughly the same size (Table 7). However, the differences are small and fall below significance when changing the separation to Def(MHI) = 0.50. Thus, the statistics in our sample are not sufficient to firmly conclude whether this trend is real.

4.3. Comparison to the HI content and evolutionary stage of the group

To study the influence of the global HI content of the group on MH2 and SFR of the individual galaxies we have classified the groups as a function of their Def(MHI) as listed in Sect. 2. The average Def(MH2) and Def(LFIR) of the galaxies belonging to these groups are given in Table 7. We find no relation between the Def(MH2) of the galaxies and the global Def(MHI) of the groups, and neither between the Def(LFIR) and the Def(MHI) of the groups.

In a similar way, we calculated the average Def(MH2) and Def(LFIR) of the galaxies belonging to HCGs in different evolutionary states, as defined by (Borthakur et al. 2010, see Sect. 2), which are also given in Table 7. The Def(MH2) of the galaxies increases slightly as the group evolves along the evolutionary sequence. This trend is also visible in the ratios MH2/LB and MH2/LK. For Def(LFIR), there is no clear relation for spiral galaxies with the evolutionary state, and we only find a trend when considering the total sample, most likely due to a changing fraction of ellipticals.

A very pronounced variation with evolutionary phase is shown by the morphological types (Fig. 11). The ratio of elliptical and S0 galaxies increase strongly in groups in phase 3. It has been proposed (e.g. Verdes-Montenegro et al. 2001; Bekki & Couch 2011) that S0 galaxies in HCGs might be stripped spirals.

thumbnail Fig. 11

Morphological type distribution for different evolutionary phases. From top to bottom, the morphological type distribution of galaxies in HCGs in evolutionary phases 1, 2, and 3 are plotted. The filled red bins correspond to the distribution for the groups in each evolutionary state, while the black line bins correspond, for comparison, to galaxies of all phases.

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4.4. Star formation rate, star formation efficiency and specific SFR

We calculate the SFR from LFIR following the prescription of Kennicutt (1998): (9)where LIR refers to the IR luminosity integrated over the entire mid- and far-IR spectrum (10–1000 μm). This expression is based on a Salpeter initial mass function (IMF). We convert it to the Kroupa (2001) IMF by dividing by a factor 1.59 (Leroy et al. 2008). In our analysis we use LFIR (Eq. (6)), which estimates the FIR emission in the wavelength range of 42.5–122.5 μm. We estimate LIR from LFIR using the result of Bell (2003) that on average LIR ~ 2 × LFIR. Taking this into account, we can calculate the SFR from LFIR as (10)

The values of the SFR of the galaxies in our sample are listed in Table 3.

Table 8

Mean log(SFE) for different samples and measurements (only spiral galaxies, T ≥ 1).

We define the SFE as the ratio between the SFR and the molecular gas mass, SFE = SFR/MH2. The SFE of the individual galaxies in our sample are listed in Table 3. Figure  6 shows a good, roughly linear correlation between LFIR and MH2 and indicates that the SFE in our sample is expected to show a fairly narrow range. To calculate the average SFE of our sample, we must consider that ASURV can only handle data showing upper or lower limits, but not both. Thus, we first calculated the average SFE considering only galaxies detected in CO with an upper limit in FIR, together with the ones detected in both bands. Separately, we considered only those detected in FIR but not detected in CO and the ones detected in both bands. The average values are listed in Table 8.

We calculated the average SFE for the AMIGA sample of isolated galaxies, taking MH2 from Lisenfeld et al. (2011) and LFIR from Lisenfeld et al. (2007), for comparison with the SFE in HCG galaxies. The values are listed in Table 8. We furthermore list the SFE derived from a spatially resolved analysis for 30 nearby galaxies from the HERACLES survey (Bigiel et al. 2011). All values are adjusted to our CO-to-H2 conversion factor, Kroupa IMF, and no consideration of helium in the molecular gas mass. Table  8 shows a slightly lower SFE in HCGs than in AMIGA galaxies, in line with the previous results of an excess in MH2  but a normal value of LFIR. In comparison to the galaxies from the HERACLES survey there is no noticeable difference. Thus, overall there are no strong indications that the process of SF occurs in a different manner in the different environments of HCGs.

The specific SFR, sSFR, is defined as the ratio between the SFR and the stellar mass of a galaxy. We calculated the stellar mass from the K band luminosity since the light in this band is dominated by the emission of low-mass stars, which are responsible for the bulk of stellar mass in galaxies. From LK we derived the stellar mass, M, by adopting a mass-to-luminosity ratio of M/ LK, ⊙  = 1.32 (Cole et al. 2001) for the Salpeter IMF, and applying a correction factor of 0.5 (from Bell et al. 2003) to change to the Kroupa (2001) IMF used in this paper. The values for the individual galaxies are listed in Table 3. The average sSFR for spiral galaxies in our sample is log (sSFR) = –10.61  ±  0.07 yr-1.

4.4.1. SFE and sSFR functions of the deficiencies of the galaxies

In Fig. 12 we display the SFE and the sSFR of the spiral galaxies in our sample as a function of their Def(MHI) and Def(MH2). There is no clear trend of the SFE with the gas deficiency of the galaxies, whether atomic or molecular. This is confirmed by the mean values listed in Table 9. This result indicates that SF proceeds with the same efficiency, regardless of whether it occurs in a galaxy with a low or high MHI deficiency.

On the other hand, galaxies with lower Def(MH2) or Def(MHI) tend to have a higher sSFR (see Fig.12, as well as Table 9 for the quantitative trends). In particular, the trend with Def(MHI) is interesting because it suggests that, although the Def(MHI) of a galaxy has no influence on the absolute SFR or SFE, it has a noticeable effect on the SFR per stellar mass.

Table 9

Mean log(sSFR) and log(SFE) as a function of Def(MHI) and Def(MH2) for spiral galaxies (T ≥ 1).

thumbnail Fig. 12

Specific SFR (sSFR) (top) and star formation efficiency (SFE) (bottom), vs. MH2 and MHI deficiencies of spiral galaxies (T ≥ 1) in HCGs. Red symbols denote upper limits in MH2 or LFIR.

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4.5. Line ratio

Figure 13 shows the CO(1–0) versus the CO(2–1) intensity for the galaxies we observed (Sect. 3.1.1). The plotted intensities are not aperture corrected. The mean ratio between both intensities is ICO(2−1)/ICO(1−0) = 1.13 ± 0.11 for the full sample and 1.13 ± 0.12 for spiral galaxies alone. To calculate this mean ratio with ASURV, we have taken galaxies into account with detections in both CO transitions, as well as those detected only in CO(1–0). These values are slightly higher than those found by Leroy et al. (2009) from CO(2–1) and CO(1–0) maps for nearby galaxies from the SINGS sample (ICO(2−1)/ICO(1−0) ~ 0.8) and than those from Braine et al. (1993) who obtained a mean line ratio of ICO(2 − 1)/ICO(1−0) = 0.89 ± 0.06 for a sample of nearby spiral galaxies. Both values are, in contrast to ours, corrected for beam-size effects.

To interpret the ratio of ICO(2−1)/ICO(1−0) one has to consider two main parameters: the source distribution and the opacity. For optically thick, thermalized emission with a point-like distribution we expect a ratio ICO(2−1)/ICO(1−0) = (θCO(1−0)/θCO(2−1))2 = 4 (with ICO in Tmb and θ being the FWHM of the beams). On the other hand, for a uniform source brightness distribution we expect ratios over 1 for optically thin gas, and ratios between about 0.6 and 1 for optically thick gas (with excitation temperatures above 5 K).

Owing to the different beam sizes of CO(1–0) and CO(2–1) in our observations we cannot distinguish these two cases. However, we can conclude that our values are consistent with optically thick, thermalized gas with an extended distribution. Our mean value is slightly higher than the (beam-corrected) values of Leroy et al. (2009) and Braine et al. (1993) which might indicate that the molecular gas is not completely uniform over the CO(1–0) beam, but slightly concentrated towards the center.

thumbnail Fig. 13

log (ICO(2−1)) versus log(ICO(1−0)) for the galaxies observed by us. Spiral galaxies (T ≥ 1) are shown as filled squares and early-type (T ≤ 0) as crosses. Red symbols indicate upper limits in either ICO(2−1) or ICO(1−0), and black symbols are detections. The y = x line is plotted as reference and does not represent a fit to the data.

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5. A possible evolutionary sequence of the molecular gas content and SFR in HCGs

In contrast to the HI content, which can be highly deficient, the mean deficiencies for both MH2 and LFIR are low and close to the values found for isolated galaxies. In the case of MH2 we even find indications for a 40–60% excess compared to isolated galaxies. The difference in deficiency between the atomic and molecular gas is most likely due to the larger extent of the HI gas, which can thus be removed more efficiently from the galaxies, while the molecular gas, which is typically more concentrated in the inner regions, is presumably less affected by the environment. Subsequently, the lower HI mass might cause lower MH2, which in turn leads to lower SFR. It is, however, remarkable that galaxies with a high HI deficiency can still contain a considerable amount of molecular gas and continue to form stars with a normal SFE. This SF is not expected to last very long because once the molecular gas is used up, no HI is available to provide fuel for future SF.

Within this general picture of the relative normality of MH2 and LFIR, we have found a relation between Def(MHI) and the sSFR, and a tentative trend with Def(MH2) and Def(LFIR). Furthermore, there is a trend of Def(MH2) with the evolutionary phase (Table  7), in the sense that galaxies in HCGs belonging to phase 1 have the highest excess in MH2. These trends might suggest that two mechanisms are in play. First, an increasing MHI deficiency can be interpreted within a scenario in which galaxies in HCGs lose part of their HI as a result of mostly tidal stripping during the initial evolutionary phase, as suggested in the evolutionary model of Verdes-Montenegro et al. (2001). On the other hand, in an early evolutionary phase the HI-to-H2 conversion rate might be enhanced as a result of the continuous interactions between galaxies, leading to the enhancement in MH2 that we observe in evolutionary phase 1. This enhancement of MH2 cannot explain the high HI deficiencies observed in most galaxies, in agreement with the conclusions of Rasmussen et al. (2008), but it could partly explain the lack of HI, especially in the galaxies with the lowest HI deficiencies.

Based on our results, we thus suggest the following scenario, which is speculative but compatible with our observations. Galaxies in an HCG start with a normal content in MH2 and MHI; i.e., they have Def(MHI) = 0 and Def(MH2) = 0. Then, during the early evolutionary phase, tidal interactions enhance the conversion from atomic to molecular gas at the same time as they strip the HI from the galaxies, which leads to Def(MHI)  >  0 and Def(MH2)  <  0. Finally, the multiple interactions within the group strip the main part of the HI in the disks, resulting in Def(MHI)  >  >  0 and, as a consequence also increase in Def(MH2). This last effect could have contributed to an increase in the fraction of lenticular galaxies along the evolutionary sequence due to HI and H2 stripping of spirals.

6. Summary and conclusions

We analyzed data for MH2, which was obtained from observations with the IRAM 30 m telescope and from the literature, LFIR from IRAS, and MHI for 86 galaxies in 20 HCGs in order to study the relation between atomic gas, molecular gas, and SFR, traced by LFIR in these galaxies. We compared these properties to those of isolated galaxies from the AMIGA project (Verdes-Montenegro et al. 2005). We adopted the same CO-to-H2 conversion factor for both samples. The main conclusions of our study can be summarized as follows:

  • The relation between MH2, LFIR, and LB in the galaxies in HCGs is not significantly different from the one found in isolated galaxies. The values of LFIR for spirals galaxies in HCGs are similar to those of the AMIGA galaxies for the same LB. For MH2 we find, however, a slight, but statistically significant, excess (~50%) of HCGs spiral galaxies relative to AMIGA galaxies. This could alternatively be explained by a higher radial concentration of the molecular gas in HCG galaxies to the center when compared with isolated galaxies, so that the extrapolation of the flux based on a similar extent (see Sect. 3.1.3) would lead to overestimating MH2 for the group galaxies. Another possible explanation could be a systematically lower CO-to-H2 conversion factor for spirals in HCGs.

  • For elliptical and S0 galaxies, the large number of upper limits do not allow strong conclusions about their MH2 or LFIR. We note, however that, while the LB range for S0s is comparable to isolated S0 galaxies, Es in HCGs are up to half an order magnitude more luminous in LB than isolated Es are.

  • Practically all of our galaxies lie in the range of LFIR/MH2 = 1–10 L/M, which is typical of normal, quiescent galaxies. The deficiencies in MH2 and LFIR are tightly correlated and cover about the same range as in isolated galaxies.

  • The MHI deficiency, calculated from the VLA data for individual galaxies, is much greater than the other deficiencies with a mean value of 0.93  ±  0.13 (12% of the expected value) for spiral galaxies and 1.31  ±  0.11 (5% the expected value) for all morphological types, and it represents the largest difference with respect to isolated galaxies. These values are significantly higher than those obtained in Verdes-Montenegro et al. (2001) since the sample in that study was biased towards HI bright galaxies, while here we present a redshift-selected sample.

  • The SFE of the spiral galaxies in the HCGs is slightly lower than in isolated galaxies, but in the range of values found for nearby spiral galaxies (Bigiel et al. 2011). We have found no relation of the SFE with either Def(MHI) or Def(MH2).

  • There is a trend for the sSFR to increase with decreasing Def(MHI) and Def(MH2). This suggests that, although the Def(MHI) of a galaxy only has a weak influence on the absolute SFR, it has a stronger influence on the SFR per stellar mass.

  • There is a trend toward decreasing molecular gas deficiency with evolutionary phase, with galaxies in groups in an early phase showing an excess in MH2. This excess goes to 0 in later phases. A similar trend might exist with Def(MHI), but is statistically only marginally significant in our sample. This is interpreted as an initial enhancement of the conversion from atomic to molecular gas due to ongoing tidal interactions, later followed by stripping of most of their HI. In these later phases, evolution of spiral to lenticular galaxies would explain both the overabundance of those morphological types and the MHI deficiency and decrease in MH2 content of the galaxies.

  • No trend in the global HI deficiency of the groups is found, which suggests that the molecular gas content and SF are more driven by one-to-one interaction than directly by the local environment.


1

While HCG 40 was classified in Verdes-Montenegro et al. (2001) as Phase 3, new VLA observations (Yun et al., in prep.) show that a significant amount of HI was missed due to a narrow spectral window, and based on these data it has been reclassified as Phase 2.

3

IRAM is supported by CNRS/INSU (France), the MPG (Germany) and the IGN (Spain).

5

Astronomical Survival Analysis (ASURV) Rev. 1.1 (Lavalley et al. 1992) is a generalized statistical package that implements the methods presented by Feigelson & Nelson (1985).

Acknowledgments

This work has been supported by the research projects AYA2008-06181-C02 and AYA2007-67625-C02-02 from the Spanish Ministerio de Ciencia y Educación and the Junta de Andalucía (Spain) grants P08-FQM-4205, FQM-0108 and TIC-114. D.E. was supported by a Marie Curie International Fellowship within the 6th European Community Framework Program (MOIF-CT-2006-40298). U.L. warmly thanks IPAC (Caltech) for their hospitality during a sabbatical stay when this work was finished. We also thank T. Bitsakis and V. Charmandaris for letting us use their Spitzer data prior to publication, and the anonymous referee for critical comments that helped to put our conclusions on a firmer ground and improved the quality of the paper. This work is based on observations with the Instituto de Radioastronomia Milimétrica IRAM 30 m and the Five College Radio Astronomy (FCRAO) 14 m. The FCRAO is supported by NSF grant AST 0838222. This research has made use of the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. We also acknowledge the usage of the HyperLeda database (http://leda.univ-lyon1.fr).

References

Appendix A: CO spectra

Figure A.1 shows the CO(1–0) spectra of the detections and tentative detections observed by us and Fig. A.1 the CO(2−1) spectra.

thumbnail Fig. A.1

CO(1–0) spectra for the detected HCG galaxies. The detection window is shown with a red horizontal line. Main beam temperature (Tmb, in K) is displayed on the Y axis, and the velocity with respect to LSR in km s-1 on the X axis. Velocity resolution is smoothed to 21 or 27 km s-1. The optical velocity of the galaxy, converted to the radio definition, is marked with an arrow.

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thumbnail Fig. A.2

CO(2–1) spectra for the detected HCG galaxies. The detection window is shown with a red horizontal line. Main beam temperature (Tmb, in K) is displayed on the Y axis, and the velocity with respect to LSR in km s-1 on the X axis. Velocity resolution is smoothed to 21 or 27 km s-1. The optical velocity of the galaxy, converted to the radio definition, is marked with an arrow.

Open with DEXTER

All Tables

Table 1

Basic parameters of the galaxies in the HCG sample.

Table 2

Observed and derived molecular gas properties.

Table 3

FIR, SFR, SFE, and sSFR.

Table 4

Correlation analysis of MH2 vs. LB, LFIR vs. LB, and MH2 vs. LFIR.

Table 5

Deficiencies of MH2, LFIR, and MHI derived from LB.

Table 6

Mean values and their errors for spiral galaxies (T ≥ 1) in HCGs and from the AMIGA sample.

Table 7

Mean values and their errors of deficiencies and ratios of MH2 and LFIR for different samples, considering only spiral galaxies (T ≥ 1).

Table 8

Mean log(SFE) for different samples and measurements (only spiral galaxies, T ≥ 1).

Table 9

Mean log(sSFR) and log(SFE) as a function of Def(MHI) and Def(MH2) for spiral galaxies (T ≥ 1).

All Figures

thumbnail Fig. 1

Morphological types, T(RC3), for the AMIGA isolated (black line) and HCGs (red filled bars) galaxies.

Open with DEXTER
In the text
thumbnail Fig. 2

Distribution of the aperture correction factor for MH2.

Open with DEXTER
In the text
thumbnail Fig. 3

MH2 and LFIR distribution of the HCG galaxies. The red filled bins show the distribution of the detected galaxies, the dashed line gives the distribution of upper limits and the full line shows the distribution calculated with ASURV.

Open with DEXTER
In the text
thumbnail Fig. 4

MH2 vs. LB for a) spiral galaxies (T ≥ 1) and b) elliptical (circles) and S0 galaxies (triangles). LFIR vs. LB for c) spiral galaxies (T ≥ 1) and d) elliptical (circles) and S0 galaxies (triangles). The full green line corresponds to the bisector fit found for HCG galaxies (fit parameters are given in Table 4), while the blue dashed-dotted line corresponds to the bisector fit found for the AMIGA isolated galaxies. Both fits are done for the entire range of morphological types. The dashed black lines are offset by the standard deviation of the correlation for the isolated galaxies, which is  ± 0.35 for the MH2 and  ± 0.4 for LFIR. Black symbols denote detections and red symbols upper limits.

Open with DEXTER
In the text
thumbnail Fig. 5

MHI vs. LB for late-type (T ≥ 1, squares) and early-type (crosses) galaxies. Black symbols denote detections and red symbols upper limits.

Open with DEXTER
In the text
thumbnail Fig. 6

MH2 vs. LFIR for late-type (T ≥ 1, squares) and early-type (E+S0, crosses) galaxies. The green line corresponds to the bisector fit found for HCGs galaxies, while the blue dashed-dotted line corresponds to the bisector fit found for the AMIGA isolated galaxies from Lisenfeld et al. (2011). The fits are detailed in Table 4. The dotted black lines correspond to the LFIR/MH2 ratios 1 (left), 10 (middle), and 100 (right) L/M. Black symbols denote detections and red symbols upper limits.

Open with DEXTER
In the text
thumbnail Fig. 7

MH2 vs. LB for early-type galaxies in HCGs (full symbols) and from the AMIGA sample of isolated galaxies (open symbols) with distances between 20 and 70 Mpc. The lines are the same as in Fig. 4. Black symbols denote detections and red symbols upper limits. Top: S0 galaxies (triangles); bottom: elliptical galaxies (circles).

Open with DEXTER
In the text
thumbnail Fig. 8

MH2 deficiency vs. LFIR deficiency for late-type (T ≥ 1) galaxies in HCGs (left) and from the AMIGA sample (right). Red symbols represent upper limits in either MH2 or LFIR, and black symbols detections. The y = x line is plotted as reference and does not represent a fit to the data.

Open with DEXTER
In the text
thumbnail Fig. 9

Def(MH2) (top) and Def(LFIR) (bottom) distribution of spiral galaxies in AMIGA (black line) and in HCGs (red filled bars), calculated with ASURV in order to take the upper limits into account.

Open with DEXTER
In the text
thumbnail Fig. 10

MHI vs. MH2 deficiencies (left) and MHI vs. LFIR deficiencies (right) for spiral galaxies (T ≥ 1). The red lines show Def(MHI) = 0, Def(MH2) = 0 and Def(LFIR) = 0, and the dashed lines give Def(MH2) = 0.75, separating low and highly HI-deficient galaxies in our analysis. Red symbols denote upper limits in MH2 or LFIR.

Open with DEXTER
In the text
thumbnail Fig. 11

Morphological type distribution for different evolutionary phases. From top to bottom, the morphological type distribution of galaxies in HCGs in evolutionary phases 1, 2, and 3 are plotted. The filled red bins correspond to the distribution for the groups in each evolutionary state, while the black line bins correspond, for comparison, to galaxies of all phases.

Open with DEXTER
In the text
thumbnail Fig. 12

Specific SFR (sSFR) (top) and star formation efficiency (SFE) (bottom), vs. MH2 and MHI deficiencies of spiral galaxies (T ≥ 1) in HCGs. Red symbols denote upper limits in MH2 or LFIR.

Open with DEXTER
In the text
thumbnail Fig. 13

log (ICO(2−1)) versus log(ICO(1−0)) for the galaxies observed by us. Spiral galaxies (T ≥ 1) are shown as filled squares and early-type (T ≤ 0) as crosses. Red symbols indicate upper limits in either ICO(2−1) or ICO(1−0), and black symbols are detections. The y = x line is plotted as reference and does not represent a fit to the data.

Open with DEXTER
In the text
thumbnail Fig. A.1

CO(1–0) spectra for the detected HCG galaxies. The detection window is shown with a red horizontal line. Main beam temperature (Tmb, in K) is displayed on the Y axis, and the velocity with respect to LSR in km s-1 on the X axis. Velocity resolution is smoothed to 21 or 27 km s-1. The optical velocity of the galaxy, converted to the radio definition, is marked with an arrow.

Open with DEXTER
In the text
thumbnail Fig. A.2

CO(2–1) spectra for the detected HCG galaxies. The detection window is shown with a red horizontal line. Main beam temperature (Tmb, in K) is displayed on the Y axis, and the velocity with respect to LSR in km s-1 on the X axis. Velocity resolution is smoothed to 21 or 27 km s-1. The optical velocity of the galaxy, converted to the radio definition, is marked with an arrow.

Open with DEXTER
In the text

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