c ESO 2009
Astronomy & Astrophysics manuscript no. 1637
May 8, 2009
Small-scale systems of galaxies. IV.
arXiv:0905.1264v1 [astro-ph.CO] 8 May 2009
Searching for the faint galaxy population
associated with X-ray detected isolated E+S pairs
R. Grützbauch1,2 , W.W. Zeilinger1 , R. Rampazzo3 , E.V. Held3 , J.W. Sulentic4 , and G. Trinchieri5
1
2
3
4
5
Institut für Astronomie, Universität Wien, Türkenschanzstraße 17, A-1180 Wien, Austria
University of Nottingham, School of Physics and Astronomy, Nottingham NG7 2RD, UK
INAF - Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5,I-35122, Padova, Italy
Department of Physics and Astronomy, University of Alabama, Tuscaloosa, AL 35487, USA
INAF - Osservatorio Astronomico di Brera, Via Brera 28, I-20121, Milano, Italy
Preprint online version: May 8, 2009
ABSTRACT
Context. In hierarchical evolutionary scenarios, isolated, physical pairs may represent an intermediate phase, or “way station”, between collapsing groups and isolated elliptical (E) galaxies (or fossil groups).
Aims. We started a comprehensive study of a sample of galaxy pairs composed of a giant E and a spiral (S) with the aim of investigating their formation/evolutionary history from observed optical and X-ray properties.
Methods. We present VLT-VIMOS observations designed to identify faint galaxies associated with the E+S systems from candidate
lists generated using photometric criteria on WFI images covering an area of ∼ 0.2 h−1
100 Mpc radius around the pairs.
Results. We found two and ten new members likely to be associated with the X-ray bright systems RR 143 and RR 242, respectively.
The X-ray faint RR 210 and RR 216, which were only partially covered by the VIMOS observations, have two and three new faint
members, respectively. The new members increase the number of associated galaxies to 4, 7, 6, and 16 for RR 143, RR 210, RR 216,
and RR 242, respectively, down to MR ∼ −12 + 5 log h100 . We derive structural properties of the faint members from surface photometry. The faint galaxy population of all the systems is dominated by disk galaxies, 40% being S0s with generally low bulge to
total light ratios. A small fraction of the faint companions show signatures of interaction. A remarkable shell system is detected in the
early-type galaxy RR 242 24532. We also derive dynamical properties and optical luminosity functions for the 4 groups.
Conclusions. The above results are discussed in the context of the evolution of poor galaxy group associations. A comparison between the Optical Luminosity Functions (OLFs) of our E+S systems and a sample of X-ray bright poor groups suggest that the OLF
of X-ray detected poor galaxy systems is not universal. The OLF of our X-ray bright systems suggests that they are more dynamically
evolved than our X-ray faint sample and some X-ray bright groups in the literature. However, we suggest that the X-ray faint E+S
pairs represent a phase in the dynamical evolution of some X-ray bright poor galaxy groups. The recent or ongoing interaction in
which the E member of the X-ray faint pairs is involved could have decreased the luminosity of any surrounding X-ray emitting gas.
Key words. galaxies: evolution – galaxies: individual: RR 143 (NGC 2305/2307), RR 210 (NGC 4105/4106), RR 216
(IC 3290/NGC 4373), RR 242 (NGC 5090/5091)
1. Introduction
The X–ray signature of a hot Intra-Group Medium (IGM) has
been detected in loose groups containing an early-type galaxy
population (e.g., Mulchaey, 2000). The pioneering work of
Zabludoff (1999) suggested that groups might fall into different classes defined by their X–ray properties: from groups
with a luminous, extended, hot IGM centred on a giant E
to groups with little or no diffuse emission. Several examples of these classes can now be found in the literature (see
e.g., Mulchaey et al., 2003; Trinchieri et al., 2003; Belsole et al.,
2003; Ota et al., 2004). In a hierarchical evolutionary scenario,
the final product of a merged group would be a luminous isolated
E with an extended X-ray halo, and a few have indeed been identified (see e.g., Mulchaey & Zabludoff, 1999; Vikhlinin et al.,
1999; Jones et al., 2003; Khosroshahi et al., 2004). Jones et al.
(2003) estimated the incidence of fossil groups. They found that
fossil systems, defined as a spatially extended X-ray source with
an X-ray luminosity from diffuse, hot gas of LX,bol ≥ 1042 h−2
50
erg s−1 , represent 8-20% of all systems of the same X-ray lumi-
nosity. However, an optical study of a sample of 100+ isolated
early-type galaxies found that almost no systems were luminous
enough to have been the product of a merger between galaxies brighter than L∗ , i.e., a merged group (Sulentic et al., 2006).
Chandra and XMM-Newton observations of optically selected
merger remnants show that their hot gas is X-ray underluminous compared with mature E galaxies into which these merger
remnants are expected to evolve (see e.g., Sansom et al., 2000;
Nolan et al., 2004; Rampazzo et al., 2006). Brassington et al.
(2007), investigating the evolution of X-ray emission during
the merger process, similarly found that just after an accretion
episode (∼1 Gyr after coalescence) merger remnants are X-ray
faint compared to a typical mature E galaxy. They suggested that
these systems will start to resemble typical elliptical galaxies at
a greater dynamical age (after ∼3 Gyr), supporting the idea that
halo regeneration takes place within low LX merger remnants.
Compact galaxy groups generally show modest diffuse Xray emission (e.g., Trinchieri et al., 2005). However, optically
selected structures (such as compact groups) generally tend to be
2
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
X-ray underluminous in comparison to X-ray selected systems
(Popesso et al., 2007; Rasmussen et al., 2006) and consequently
the modest diffuse X-ray emission is not necessarily associated
with a recent galaxy merger. Rasmussen et al. (2006) found that
low level IGM emission could be an indication that the group is
in the process of collapsing for the first time. Other possibilities
include either that the gravitational potentials are too shallow for
the gas to emit substantially in X-rays or that there is simply little
of no intra-group gas present in those groups.
We have extended the optical and X-ray studies to isolated
physical pairs of galaxies which are simple, and rather common galaxy aggregates in low density environments (LDEs).
Among pairs, the mixed E+S binary systems are particularly
interesting in the context of an evolutionary accretion scenario
(see e.g., Rampazzo & Sulentic, 1992; Hernandez-Toledo et al.,
2001; Domingue et al., 2003, 2005), since the luminous E components might be merger products. The study of such relatively
simple structures may then shed light on a possible evolutionary link between poor groups and isolated Es. One of the most
spectacular examples involves the optical and X-ray bright isolated E+S pair (CPG127=Arp114=NGC2276+2300) with only
two known dwarf companions (Davis et al., 1996).
In this context, we initiated an optical and X-ray study of
four E+S systems: RR 143, RR 210, RR 216, and RR 242
(Grützbauch et al., 2007, hereafter Paper III). In contrast to their
similar optical characteristics, their X-ray properties (see also
Trinchieri & Rampazzo, 2001) indicate that their X-ray luminosities, LX /LB ratios and morphologies are very different, which
implies that they have different origins and/or represent different
evolutionary stages of the systems. X-ray emission in RR 143
and RR 242 is centred on the elliptical but is much more extended than the optical light. The emission in RR 143, although
centred on the elliptical, shows an asymmetric elongation towards the late-type companion. The total extension is r ∼ 500′′
(90 h−1
100 kpc). The extended emission from RR 242 is more regular and has an extent of 700′′ (120 h−1
100 kpc). RR 210 and RR 216
show relatively faint and compact (i.e., within the optical galaxy)
X-ray emission, consistent with an origin in an evolved stellar
population. The emission in RR 143 and RR 242 can be argued
to be related to a group potential (as in CPG127) rather than to an
individual galaxy. In such a scenario, RR 210 and RR 216 could
represent the active part of very poor and loose evolving groups
(see e.g., Rampazzo et al., 2006). The activity is reflected by the
optically extended and distorted morphologies (see Paper III).
This paper presents results of VLT-VIMOS observations of
the faint galaxy populations around the above four RR E+S
pairs. Candidates were selected based on their magnitude, (V−R)
colour, and size in Paper III. The new observations allow us
to determine the redshifts of these faint galaxies and consequently their membership in the E+S systems. These measurements enable us to discuss these groups in the context of previous work by Zabludoff & Mulchaey (1998). They found a significantly higher number of faint galaxies (∼ 20 - 50 members with
MB ≤ −14 + 5 log h100 ) in groups with a significant hot IGM
compared to groups without this component. We estimate optical luminosity functions (OLF hereafter) for the combined X-ray
bright and X-ray faint groups, respectively, and evaluate them in
the context of group evolution (see e.g., Zabludoff & Mulchaey,
2000). We also compare the characteristics of the galaxy population in the E+S pairs’ environments with those of other X-ray
detected groups (see e.g., Tran et al., 2001).
The paper is organised as follows. Section 2 describes the
VLT-VIMOS observations as well as data reduction techniques.
Results are presented in section 3. A discussion of the results in
the light of current literature is given in section 4.
2. Observations and reduction
The colour selection applied to WIde Field Imager (WFI) images
described in Paper III permitted us to isolate a sample of faint
galaxies possibly associated with the E+S systems. This sample
is referred to as the candidate sample in the following. We summarize briefly the selection criteria used to construct this sample.
Galaxy colours were obtained with SExtractor (Bertin et al.,
1996). The source extraction was completed for the R-band and
V-band images simultaneously, using the same extraction criteria for both bands. The MAG AUTO output magnitudes from
SExtractor were then calibrated using the photometric equations given in Paper III (Section 4.2). The colour selection was
based on the colour-magnitude relation of the Virgo Cluster
(Visvanathan & Sandage, 1977), from which the expected location of the red sequence at the pairs’ distance was computed.
The bright member galaxies in the four groups indeed follow
this red sequence or are a bit bluer (see Fig. 12 in Paper III).
Fukugita et al. (1995) found from synthetic galaxy colours that
the K-correction for a typical elliptical galaxy at z ∼ 0.2 corresponds to a shift in colour of ∆(V − R) ∼ 0.2 mag. Adopting an
intrinsic colour of bright ellipticals of (V − R) ∼ 0.7, galaxies
with a colour of (V − R) > 0.9 are already most likely to be in
the background. Therefore, for the sake of simplicity, all objects
with (V − R) ≥ 1 were excluded from the candidate sample. We
further applied a general magnitude limit of mR = 21, finter than
which the star-galaxy classifier of SExtractor becomes unreliable, and a size cut-off at a detected semi-major axis of a = 1.5′′
(see below). This corresponds to MR ∼ −12 + 5 log h100 and
a ∼ 400 pc at the distance of the farthest pair (RR 242).
The candidates are found all over the WFI images (34′ ×
′
34 ), although in some fields they show peculiar, e.g., clumpy,
distributions (see Fig. 14 in Paper III).
To cover the entire WFI field of view and obtain spectra with
a signal-to-noise ratio adequate for measuring the redshifts of
our faint magnitude candidates, we used VIMOS (VIsible MultiObject Spectrograph) (Le Fèvre et al., 2003) at the Very Large
Telescope (VLT) of ESO located at Cerro Paranal, Chile. The
instrument is mounted on the Nasmyth focus B of UT3 Melipal
and has four identical arms, which correspond to the 4 quadrants
covering the entire field, each having a field of view of 7′ × 8′ .
The gap between each quadrant is ∼ 2′ .
Spectroscopic observations were tailored to derive both the
candidate-galaxy redshift and, in a subsequent study, the absorption line–strength indices of member galaxies to investigate their
average age and metallicity and infer their star-formation history
(see e.g., Grützbauch et al., 2005). We adopted the HR (high resolution) blue grism, which permits the coverage of the spectral
region containing the Hβ, Mg2, and Fe (λ 5270 Å, λ 5335 Å) absorption lines with a resolution of R=2050 (1′′ slit) and a dispersion of 0.51 Å pixel−1 . Spectrophotometric and Lick standardstars were either observed or extracted from the VIMOS data
archive with the same instrument set-up. This configuration allows only one slit in the dispersion direction, i.e., each single
spectrum covers the full length of the detector. The wavelength
interval depends on the slit position. At the CCD centre, the
wavelength interval is 4150 – 6200 Å. At the upper CCD edge
(+4′ ), the interval is 4800 – 6900 Å and 3650 – 5650 Å at the
lower edge (−4′ ).
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
3
Table 1. Log of VIMOS observations
Field
OB ID
RR 143-a
RR 143-a
RR 143-b
RR 143-b
RR 143-c
RR 143-c
RR 143-d
RR 143-d
Moon dist.1
[◦ ]
90.7
90.8
85.2
85.2
92.1
92.1
85.1
85.1
Lunar illum.2
1.298
1.308
1.297
1.306
1.359
1.408
1.319
1.304
Seeing
′′
[ ]
0.64
0.65
0.84
0.78
0.85
0.69
0.83
0.82
1800
1800
1800
1800
1.018
1.006
1.073
1.135
0.66
0.66
1.03
1.05
150.3
150.1
121.6
121.9
0.041
0.042
0.138
0.136
2006-03-29
2006-03-29
1800
1800
1.052
1.081
1.03
1.03
143.1
143.3
0.000
0.000
2006-03-24
2006-03-24
2006-03-09
2006-03-09
2006-03-25
2006-03-26
2006-03-26
2006-03-26
1800
1800
1800
1800
1800
1800
1800
1800
1.104
1.150
1.093
1.175
1.208
1.069
1.132
1.191
0.82
0.94
0.82
0.73
1.40
1.17
1.06
1.09
76.3
76.6
110.0
109.5
89.2
100.5
101.9
102.1
0.328
0.326
0.730
0.734
0.221
0.141
0.132
0.131
Exposure time
[sec]
1800
1800
1800
1800
1800
1800
1800
1800
Airmass
218955
218955
218973
218973
218964
218964
218946
218946
Observation date
JJJJ-MM-DD
2006-01-29
2006-01-29
2006-01-23
2006-01-23
2006-01-30
2006-01-30
2006-01-22
2006-01-22
RR 210-a
RR 210-a
RR 210-b
RR 210-b
219249
219249
219258
219258
2006-03-31
2006-03-31
2006-03-26
2006-03-26
RR 216-a
RR 216-a
219285
219285
RR 242-a
RR 242-a
RR 242-b
RR 242-b
RR 242-c
RR 242-c
RR 242-d
RR 242-d
219759
219759
219750
219750
219741
219741
219732
219732
1
2
0.005
0.004
0.446
0.444
0.006
0.006
0.550
0.547
Angular distance of the moon on the sky.
Fractional illumination of the moon.
Each WFI field is covered by four VIMOS observing blocks,
one VIMOS pointing for each quadrant of the WFI field of view.
The observations of each single block were divided into two exposures. Bias, flat-field, and standard-star calibration files were
associated with each observing block as well as the helium-argon
lamp spectrum for wavelength calibration. Observations were
obtained in service mode to guarantee optimal observing conditions. Unfortunately, the complete four quadrants were obtained
only for RR 143 and RR 242. Two quadrants were observed for
RR 210 and only one for RR 216. Table 1 provides the observing log for the four E+S systems. Figure 1 shows the WFI fields
with the results of our VIMOS observations.
The basic CCD reduction of each frame containing spectra as
well as the wavelength calibration was completed using the ESO
software environment Esorex. The 2D spectrum of each slit was
coadded to the corresponding one from the second exposure. The
object(s) in each slit was(were) then extracted into 1D objectspectra containing the total light of each target. Finally, each
wavelength-calibrated spectrum was stored as a single FITS-file
for further processing.
Redshifts were measured using the cross-correlation technique (see e.g., Tonry & Davies, 1979). To provide a reliable
estimate of radial velocities and their uncertainties, 5 stellar template spectra were used. The IRAF1 package rvsao provides the
xcsao task to measure radial velocities via cross-correlation.
During this interactive cross-correlation procedure, the result of
the cross-correlation was always directly inspected to avoid spurious results. In some cases, usually for emission-line dominated
spectra, the lines could not be identified by xcsao and were processed using the IRAF task splot. With this task a Gaussian fit
Table 2. Observation statistics of the candidate samples.
1
IRAF is distributed by the National Optical Astronomy
Observatories, which are operated by the Association of Universities
for Research in Astronomy, Inc., under cooperative agreement with the
National Science Foundation.
Table 2 reports the statistics of the observational campaign.
For the newly confirmed members of the E+S systems, we obtained accurate surface photometry from our WFI images. The
methods adopted are fully explained in Paper III.
Candidate sample
Observed
Measured z total
Measured z xcsao
Measured z splot
Members
Members corr.
RR 143
171
106 (62%)
84 (79%)
69
15
2 (2%)
4.4
RR 210
190
55 (29%)
52 (95%)
50
2
2 (4%)
4.9
RR 216
178
24 (13%)
22 (92%)
22
–
3 (14%)
15
Total number of spectroscopically observed objects:
Total number of measured redshifts:
RR 242
118
73 (62%)
55 (75%)
45
10
10 (18%)
15.4
258
213 (83%)
Row 1: Number of galaxies in the candidate sample.
Row 2: Number of galaxies observed with VIMOS, in brackets the percentage
of spectroscopically observed galaxies out of the candidate sample.
Row 3: Number of galaxies for which a redshift could be measured, in brackets
the percentage of galaxies with measured redshift out of all spectroscopically
observed galaxies.
Row 4: Number of redshifts obtained by cross-correlation with xcsao.
Row 5: Number of redshifts obtained by fitting single lines with splot.
Row 6: Number of newly discovered members, in brackets the percentage of
newly found member galaxies out of all galaxies with measured redshifts.
Row 7: Completeness corrected number of members (see Section 3).
to each spectral line can be performed. The adopted redshift is
then the average of all fitted lines, and its error is given by the
standard deviation in the different redshift values.
4
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
Fig. 1. Results of VIMOS observations superposed on the WFI images: RR 143 (top left), RR 210 (top right), RR 216 (bottom
left), and RR 242 (bottom right). Each field is centred on the bright E galaxy of the pair. Due to an error in the service-mode
observations, the field of RR 143 was not dithered and thus shows the gaps between the single CCDs. Each WFI field was covered by
4 VIMOS pointings, one for each quadrant of the WFI field of view. However, not all four quadrants were observed for each group,
due to incomplete service-mode observations. Marked with different symbols are: new group members (green circles), confirmed
background galaxies (red squares), spectroscopically observed candidates too faint for redshift measurement (blue triangles), and
spectroscopically non-observed objects (black diamonds). The newly identified group members are also labelled with their object
ID (see Table 3). The different spectroscopic coverage of each field can be seen, since the not-covered quadrants only contain black
diamonds.
3. Results
Redshift measurements allowed us to identify faint galaxies
likely to be physically associated with each E+S pair system.
However, the spectroscopic coverage of the fields around the
four pairs is not uniform, due to incomplete service-mode observations. In spite of two approved observing programs, only
∼ 2/3 of the total area (the 4 WFI fields) was observed with
VIMOS. Additionally, the incompleteness differed between the
4 fields with RR 143 and RR 242 being covered completely,
while for RR 210 and RR 216 only 50% and 25%, respectively,
of their fields were covered.
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
5
Fig. 2. Histogram of galaxy redshifts measured in the WFI fields. Large windows show the velocity range from 0 − 30 000 km s−1 ,
i.e. up to ∼ 10 times the group systemic velocity. The whole range of redshifts (up to ∼ 130 000 km s−1 ) is shown in small windows
in the upper right corner of each panel. Galaxies considered to be group members (new and previously known, see text) are shaded
in red. The bin size is 500 km s−1 .
Figure 2 shows the distribution of measured radial velocities
for each field. Each group panel shows the velocity range up to
30 000 km s−1 (∼10 times the group velocity), whereas the full
range of velocities found up to ∼ 130 000 km s−1 is plotted in
the small window in the upper right corner of each group panel.
Galaxies belonging to each group are shaded in red. To show the
group structure in greater detail,we included the already known
group members inside the WFI field of view in each histogram.
RR 143 (top left panel) and RR 210 (top right panel) show
prominent concentrations at ∼ 40 000 km s−1 , while RR 242
(bottom right panel) shows a weaker concentration at ∼ 16 000
km s−1 . Galaxies in the field of RR 216 (bottom left panel) seem
to show a less clustered distribution in redshift space. However,
only one quadrant of the WFI field was observed with VIMOS
in the latter case, resulting in a smaller number of galaxies with
measured redshifts than in the other fields (see Table 2).
3.1. Identification of faint members and completeness
correction
Different membership criteria are considered in the literature. A
galaxy is often considered to be a member of a structure if the velocity difference between the galaxy and the structure’s systemic
velocity is lower than a certain value. For instance, Karachentsev
(1990) and Hickson et al. (1992) adopted |(vgalaxy − vgroup )| ≤
1000 km s−1 , while Ramella et al. (1994) used ≤ 1500 km s−1 .
Group membership may also be defined in terms of the group
velocity dispersion, σgroup , leading to a selection that is more capable of being adapted to the true group’s properties. In this case,
a limit of |(vgalaxy − vgroup )| ≤ 3σgroup has been used, reflecting
the approximate dynamical boundaries of the group (see e.g.,
Mulchaey, 2000; Cellone & Buzzoni, 2005; Firth et al., 2006).
Different membership criteria applied to our sample yield the
same result: there are no galaxies close to the velocity limits
set by the above criteria (the group velocity dispersions used in
the flexible group membership criterion are listed in Table 5).
Line 6 of Table 2 gives the number of newly identified members for each of our groups. Only a small fraction of the candidates turned out to be new members of our four E+S systems. However, the E+S systems are clearly defined structures in
redshift space suggesting that they are real, albeit sparse, physical associations. Coordinates, redshifts with estimated uncertainty, projected distance from the bright E, total R-band magnitude, and (V − R) colour of the new members are presented
in Table 3. The structural parameters (Sersic index n, effective
radius re , and central and effective surface brightness µ0 and
µe ) given in the rightmost 4 columns of Table 3 are obtained
6
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
Table 3. Summary of the properties of the new members of the E+S systems.
ID
3.34
3.87
D2
[Mpc]
0.122
0.269
mR
[mag]
17.25
16.95
(V − R)
[mag]
0.56
0.34
1927 ± 183
2008 ± 45
3.08
9.08
0.117
0.080
17.59
14.26
-39 33 45.6
-39 34 47.1
-39 44 36.4
3716 ± 134
3193 ± 50
2697 ± 62
2.13
9.63
5.32
0.165
0.209
0.054
-43 32 55.6
-43 36 58.1
-43 38 53.2
-43 42 27.5
-43 44 20.2
-43 39 43.0
-43 45 43.2
-43 46 28.0
-43 43 21.6
-43 55 38.4
3441 ± 60
3172 ± 45
3731 ± 81
3045 ± 52
3237 ± 70
3400 ± 88
2697 ± 45
2655 ± 84
3017 ± 87
2997 ± 40
2.58
3.42
7.20
6.63
9.63
0.00
10.05
4.26
6.80
3.19
0.202
0.080
0.050
0.120
0.054
0.058
0.050
0.239
0.030
0.245
-64 10 55.1
-64 29 15.8
vrad
[km s−1 ]
3296 ± 46
3724 ± 73
12 07 43.28
12 06 50.26
-29 43 10.9
-29 36 23.5
RR216 03519
RR216 040524
RR216 12209
12 25 17.80
12 26 11.30
12 25 37.04
RR242
RR242
RR242
RR242
RR242
RR242
RR242
RR242
RR242
RR242
13 22 10.44
13 21 03.88
13 21 16.54
13 20 26.71
13 20 55.54
13 21 33.01
13 21 14.82
13 19 44.25
13 21 22.55
13 20 21.80
α
δ
RR143 09192
RR143 24246
06 47 53.25
06 47 07.47
RR210 11372
RR210 134933
1
2
3
4
5
6
08064
13326
15689
20075
223275
23187
243526
25575
28727
36267
R1
n
1.05
1.57
re
[kpc]
1.30
1.40
µ0
µe
[mag arcsec−2 ]
21.30
23.23
19.73
22.80
0.50
0.52
0.83
1.98
0.97
1.30
22.24
18.03
23.68
21.97
17.26
14.57
15.58
0.51
0.60
0.59
1.02
1.07
0.74
1.22
2.52
2.17
21.41
20.18
21.57
23.28
22.15
22.82
17.85
16.15
17.73
16.33
15.06
17.37
13.71
16.42
17.05
18.67
0.64
0.60
0.74
0.55
0.69
0.44
0.62
0.60
0.64
0.50
1.21
1.42
1.38
1.51
1.34
0.75
2.35
2.24
1.16
0.89
1.32
2.11
0.73
1.38
0.98
0.95
1.68
1.73
0.69
1.07
21.88
20.15
20.47
20.02
18.38
20.54
16.00
19.02
20.08
22.32
24.15
22.87
23.11
22.95
20.94
21.82
20.76
23.53
22.25
23.90
confidence parameter R from the cross-correlation procedure. If R = 0 then the lines were measured by hand with splot.
projected distance from the E member of the pair.
RR210 13493 = 2MASX J12065029-2936236.
RR216 04052 = 2MASX J12261133-3934474.
RR242 22327 = ESO 270- G 001.
RR242 24352 = ESO 270- G 003.
from one-component Sersic-model fits completed with GALFIT
(Peng et al., 2002, see Paper III).
Figure 3 shows the distribution of R-band magnitudes for the
candidate list (Paper III), for all objects observed spectroscopically (red), all galaxies with measured redshifts (magenta), and
for those adopted as group members (black). Both RR 143 and
RR 242, with complete VIMOS pointings, show a reasonable
degree of completeness. In contrast, for both RR 210 (2 pointings) and especially RR 216 (1 pointing) the number of candidates without spectroscopy is very high (see also Table 2). In
addition to the missing pointings in RR 210 and RR 216, with
50% and 25% coverage, respectively, we must also consider
two additional sources of incompleteness affecting the number
of spectroscopically observed candidates. The major source of
incompleteness is caused by instrumental constraints: (a) the
gaps between the 4 CCDs in the VIMOS field of view reduce
the analysed area by about 22%, and (b) the HR blue grism
allows only one slit to be placed along the dispersion direction, i.e., galaxies with similar declinations (closer than the slit
length) cannot be observed with a single pointing. The second
type of incompleteness depends on the source magnitude, i.e.,
redshifts of fainter objects become increasingly difficult to measure with the adopted exposure times. The magnitude-dependent
incompleteness starts at R ∼ 18.5 mag, while at brighter magnitudes the incompleteness is determined by the instrumental constraints. Therefore, a magnitude-limited completeness correction
was adopted.
Our photometric selection criteria (see Section 2 ) may
also bias the number of faint member galaxies. When selecting
the candidate objects, our first goal was to complete reliably
the surface photometry to check if their structural properties
are consistent with them being faint galaxies associated with
their respective pair. To obtain a reliable estimate of the surface photometric parameters, the galaxy size (or the effective
radius) should exceed the size of the seeing disk, which has
60
RR 143
RR 210
60
40
40
20
20
0
50
0
RR 216
RR 242
30
40
30
20
20
10
10
0
0
15
20
15
20
Fig. 3. Distribution of R-band magnitudes in all 4 group samples: RR 143 (top left), RR 210 (top right), RR216 (bottom left)
and RR 242 (bottom right). The different histograms show the
full candidate sample (unshaded), spectroscopically observed
galaxies (red, mildly shaded), galaxies with measured redshift
(magenta, shaded), and new member galaxies (black, heavilyshaded). They correspond to the samples given in lines 1, 2, 5,
and 6 in Table 2. The bin size is 1 magnitude.
a FWHM < 1′′ in all images (see Table 4 in Paper III). The
detection size a given by SExtractor is not directly related
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
RR 143
RR 210
10
10
5
5
0
RR143
RR210
RR216
RR242
0
RR 216
RR 242
10
10
5
5
0
0
-22
7
-20
-18
-16
-14
-12
-10
-22
-20
-18
-16
-14
-12
-10
Fig. 4. Completeness correction of the number of member galaxies in each bin of absolute magnitude MR . The bin size is 1 magnitude. The true number of identified members per magnitude
bin is shaded in red, whereas the number of expected members
is plotted in white. The numbers over each bin represent the percentage of candidates that have a measured z in each magnitude
bin, i.e., the spectroscopic completeness in that bin. If there is no
number above a bin, then no candidates are located in this bin.
to the galaxy’s effective radius or the FWHM, but it gives the
semi-major axis length of the detection ellipse, which is most
likely larger than the galaxy’s effective radius (which contains
only half of the light). To ensure that we selected galaxies with
effective radii larger than the area affected by the seeing, we
chose a generous detection size limit of 1.5′′ , which corresponds
to a physical size of ∼ 400 pc at the distance of RR 242. This
limit is also reasonable in a physical sense, since it is smaller
than the smallest Local Group dwarf galaxies found by Mateo
(1998), and reaches the domain of dwarf spheroidal galaxies
(dSph), which are typically smaller than 500 pc. Since small
galaxies tend to be faint, the size limit leads to incompleteness
of the candidate sample at faint magnitudes. Plotting detected
size against magnitude, we found that this incompleteness
starts at mR ∼ 18. This is also visible in Fig. 3, where the
magnitude histogram of the candidate sample bends over at
around this magnitude. This incompleteness, however, does not
affect galaxies at brighter magnitudes, since we did not exclude
any objects brighter than mR ∼ 18 via this size cut-off. This
apparent magnitude corresponds to an absolute magnitude of
MR ∼ −15 + 5 log h100 , which is the limit to which we compare
our results with the literature in the discussion (also because
the spectroscopic completeness is higher than 50% for brighter
magnitudes, see below). Therefore, we can conclude that our
size cut-off does not affect the results we present here.
We adopt a simple completeness-correction criterion based
on the assumption that the same fraction of members is present
in the sample of spectroscopically observed and not observed
objects. The fraction of confirmed members in the sample with
Fig. 5. Colour-magnitude relation for the candidate samples in
the fields of the 4 E+S systems. Members are plotted as red triangles, previously known group members are additionally labelled
with their ID (see Fig. 6-9 in Paper III). Spectroscopically observed candidates that were found to be in the background are
plotted as green squares, while candidates without measured z
are plotted as blue crosses. The solid line is a fit to the redsequence of the Virgo-Cluster (Visvanathan & Sandage, 1977),
whereas the dashed vertical line represents the colour cut-off applied to the candidate sample.
measured redshifts is computed for each magnitude bin and
multiplied by the number of galaxies in the candidate sample
found in the respective bin. This yields the total completenesscorrected number of members in each magnitude bin. The sum
over all magnitude bins then gives the corrected number of
group members given in line 7 of Table 2. Figure 4 shows the
completeness-corrected number of members in each magnitude
bin (white histogram) with the true number of confirmed members indicated in red. The number printed above each bin is the
percentage of the spectroscopic completeness in the respective
bin. So “100” means that the completeness is 100% and all candidates in this bin have a measured redshift. If there is no number
above a bin then this bin is empty, i.e., there are no candidates
within this magnitude range.
Poor statistics in the RR 216 field makes the application
of the above criterion very uncertain. We note, e.g., the high
number of expected members in the −15 > MR > −16 magnitude bin: the only candidate with measured redshift in this
bin was confirmed as a member. Following the above criterion,
all galaxies in the candidate sample at this magnitude were
assumed to be members. However, comparing the number of
confirmed members with the number of galaxies with measured
redshifts (Table 2) also suggests a higher number of member
galaxies for RR 216, approaching a number similar to that of
the X-ray bright RR 242.
Incompleteness effects are certainly an issue in our sample,
although they are not expected to play a major role for abso-
8
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
Table 4. Results of the Bulge-disk decomposition.
Object
Morph. type
type
B/TS
∆ B/T
∆χ2ν
0.16
0.29
0.77
0.29
-0.6
-0.01
30.6
0.9
RR143 09192
RR143 24246
SB0
S0
RR210 11372
RR210 13493
dE
SB0
...
15.11
17.58
14.87
...
16.18
...
17.30
17.58
14.53
0.00
0.44
0.00
0.23
0
0.22
0.36
10.6
RR216 03519
RR216 04052
RR216 12209
dE
SB0
Sc
21.95
17.95
22.07
17.29
14.67
15.63
20.69
17.39
21.79
...
...
...
17.32
14.70
15.63
0.01
0.05
0.00
0.04
0.08
0.00
-0.03
-0.03
0
0.1
1.5
0.1
RR242
RR242
RR242
RR242
RR242
RR242
RR242
RR242
RR242
RR242
dE
S0
dE
dE
S
SB0
SB0
SB?
dE
dE
19.25
17.17
19.52
17.39
...
16.72
14.59
16.27
18.25
21.31
18.18
16.54
18.10
16.79
...
15.51
14.52
17.70
17.38
17.92
19.28
18.64
20.93
18.74
17.63
17.58
15.31
17.62
17.95
23.06
...
...
...
...
...
16.56
17.36
...
...
...
18.30
16.29
17.95
16.53
17.96
15.79
14.17
16.77
17.40
17.89
0.27
0.36
0.21
0.37
...
0.25
0.49
0.79
0.31
0.04
0.29
0.10
0.06
0.12
0.57
0.41
0.29
0.31
0.38
0.01
-0.02
0.26
0.15
0.25
...
-0.16
0.2
0.47
-0.07
0.03
0.3
0.1
0.3
0.8
...
2.2
21.4
0.8
-82.9
22.5
08064
13326
15689
20075
23187
22327
24352
25575
28727
36267
bulge
mR
17.31
18.47
Sersic bulge
bar
disk
mR
mR
20.70
18.55
...
17.51
B/TdeV
de Vaucouleurs bulge
bulge
disk
mR
mR
19.10
17.33
18.36
17.37
Notes: Bulge and disk magnitudes and bulge-to-total light (B/T) ratios of the fit with de Vaucouleurs bulges (columns 3-4)
and Sersic bulges (columns 5-7) respectively. Columns 8-10 give the Bulge-to-Total light ratios of the two models respectively
and their difference. The last column gives the difference between the χ2ν of the de Vaucouleurs and the Sersic-model fit.
lute magnitudes as faint as MR ∼ −17 + 5 log h100 . Due to instrumental constraints, our spectroscopy missed 3 “bright” candidates in RR 143, visible as the 2 first bins labelled with “0”
in Fig. 4. In all the other groups, the candidates brighter than
MR = −17 + 5 log h100 without measured redshift are accounted
for in the completeness correction: 1 object in RR 210 (bin labelled with “50”, 1 object added by correction), 2 objects in
RR 216 (2 bins labelled with “50”, 2 objects added by correction) and 1 object in RR 242 (bin labelled with “67”, 1 object added by correction). Figure 4 also shows that, apart from
RR 216 (where only one quarter of the field was covered), the
spectroscopic completeness is above 50% in all magnitude bins
down to MR = −15+5 log h100 . Any information fainter than this
magnitude is not used in the comparison of our results with the
literature and does not affect our conclusions.
3.2. The photometric and structural properties of faint
members
In general, we find that the confirmed companions tend to be of
intermediate luminosity, which is a domain populated by faint
S0, spirals, and dwarf elliptical galaxies.
Figure 5 shows the colour-magnitude relation of each group.
Confirmed member galaxies are indicated by red triangles. The
solid line represents a fit to the red sequence of the Virgo Cluster
(Visvanathan & Sandage, 1977) shifted to the pairs’ distance,
while the dashed line shows the colour cut-off applied to the
candidate sample. Galaxies follow the red sequence for earlytype galaxies even at faint magnitudes. The new members have
very uniform colours with no blue star-forming dwarfs found
in our sample. This might be partially caused by the selection
criteria used to construct the candidate sample, especially by
the size cut-off. As discussed in section 3.1, this size cut-off
leads to the loss of galaxies below mR ∼ 18 (corresponding to
MR ∼ −15 + 5 log h100 at the farthest pair’s distance). However,
also above that magnitude, where the candidate samples are
complete (photometrically), they do not contain blue galaxies.
Blue galaxies fainter than mR ∼ 18, are not abundant in our sample, but those observed spectroscopically were found to be in the
background.
In the following, we consider only galaxies identified as group members according to the redshift measurements. Figure 6 shows R-band images of new confirmed members. Galaxy morphologies were investigated with
ELLIPSE and GALFIT (see Paper III for a full explanation).
Results are compared with the galaxy morphologies of the
Zabludoff & Mulchaey (2000) X-ray detected groups discussed
in Tran et al. (2001).
In Fig. 7, we display residual images after subtraction of a galaxy model constructed from the isophotal fit with ELLIPSE. Different galaxy substructures are
clearly visible in the residual images including: asymmetries (RR143 09192, RR242 23187, RR242 25575), bars
(RR143 09192, RR216 04052, RR242 22327), filaments
(RR242 13326), and shells (RR242 24352). The system of
shells in RR242 24352 extends to a radius of ∼ 6 h−1
100 kpc
and is not aligned (∆PA∼30◦ ) with the semi-major axis of the
galaxy. Different formation scenarios for stellar shells have been
proposed including weak interactions, accretion of companions
and major/minor mergers (see e.g., Dupraz & Combes, 1986;
Hernquist & Quinn, 1987a,b). In any case, they are considered
clear evidence of environmental influence on galaxy evolution,
which is apparently found even in rather sparse groups such as
those in our sample.
Radial profiles of surface brightness (µ), ellipticity (ǫ), position angle (PA), and higher order coefficients of the Fourier
expansion a3 , b3 , a4 , and b4 are shown for each galaxy in Figs. 8
and 9. Isophotal profiles and the GALFIT two-component models (see next paragraph) are shown. Dotted, red lines are Sersic
(bulge) and exponential (disk) models, while solid lines represent the resultant total galaxy model. The figures show that
the two-component models yield good fits to the real profiles
in most cases. There are some residuals for RR242 22327 and
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
9
Fig. 6. R-band thumbnails of member galaxies of all 4 groups. From left to right: 1 st row from the top: RR143 09192,
RR143 24246, RR210 11372, and RR210 13493; 2nd row: RR216 03519, RR216 04052, RR216 12209, and RR242 08064; 3rd
row: RR242 13326, RR242 15689, RR242 20075, and RR242 22327; 4th row: RR242 23187, RR242 24352, RR242 25575, and
RR242 28727; last row: RR242 36267. The scale bar has a length of 2 arcseconds in each image.
RR242 24352, which exhibit very complex structures and are
not well represented even by 3 components. The deviations from
the model in the outskirts of RR242 22327 are caused by an underestimated ellipticity of the outer isophotes. The ellipticity had
to be fixed because of the high percentage of masked area due
to some brighter foreground stars in the field. This was also the
case for RR143 09192: a foreground star close to the galaxy cen-
tre led to the loss of a considerable part of the galaxy image. This
galaxy is not included in the B/T analysis in Fig. 10.
The morphological classification listed in Table 4 was completed using the surface photometric profiles in Figs. 8 and 9.
The presence of bulge and disk components is clearly evident as
a bend in the surface-brightness profile at the transition between
the two components, due to the different shapes of the radial
surface-brightness profiles of spheroids and disks. The surface
10
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
Fig. 7. Residual images after subtraction of a model constructed from the isophotal fit with ELLIPSE. Objects are shown in the same
order and with the same scale as in Fig. 6.
brightness µ declines with radius as µ ∼ r1/n . Disks (and dwarf
ellipticals) have an exponential profile with n ∼ 1, whereas
spheroids are characterised by a higher value of n. Additionally,
a disk is characterised by a constant ǫ and PA, whereas along
the bulge, both ǫ and PA can vary. In this way, ellipticals (being
pure bulges) and lenticular galaxies (having a bulge and a disk)
can be distinguished easily. Bars can also be identified in the surface photometric profiles, showing a constant surface brightness
and PA in combination with a high ǫ. We find a number of bars
in our faint galaxy sample. In 5/17 objects a bar is clearly dis-
tinguishable in the radial profiles. A weak bar is also suspected
in (RR242 25575). Bars can also be an indication of on-going
interaction as shown by Noguchi (1987).
In order to investigate the bulge-to-total light (B/T) ratios a bulge-disk decomposition was attempted with GALFIT
(Peng et al., 2002). Tran et al. (2001) found that the galaxy population of poor X-ray detected groups was well described by
a 2-component model composed of a de Vaucouleurs bulge
(µr ∼ r1/4 ) and an exponential disk (µr ∼ r). However, the shape
of the profile determined by the exponent 1/n is supposed to
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
11
20
22
22
22
24
24
24
26
26
26
28
0.35
0.4
0.3
0.3
0.25
0.2
0.2
-30
0.1
0.45
0.4
0.35
0.3
80
-40
70
78
-50
60
76
-60
0
-0.1
0
-0.1
74
0.02
50
0.1
0
0
0
0
0.05
0
0
-0.05
0
0.1
0
-0.02
0
0
5
10
15
0
0
5
10
15
0
5
10
15
20
18
20
22
20
22
22
24
24
24
26
26
26
28
28
0.25
0.4
0.1
0.3
0.2
0.2
0.05
-30
0.15
70
-40
60
0.1
180
160
-50
50
-60
140
0
0
0
-0.05
0
0
-0.05
0
0.05
0.05
0
0
0
0
0
0
-0.05
0
10
20
30
0
5
10
15
0
10
20
30
40
20
22
22
22
24
24
24
26
26
26
28
28
28
0.3
30
0.3
0.2
0.25
0.5
0.4
0.2
0.1
0.3
0.15
100
-45
-70
-50
50
-55
-75
-60
0.05
0.5
0
0
0
-0.05
0.1
-0.2
0.2
0
0
0
0
0
0
0.2
0
0
0
0
10
20
30
-0.05
0
5
10
15
20
25
0
10
20
30
Fig. 8. Surface photometry. The radial profiles of surface brightness, ellipticity, position angle, and higher order coefficients of the
Fourier expansion a3 , b3 , a4 , b4 of faint member galaxies in the R-band obtained by ELLIPSE. Each successive isophote is plotted
as a small cross, the error bars are the errors given by ELLIPSE. The surface-brightness panel (top panel in each plot) also shows
the profile of the 2-component model obtained from GALFIT. The bulge and disk components are plotted as red dotted lines. If a
third component (bar) improved the fit, it is also plotted. The total 2 (or 3) component galaxy model is then plotted as solid black
line. Objects shown are RR143 09192, RR143 24246, RR210 11372, RR210 13493, RR216 03519, RR216 04052, RR216 12209,
RR242 08064, and RR242 13326.
12
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
18
20
20
22
22
22
24
24
24
26
26
26
28
28
30
0.15
0.06
0.2
0.04
0.1
0.1
0.02
0.05
-55
60
40
20
0.2
50
-56
0
-57
-50
-58
0
0
0.2
-0.1
0
0.2
0.1
0.05
0
0.05
0
0
0
0
-0.2
0
-0.1
0
0
0
-0.1
-0.1
2
4
6
8
10
12
0
5
10
15
20
20
0
10
20
20
18
22
22
20
24
22
24
26
24
26
28
26
0.3
0.6
0.3
0.2
0.25
0.5
0.1
0.2
0.15
20
-52
-70
-54
0
-75
-56
-58
-80
-20
-85
0.1
-40
0
0
0
0
0
-0.5
0.1
0
-0.2
0
0
-0.2
0.1
-0.2
0.1
0
0
0
5
10
15
20
0
0
0
10
20
20
30
40
0
5
10
15
20
22
22
24
24
26
26
28
28
0.03
0.08
0.02
0.06
0.01
0.04
20
50
0
0
-20
-50
-40
0.05
0
0
0.05
0
0
0
0
-0.1
0
0
2
4
6
8
10
12
0
5
10
Fig. 9. Surface photometry as in Fig. 8. Objects shown are RR242 15689, RR242 20075, RR242 22327, RR242 23187,
RR242 24352, RR242 25575, RR242 28727, and RR242 36267.
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
13
20
22327
24352
23187
13493
22
28727
4052
RR143a
24246
12209
13326
20075
9192
3519
25575
11372
36267
8064
RR210a
15689
24
RR216b
RR242a
26
0
1
2
Fig. 11. Hamabe-Kormendy relation (Hamabe & Kormendy,
1987, dashed line) for new group member galaxies. The pair
ellipticals are also plotted for comparison. All galaxies are labelled with their ID. The area left of the vertical dotted line
represents the region inhabited by “ordinary” galaxies (see
Capaccioli et al., 1992).
Fig. 10. Bulge-to-total light ratios (B/T) of member galaxies of
the 4 groups. Black and open symbols indicate B/T obtained
from fitting a Sersic bulge, red and solid symbols indicate the
de Vaucouleurs bulge. Top: histogram of B/T ratios for the fit
with Sersic bulges (solid black line) and de Vaucouleurs bulges
(red dotted line). Centre: B/T versus radius from the pair elliptical. Lines are least squares fits to the data (black solid – Sersic:
red dashed – de Vaucouleurs). Bottom: B/T versus local projected density. Values are computed with the closest 5 galaxies
(d5 , circles) and the closest 3 galaxies (d3 , triangles). Lines are
least squares fits to d5 (dashed-dotted – Sersic; red dashed – de
Vaucouleurs) and d3 (long-dashed – Sersic; red short-dashed –
de Vaucouleurs). The vertical lines in the upper right corner represent the rms scatter of each fit.
vary significantly with mass. This is valid not only for low-mass
ellipticals but also for the bulges of low-mass galaxies, which
are expected to show a different profile shape, i.e., a lower n.
Both bulge/disk combinations with n = 4 (de Vaucouleurs bulge)
and n as a free parameter (Sersic bulge) were fit to our member galaxies to investigate differences between these two models. Bulge and disk magnitudes as well as B/T ratios for the two
different models can be found in Table 4. Differences in B/T can
be significant especially for bulge-dominated galaxies. The difference between the χ2ν of the two-component fits from GALFIT
is given in the last column of Table 4. We find that the Sersic
bulge provides an accurate representation of faint galaxy bulges:
apart from one galaxy (RR242 28727), the Sersic fit always has
a lower χ2ν than the de Vaucouleurs fit. The de Vaucouleurs fits of
RR143 09192, RR210 13493, RR242 24352 and RR242 36267
are unsatisfactory, while RR242 23187 was fitted badly by the
Sersic model.
Figures 8 and 9 show the results of the Sersic-model fit and
the observed surface-brightness profile in the R-band. The dotted, red lines represent the bulge and disk model, respectively. A
model for the bar was added for the brighter bars. The solid line
represents the resulting galaxy model.
Figure 10 (top panel) shows the distribution of B/T ratios
for Sersic (black, shaded) and de Vaucouleurs models (red, dotted line). Both distributions show that our sample is dominated
by low B/T - i.e., disk-dominated - galaxies. An automatic classification by the B/T ratio was proposed by Marleau & Simard
(1998) and also used by Tran et al. (2001). This automatic classification divides late (S, disk-dominated) and early-type (S0 and
E, bulge-dominated) galaxies at B/T = 0.4. This may work in
distinguishing between bright E and S galaxies, but is problematic for faint S0 galaxies. A comparison between visual (based
on surface photometric profiles) and automatic classification
(based on the B/T ratio) shows the problem: visually classified
S0 galaxies show a wide range of B/T ratios and are not necessarily bulge-dominated systems. We found a high fraction of
S0s from our visual classification (7 out of 17, 40%), but only
3 of those galaxies have a B/T ≥ 0.4 (18%). This would yield
an early-type fraction similar to that of the field, while the visually estimated S0-fraction is more typical of galaxy clusters and
X-ray luminous galaxy groups (Tran et al., 2001, and references
therein).
Figure 10 (middle panel) plots B/T ratios versus the projected distance from the dominant E member. The individual
galaxies are plotted as small dots, while the mean and dispersion are computed in bins of 30 h−1
100 kpc and plotted as open
circles with respective error bars (red – de Vaucouleurs bulge;
black – Sersic bulge). The lines are a least squares fit to the individual data points (solid black line – Sersic; dashed, red line
14
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
– de Vaucouleurs). A morphology-radius relation appears to be
present in our sample: galaxies with higher B/T ratios tend to
be more centrally concentrated than pure disks. This result was
also found for the groups studied by Tran et al. (2001). However,
the relation is very flat and the scatter is large, probably due to
projection effects and the small number of galaxies studied. We
computed Spearman rank coefficients for the two data sets resulting in ρ sersic = 0.14 and ρdeV = 0.32, values that are not
considered to be statistically significant.
The bottom panel of Fig. 10 shows the dependence of B/T
on the local projected number density. This local density was
computed using the distance to the 5th (d5 ) and 3rd (d3 ) closest
group member. The results are plotted as circles (d5 ) and triangles (d3 ), and in black and red for Sersic and de Vaucouleurs
bulges, respectively. The lines are the least squares fit to the data
of d5 (dashed-dotted – Sersic; red dashed – de Vaucouleurs) and
d3 (long-dashed – Sersic; red short-dashed – de Vaucouleurs).
Using only the area occupied by the 3 closest galaxies changes
the results slightly but significantly: the correlation for d5 is
stronger than for d3 , the Spearman rank coefficients being
ρd3 = 0.46 and ρd5 = 0.56. For our sample size (16 galaxies),
the latter value is higher than the critical value of ρ at the 0.05
(2σ) level of significance. Hence, there seems to be a positive
correlation between the projected number density (d5 ) and the
morphology (expressed by the B/T ratio).
The Hamabe-Kormendy Relation (Hamabe & Kormendy,
1987, HKR hereafter) in the log re − µe plane is one projection of the Fundamental Plane for early-type galaxies. Faint Es,
S0s, or dwarf ellipticals (dEs) do not follow this relation but are
distributed in the re − µe plane below the HKR and also below
re = 3 kpc. They are considered a distinct family of galaxies,
the so-called ordinary galaxies, whereas galaxies above re = 3
kpc belong to the family of bright ellipticals (Capaccioli et al.,
1992). The ordinary family are often considered to be the building blocks of galaxies of the bright family, although more recent
simulations of Evstigneeva et al. (2004) show that by merging,
the galaxies evolve along tracks that are parallel to the HKR.
The re - µe plane for our four E+S groups is shown in Fig. 11
with the dashed line indicating the HKR (in the R-band). The
dotted line at log re = 0.5 separates the bright and ordinary families (Capaccioli et al., 1992). The four bright Es lie on the HKR
clearly in the bright galaxy domain, while the newly identified
faint members are distributed in the log re − µe plane below the
limit of ordinary galaxies. The brightest galaxy of the newly discovered galaxy population RR242 24352 is closest to the HKR,
while the faintest galaxies are off the relation in the expected
vertical strip of ordinary galaxies (Capaccioli et al., 1992).
3.3. The spectral properties of faint members
VIMOS spectra of the new member galaxies are presented in
Figs. 12 and 13. As already suggested by the colour-magnitude
relation, the spectra are characterised by a relatively old stellar population. Strong metal lines (Mg I and Fe) are present in
most of the spectra, although many also exhibit strong Hβ absorption suggesting the presence of a younger or intermediateage population. Emission lines are detected in only two galaxies;
RR242 23187, where Hγ and Hδ emission suggests recent or ongoing star-formation activity, and RR143 24246, where [O II] λ
3727-29 Å emission is detected. The forbidden lines of [O III] λ
4959 Å and λ 5007 Å as well as Hβ emission is present in both
galaxies, although on top of a substantial absorption compo-
nent. The detailed analysis of the line-strength indices (see e.g.,
Rampazzo et al., 2005; Grützbauch et al., 2005; Annibali et al.,
2007) will be treated in a forthcoming paper.
4. Discussion
In the following discussion, we consider only the spectroscopically confirmed member galaxies. Their basic properties are
summarised in Table 3. The new members are used, along with
the previously known group members (see Paper III), to infer group properties. Table A.1 lists all known members of
the groups: i.e., new members found within the WFI field of
view and already known members found in the NASA/IPAC
Extragalactic Database (NED). To investigate the effects of the
WFI’s small field of view, we compare velocity and magnitude
distributions of the full sample (90′ ) and the WFI subsample
(34′ ×34′ ).
Group kinematics and dynamics are discussed in the first
section using a luminosity-weighted approach for the determination of all mass-related quantities (see e.g., Firth et al., 2006).
However, we also compute uniformly weighted quantities and
discuss the results of the different weighting. The dynamical
formulae are given in the Appendix. Distribution of members,
group compactness, crossing times, and mass-to-light ratios are
analysed and compared with the literature. OLFs of the individual and combined group samples are presented in the next
section. The OLFs of our X-ray faint and X-ray bright groups
are compared with OLFs for samples of: 1) X-ray detected poor
groups, 2) simulated and observed fossil groups, and 3) the OLF
of the local field. We then attempt to investigate the possible relation between the group dynamical characteristics as well as the
group “activity” and the X-ray luminosity of the E members.
4.1. The E+S system kinematics and dynamics
Distributions of radial velocities for the full sample and the
WFI-subsample (red-dashed) are shown in Fig. 14. The vgroup
of each sample is plotted as a vertical, dashed line, while the
horizontal line above each histogram indicates the 3 σr limits
as an approximate dynamical boundary for each group. The
mean velocities of the two samples do not differ significantly,
apart from RR 216, where the velocity of the WFI-subsample
is dominated by the very bright pair elliptical (no other bright
members are present in the WFI-field) and is therefore biased
towards a higher value. For this group, the group velocity for
the full sample is v90arc = 3223 ± 49 versus vWFI = 3378 ± 24
for galaxies in the WFI field of view. The unweighted velocity
dispersions of the two samples are comparable within the errors
for all four systems. The luminosity-weighting significantly
changes the velocity dispersion only in RR 216 where the
velocity difference between the 2 dominating pair galaxies is
very small, biasing the dispersion within the WFI field towards
a low value (σ90arc = 241 ± 35 versus σWFI = 56 ± 18).
The peculiar velocity is the difference between an individual
galaxy velocity and the group centre velocity (vgalaxy − vgroup ).
It is usually normalised by the group’s velocity dispersion for
comparison with other groups (see e.g., Zabludoff & Mulchaey,
2000). Figure 15 shows the peculiar velocities of galaxies versus their projected radial distance from the optical group centre. The group members are separated into giants (big symbols)
and dwarfs (small symbols) following Zabludoff & Mulchaey
(2000). Galaxies with absolute magnitudes brighter than MR =
−19 + 5 log h100 are considered giants, and fainter objects are
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
15
Fig. 12. Rest frame spectra of member galaxies observed with VIMOS. Objects RR143 09192, RR143 24246, RR210 11372,
RR21 13493, RR216 03519, RR216 04052, RR216 12209, RR242 08064, and RR242 13326. The most prominent absorption
and emission lines are marked.
defined as dwarfs. The four group samples are combined in
analysing the dependence of velocity dispersion on the projected
distance. The mean velocity and velocity dispersion is computed
in radial bins of 100 h−1
100 kpc. They are plotted as red squares
(mean velocity) and respective red error bars (velocity dispersion) for each bin.
Figure 15 suggests that the velocity dispersion is not constant with projected radius. The maximal dispersion is reached
at around 0.2 h−1
100 Mpc (∼ the border of the WFI field of view),
from where it starts to decrease out to ∼ 0.5 h−1
100 Mpc. At greater
radii, the dispersion rises again, which could indicate the transition between the potential of the group concentrated around the
E and the influence of the global large-scale density. We computed the statistical errors in the velocity dispersion in each bin
following Osmond & Ponman (2004, their equation 4). The σmin
and σmax quoted below are normalised velocity dispersions obtained by dividing the peculiar galaxy velocities by the respective group’s velocity dispersion. The maximum of σv at 0.15 h−1
100
Mpc is σmax = 1.24 ± 0.20, while the minimum at 0.45 h−1
100 Mpc
is σmin = 0.45 ± 0.14. Hence, the velocity dispersion is not constant within the errors. However, this is a tentative result due to
the intrinsically low number of members in our groups.
To investigate the effect of incompleteness and low number
statistics in our sample, a set of Monte Carlo simulations was
performed. The question is whether the detected drop in velocity
dispersion is significant, or if our measured σ cannot be distinguished from a constant velocity dispersion. Therefore, in each
radial bin a Gaussian velocity distribution of the same σ (the
maximum velocity dispersion found in the second radial bin:
σmax = 1.24) was assumed. Then, a random sample of n velocities was taken from this Gaussian distribution, where n represents the number of galaxies of our sample in the respective
bin. After 1000 iterations, the mean velocity and velocity dispersion and their deviations in each bin was computed. The 1σ
deviations are plotted in Fig. 15 as grey (mean velocity) and red
(velocity dispersion) shaded areas. The result is that the mean
velocity is consistent with the group velocity in all bins, whereas
the velocity dispersion is lower than the expected 1σ deviation
from the constant value in all bins out to 0.5 h−1
100 Mpc. If the velocity dispersion were constant, the error bars would reach into
the red shaded area. This result is significant at > 2σ in the bin
between 0.4 - 0.5 h−1
100 Mpc. We repeated this analysis with different radial binnings (75, 100, and 150 h−1
100 kpc) and the drop in
velocity dispersion was always above the 1σ significance level.
We therefore suggest that the velocity dispersion decreases with
radius and that the kinematics in the region outside 0.5 h−1
100 Mpc
may not trace the group potential (as traced by the IGM). The
dynamical quantities were also calculated by excluding galaxies
lying outside this radius.
Another interesting implication of Fig. 15 is that the position
of the bright E+S pair does not coincide with the optical centre
16
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
Fig. 13. Rest frame Spectra of member galaxies observed with VIMOS. Objects RR242 15689, RR242 20075, RR242 22327,
RR242 23187, RR242 24352, RR242 25575, RR242 28727, and RR242 36267. The most prominent absorption and emission
lines are marked.
of each group. The projected offsets for RR 216, RR 242, and
RR 143 are around 0.1 h−1
100 Mpc, and only RR 210 is located
very close to the projected group centre.
The R-statistic was developed by Zabludoff & Mulchaey
(2000) to facilitate comparison between the distribution of members in both projection and velocity space. It is defined as R2 =
(d/δd )2 + (|v pec |/δ|v pec | )2 , where δd and δ|v pec | denotes the rms deviations of the entire sample in projected distance and peculiar
velocity, respectively. A galaxy with a large distance d from
the group centre or a large peculiar velocity v pec will yield a
high value of R, while an average member should have R ∼ 2.
We compared the distribution of R values for three magnitudelimited subsamples: 1) the brightest group galaxies (BGGs) with
MR < M ∗ , 2) giants with MR < −19 + 5 log h100 , and 3) dwarfs.
We also compared the full sample with galaxies within 0.5 h−1
100
Mpc of the group centre. Figure 15 (bottom panels) shows the
distribution of R for the BGGs (heavily shaded), giants (shaded),
and dwarfs (unshaded) in the full sample (left panel) and for the
galaxies within 0.5 h−1
100 Mpc (right panel). In both samples, the
BGGs are more centrally concentrated than dwarfs and giants
with hRBGG i = 1.3 ± 0.7 and hRBGG,0.5 i = 1.4 ± 0.8. Considering
the full sample, the dwarf population follows a different distribution in phase-space than the giants with hRdwar f s i = 2.0 ± 0.9
and hRgiants i = 2.3 ± 1.1.
A K-S test is used to check whether the R distributions of
the three subsamples differ significantly. It gives the following
probabilities that the 3 distributions are the same: PBGG,DWARF =
0.13, PBGG,GIANT = 0.23, and PDWARF,GIANT = 0.29. Those values are significant above the 1σ level. Inside the 0.5 h−1
100 Mpc
radius, the difference between dwarfs and giants vanishes (both
have hRi ∼ 2.2), but the R distributions of BGGs and dwarfs
have a probability of originating in the same distribution of
PBGG,DWARF = 0.04. This suggests that the 2 distributions are
different above the 2 σ significance level. This is not only caused
by the optical group centre moving closer to the BGGs, since
both, hRBGG,0.5 i and hRDWARF,0.5 i are higher than the respective
values for the full sample.
These findings are consistent with the results of
Zabludoff & Mulchaey (2000), who argued that the three
galaxy populations (BGGs, giants, dwarfs) move on different
orbits and have not yet mixed. From the present data, this also
seems to be the case for the BGGs and dwarfs at the centres of
our 4 groups, although we note that outside the WFI field of
view we lack information about faint member galaxies.
The dynamical properties of the four E+S systems are given
in Table 5, where the first row provides data for the entire sample and the second row for group members inside a radius of 0.5
h−1
100 Mpc. Errors were estimated using Monte Carlo simulations.
For velocity-related quantities (vgroup , velocity dispersion), a
Gaussian distribution was assumed. A set of 1000 groups with
the respective number of members was constructed for both the
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
17
Table 5. Dynamical properties of the E+S systems.
Group
RR 143
Nr. of
members
Distance
(Modulus)
[Mpc (mag)]
Optical group centre
α (2000)
δ (2000)
[h:m:s]
[◦ :′ :′′ ]
vgroup
[km s−1 ]
Velocity
dispersion
[km s−1 ]
3D velocity
dispersion
[km s−1 ]
6
5
6
5
49.6 (33.5)
49.5 (33.5)
51.1 (33.5)
50.8 (33.5)
06:48:05.6
06:47:59.2
06:48:28.9
06:47:22.7
-64:10:54
-64:11:57
-64:02:09
-64:11:57
3321±52
3317±53
3426±72
3404±82
126±39
123±42
180±54
189±66
214
210
304
319
RR 210
Uniform wi
23
23
30.1 (32.4)
31.5 (32.5)
12:06:55.8
12:06:42.8
-29:46:47
-29:44:27
2018±26
2113±34
123±18
165±24
210
273
RR 216
23
12
23
12
48.1 (33.4)
47.6 (33.4)
48.6 (33.4)
47.9 (33.4)
12:26:05.5
12:26:01.5
12:26:17.1
12:25:56.4
-39:38:38
-39:33:50
-39:40:19
-39:37:55
3223±49
3191±67
3255±52
3206±80
241±35
229±48
241±36
273±58
414
393
411
463
28
22
28
22
52.5 (33.6)
52.9 (33.6)
50.1 (33.5)
50.1 (33.5)
13:20:20.7
13:20:54.8
13:20:21.6
13:20:44.1
-43:42:06
-43:42:08
-43:38:07
-43:37:00
3520±60
3546±65
3356±70
3354±76
323±45
314±48
368±51
363±56
537
529
623
617
Group
Harmonic
radius (RH )
[Mpc]
Virial
radius (Rvir )
[Mpc]
Crossing
time
[tc H0 ]
Virial
mass
[1012 M⊙ ]
Projected
mass
[1012 M⊙ ]
Group
luminosity
[1011 L⊙ ]
M⊙ /L⊙
RR 143
0.170±0.001
0.161±0.001
0.43±0.04
0.31±0.03
0.267±0.001
0.252±0.001
0.68±0.06
0.49±0.05
0.14±0.02
0.09±0.02
0.16±0.09
0.09±0.07
1.5±0.7
1.3±0.7
7.6±5.0
6.1±4.8
7.8±6.5
5.7±4.7
30.2±20.6
21.8±10.9
1.06±0.005
1.04±0.004
1.13±0.005
1.10±0.005
13±1
12±1
66±5
55±5
0.053±0.0001
0.44±0.06
0.083±0.0002
0.70±0.10
0.14±0.01
0.17±0.21
0.4±0.2
6.5±4.3
6.0±4.2
16.6±12.0
1.56±0.001
1.71±0.001
2±0.2
38±4
0.35±0.01
0.24±0.01
0.87±0.09
0.45±0.06
0.55±0.01
0.37±0.01
1.36±0.14
0.71±0.10
0.24±0.03
0.10±0.02
0.24±0.19
0.10±0.09
11.1±3.3
6.7±2.8
27.3±10.0
18.3±9.4
63.2±39.9
32.0±22.0
65.4±42.8
37.5±20.4
3.84±0.010
3.14±0.008
3.92±0.010
3.71±0.009
28±3
21±3
69±10
57±9
0.221±0.001
0.141±0.001
0.41±0.02
0.28±0.01
0.346±0.001
0.221±0.001
0.64±0.02
0.44±0.02
0.09±0.01
0.05±0.01
0.10±0.03
0.06±0.02
12.5±2.8
7.6±1.8
29.9±9.1
20.3±6.9
92.3±85.0
31.6±17.7
119.0±97.3
59.3±23.8
2.19±0.005
1.79±0.005
1.99±0.004
1.61±0.004
57±3
42±2
150±9
125±7
Uniform wi
Uniform wi
RR 242
Uniform wi
Uniform wi
RR 210
Uniform wi
RR 216
Uniform wi
RR 242
Uniform wi
Notes: The first two rows give the luminosity weighted results for the whole sample (1 st row) and for galaxies within
nd
0.5 h−1
100 Mpc (2 row). In line 3 and 4 the same values are calculated with uniform weights.
full sample and galaxies within 0.5 h−1
100 Mpc (see Appendix A).
Each of these groups consisted of galaxies with velocities that
were taken randomly from a Gaussian distribution of respective
mean velocity and dispersion given in Table 5. The rms of the
mean velocity and velocity dispersion of this set of groups is
the error given in Table 5. The position and luminosity related
quantities were treated in a different way. Here, the number of
additional galaxies expected from the completeness correction
was added to the existing group. These additional galaxies were
selected randomly from the sample of candidates without a measured redshift. After 1000 iterations, the rms of the values (RH ,
Rvir , r⊥ , group luminosity) of this set of “complete” groups were
computed. The errors in virial mass Mvir , projected mass MP ,
and crossing time tc are a combination of this two methods, since
they are dependent on velocity and position of the objects.
Both luminosity-weighted and uniformly-weighted results
are provided in Table 5. Different weightings influence the results dramatically, and especially in RR 210 where the parameters of the bright pair dominate the resulting values. The pair
components are separated by a projected distance of only 6 h−1
100
kpc and are close to the centres of both mass and velocity. This
leads to an underestimation of Mvir and a very low luminosityweighted value of M/L for this pair. RH and Mvir rise by an order of magnitude if the members are uniformly weighted. Values
change by a factor of ∼ 2 for the other systems. The group ve-
locity and velocity dispersion are less affected by the weighting.
Weights given in Table A.1 illustrate the dominance of the E
member. RR 143a and RR 210a contribute more than 50% to the
total luminous mass, while the spiral companion is clearly the
second brightest object. RR 216b and RR 242a contain about
1/3 of the luminous mass and have a few other massive objects brighter or comparable to the spiral pair member in their
close environment. The unweighted value of RH for RR 216
is very high and indicative of a higher large-scale density in
which this pair is embedded and a probable dynamical link to
the Hydra-Centaurus cluster. However, it also means that the
brighter members are more concentrated towards the pair than
the intermediate-luminosity galaxies.
Differences between Mvir and MP are quite high with MP being generally higher than Mvir by a factor of 2-5. This effect is
expected in systems where individual galaxies are close to each
other in projection. Mvir was found to underestimate the mass
by a factor of ∼ 3 or more in such systems (Heisler et al., 1985).
An alternative reason for this difference could be that the groups
are not virialised. On the other hand, the crossing times of the
groups are a small fraction of a Hubble-time and usually shorter
than 0.2 (apart from RR 216), even when calculated with the entire sample. This indicates that the groups, or at least their centres, are virialised (Fergusson & Sandage, 1990). RR 216 and
RR 242 have comparable Mvir but very different M/L due to
18
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
RR 143
RR 210
10
10
5
5
0
ALL GROUPS
0
1000
2000
3000
1000
4000
RR 216
3000
4000
2000
3000
4000
RR 242
10
10
5
5
0
0
1000
2000
2000
3000
4000
1000
Fig. 14. Distribution of radial velocities in the 4 groups, including previously known and new member galaxies. The
luminosity-weighted mean velocity and 3σ velocity dispersion
are plotted as vertical dashed and horizontal solid line respectively. The WFI-subsample is plotted in red (dashed), the line
indicating the 3σ velocity dispersion of the WFI-subsample is
drawn above the black line indicating the 3σ velocity dispersion
of the whole sample within 90′ from the pair E. The velocity of
the elliptical pair member is indicated by the shorter dotted line.
the very high group luminosity of RR 216, which is twice the
group luminosity of RR 242. Since the number of giant galaxies
is the same in both systems (Ngiant = 10), the individual giants of
RR 216 contain more (luminous) mass than the ones in RR 242.
The M/L ratios are based on Mvir , motivated by the common
use of Mvir in the literature. The M/L ratios of the groups are
consistent with typical values found for poor groups of galaxies
(see e.g., Firth et al., 2006), although at the lower limit, indicating that the virial mass is indeed underestimating the total mass
of the systems (as expected by the definition of Mvir ).
The positions of group members within our E+S systems is
presented in Fig. 16. The groups are moved to a common distance: the field of view shown in the figure corresponds to ∼
2.5 h−1
100 Mpc on each side. We plot the large-scale environment
search radius of 90′ (see Appendix A) centred on the E member (solid line), the WFI field of view (central square), the 0.5
h−1
100 Mpc radius (dashed line), and the mean harmonic radius RH
(dotted line), the latter two centred on the optical group centre.
The central, solid-line circle is RH calculated only for members
inside 0.5 h−1
100 Mpc and centred on the respective group centre.
The plot makes clear that the optical group centres do not coincide with the positions of the pairs except for RR 210. This
group also appears to be the most compact system with small
RH . This may be caused by the higher number of galaxies with
intermediate luminosity identified by NED (due to the proximity of the group). This system also involves the pair with the
smallest projected separation, which also causes the luminosityweighted RH to be smaller. The unweighted RH is comparable
Fig. 15. Analysis of galaxy offsets in projection and velocity.
Upper panel: Peculiar galaxy velocities (vgalaxy − vgroup ) normalised by the line-of-sight velocity dispersion σr as a function of the projected radial distance from the optical group centre. Members of all 4 group are plotted with different symbols:
solid triangles: RR 143; open circles: RR 210; open triangles:
RR 216; solid squares: RR 242. The size of the symbols indicate
dwarf (small symbols) and giant (big symbols) group members.
The galaxies are binned in projected radial distance every 0.1
h−1
100 Mpc. The mean velocity and velocity dispersion in each bin
is plotted as red squares and red error bars. The shaded areas
show the 1σ deviations expected from the Monte Carlo simulation (see text). The mean velocity rms is mildly shaded (in grey),
while the area of the 1σ deviation from a constant velocity dispersion is heavily shaded (in red). Bottom panels: The R-statistic
quantifies the offset in velocity and distance of each galaxy from
the group mean velocity and from the optical group centre (see
text). BGGs (heavily shaded), giants (shaded), and dwarfs (unshaded) show different distributions of R suggesting that they
occupy different orbits.
to that of the other groups (see Table 5 for values). The galaxies
around RR 216 are spread over the full area investigated without any central concentration, which is indicative of the higher
large-scale density of its environment (outskirts of the HydraCentaurus Cluster).
The spatial distribution of faint members can only be investigated for the group centres (within the WFI field of view ∼ 250
h−1
100 kpc radius around the E member). The redshift information
of galaxies outside this field is taken only from NED, which is
highly incomplete at fainter magnitudes. The radial density profile of the faint group members (MR > −19+5 log h100 ) is plotted
in Fig. 17 for all four groups. The number of group members is
counted in radial bins of 20 h−1
100 kpc width and divided by the
area of that bin. The radius is measured from the field centre, i.e.,
the E member. The first bin 0 - 20 kpc is not plotted because no
galaxies are found there because of the large extent of the bright
elliptical. The effective radius of the smallest of the four Es is
∼ 12 kpc. The faint galaxies in RR 242 are clearly concentrated
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
RR 143
0.0015
RR 210
-63
19
-28
-64
E
S
E
S
-30
-65
0.001
-32
-66
106
-38
104
102
100
98
RR 216
184
182
180
178
RR 242
-42
0.0005
-43
-39
S0
S
-40
E
S
-44
-41
-45
0
-42
188
186
184
203
202
201
200
199
198
Fig. 16. Position of member galaxies of the 4 groups moved to
a common distance. The field of view of each plot is ∼ 2.5 h−1
100
Mpc on each side. The positions of all galaxies with concordant
radial velocity found in NED within 90′ are indicated with each
galaxy’s morphological type. The 90′ search radius around the
E pair member is indicated by the solid circle. New members
found in the WFI field (central square) are plotted as crosses. The
dashed circle represents the 0.5 h−1
100 Mpc limit found in Fig. 15,
beyond which the velocity dispersion starts to rise again. The
central circle is the luminosity-weighted mean harmonic radius
RH centred on the optical group centre calculated with group
members inside the 0.5 h−1
100 Mpc radius (solid line) and with the
whole sample (dotted line), see Table A.1 and Table 5 for the
values.
around the bright elliptical, while around RR 143 there is no central concentration. The situation for the other two pairs is unclear
due to the high and asymmetrical incompleteness. In Fig. 17, the
radial density profiles of the faint members of our E+S systems
are compared with the profile measured for the isolated elliptical
and fossil group candidate NGC 1132 and the composite profile
of five X-ray bright groups (see Fig. 3 in Mulchaey & Zabludoff,
1999). In contrast to our radial density profiles that are computed
using only spectroscopically confirmed member galaxies, these
two last profiles were compiled by assuming that all the detected
faint galaxies in each group field belong to the group. Since the
background galaxies are not expected to concentrate around the
bright elliptical but instead be uniformly distributed in the field,
we expect a roughly constant shift in number density throughout
the field. This is indeed the case for our richest system RR 242,
in which the radial density profile is similar to that of the five
Mulchaey & Zabludoff (1999) groups.
4.2. The E+S system luminosity function and dwarf-to-giant
ratio
To simplify comparison with luminosity functions in the literature, absolute magnitudes are computed in +5 log h100 mag.
We consider separately the distributions of absolute magnitudes
50
100
150
Fig. 17. Radial density profile of faint group member galaxies.
The groups are colour-coded and marked with different symbols.
Note the concentration of faint galaxies towards the centre (the
pair E member) in RR 242, but not in RR 143. For comparison
are shown the radial distribution of galaxies in the groups studied
by Zabludoff & Mulchaey (1998) and of the galaxies in the fossil
group around NGC 1132 (see Mulchaey & Zabludoff, 1999).
for all member galaxies: a) within the WFI field of view, and
b) within the 90′ radius (see Appendix A). OLFs are computed
with the completeness-corrected galaxy counts only within 0.5
h−1
100 Mpc in order to compare the same physical region in all
groups. Figure 18 shows the distribution of absolute magnitudes
in the four groups, both for members located within the WFI
field of view and for all members within a 90′ radius as given in
Table A.1. The mean absolute magnitudes of the WFI field subsamples do not differ significantly from the entire group values.
The values are consistently around MR ∼ -18 mag apart from
RR 216, which has a higher value due to both its brighter elliptical pair member and other very bright members in the larger
scale environment.
It is remarkable that very few faint members were found.
Confirmed companions tend to show intermediate luminosities
in the supposed transition region from “normal” to dwarf galaxies (MR ∼ −18 (Fergusson & Binggeli, 1994)). The intermediate region tends to be populated by faint S0s, spirals, and
dwarf ellipticals. Despite the high number of very faint candidates (the median of all 4 candidate samples is fainter than
MR ∼ −14 + 5 log h100 ), no faint dwarf-irregular galaxies were
found.
The dependence of galaxy magnitude on position within a
group is also investigated in Fig. 18. The trend is that the brightest galaxies in the groups are more centrally concentrated than
the galaxies of intermediate luminosity, as already discussed in
the analysis of peculiar velocities (Fig 10). In contrast to RR 143,
RR 216 and RR 242 contain some brighter members outside ∼
0.6 h−1
100 Mpc. This region lies outside the investigated area for
RR 210. The lack of faint galaxies outside ∼ 0.2 h−1
100 Mpc is
easily explained since it corresponds to the approximate size of
20
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
10
RR 143
10
5
5
0
0
-20
10
RR 210
-15
RR 216
-20
10
5
5
0
0
-20
-15
-15
RR 242
-20
-15
Fig. 18. Distribution of absolute magnitudes. Left: Absolute magnitude histogram of members of the 4 groups in the WFI field (red)
and in the 90′ radius. The vertical dashed lines indicate the mean absolute magnitude of group members within the WFI field (red,
long dashed) and the 90′ radius (black, short dashed). Right: Absolute magnitude vs. projected radial distance from the group centre.
Galaxies of each group are marked with different symbols: solid triangles: RR 143; open circles: RR 210; open triangles: RR 216;
solid squares: RR 242.
the WFI field of view over which we searched for faint member
galaxies.
Figure 19 shows the combined OLFs of the two groups
RR 143 and RR 242 with extended hot IGM (labelled X-ray
bright in the figure) and the two groups RR 210 and RR 216,
which are X-ray faint. The galaxy counts per magnitude are
completeness-corrected but not normalised to the surveyed area.
To compare the same physical area in all groups, only galaxies within 0.5 h−1
100 Mpc projected distance from the pair elliptical were considered (the area sampled in the closest group
RR 210). Before combining them, the OLF of each group was
normalised to have the same number of members brighter than
a certain magnitude. We used MR = −15 + 5 log h100 for the
X-ray bright groups, since the spectroscopic completeness was
higher than 50% in all bins down to this magnitude. For the Xray faint groups, only the numbers above MR = −17 + 5 log h100
are used due to the high incompleteness in RR 216 below MR =
−17 + 5 log h100 (< 10% completeness, see Fig. 4). The OLFs of
both groups were then averaged. This procedure ensures that the
shape of the combined OLF is not weighted towards the richer
group.
An analytic form of the luminosity function fitted to the observed distribution of magnitudes was described by Schechter
(1976). It has a steeply rising slope at the bright end that then
levels off or even decreases at fainter magnitudes. It has 3 parameters: M ∗ , the transition between bright and faint end slope, α,
the faint end slope, and Φ∗ , the galaxy density at M ∗ , which gives
the normalization of the OLF. This latter parameter gives the
number of galaxies per Mpc3 and is difficult to compute, since
it requires a survey covering an extensive area. Since it does not
affect the shape of the OLF, it is often just adjusted to fit the number of observed galaxies. Our spectroscopic completeness drops
below 50% for magnitudes fainter than MR = −15 + 5 log h100 .
Therefore, we decided to adjust the OLFs from the literature to
our number of galaxies brighter than MR = −15 + 5 log h100 .
Galaxies fainter than this limit are not considered further here.
For the X-ray faint groups, the OLFs are adjusted to the galaxy
number brighter than MR = −17 + 5 log h100 . Since we lack information about faint galaxies outside the WFI field of view,
we compared the OLFs from the literature to our galaxy numbers within the WFI field (red squares). For comparison, the
−2
counts for all galaxies within a 0.5 h−1
100 Mpc radius (∼ 0.8 h100
2
Mpc ) are also plotted (black triangles). The OLF of a sample of simulated fossil groups of galaxies (into which our pairs
could evolve) is taken from D’Onghia et al. (2005) (magenta
crosses). The solid line represents the OLF for X-ray detected
poor groups found by Zabludoff & Mulchaey (2000). They investigated a sample of five poor galaxy groups with X-ray halos of comparable luminosity to our two X-ray bright pairs
and found that their OLF can be fitted by a Schechter function with M ∗ = −21.6 + 5 log h100 and α = −1.3. These values
are comparable to those measured for nearby rich cluster OLFs
(see e.g., Trentham, 1996; Driver et al., 1998). The short-dashed
line shows the OLF of the local field from the Las Campanas
Redshift Survey (LCRS hereafter) found by Lin et al. (1996)
with M ∗ = −20.29+5 log h100 and α = −0.7. LCRS results agree
closely with those for a large sample of the most isolated galaxies in the northern sky (AMIGA sample, Verdes-Montenegro et
al. 2005). Finally, the long-dashed line represents the OLF of the
spectroscopically confirmed members of a fossil galaxy group
from Mendes de Oliveira et al. (2006).
The possible effects of incompleteness and low number
statistics were simulated again using the Monte Carlo method.
A set of 1000 fake groups with the same number of members
as our combined group samples and whose members follow a
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
21
Fig. 19. Optical Luminosity Functions (OLF) of the 4 groups. The OLFs of the 2 X-ray bright and the 2 X-ray faint groups are
combined respectively. The OLFs are computed for the large scale sample (black triangles) and the WFI field subsample (red
squares). Note that outside the WFI field the radial velocity information comes from NED and is highly incomplete at fainter
magnitudes. The magenta crosses show the OLF of a sample of simulated fossil groups from D’Onghia et al. (2005). The solid line
is the OLF found for a sample of X-ray bright poor groups by Zabludoff & Mulchaey (2000). The short dashed line shows the OLF
of the local field (Lin et al., 1996), whereas the long dashed line is the OLF of a fossil group observed by Mendes de Oliveira et al.
(2006). All OLFs use spectroscopically confirmed members only and are completeness corrected in a similar way as our counts.
The shaded regions around each OLF shows the 1σ deviations expected due to our low number statistics obtained from a set of
Monte Carlo Simulations (see text).
certain Schechter OLF was constructed. To do this, a random
pair of values in the plane of galaxy magnitude versus galaxy
number was produced. If this data point happened to lie under
the curve of the respective OLF, this galaxy magnitude was included in the group sample. This procedure continued until the
desired number of members was reached and then repeated for
the next group. Finally, the magnitude histogram was computed
for each group and the mean and rms of all group histograms
were computed in each bin. This was done for the group, field
and fossil group OLFs with the number of members of both, the
X-ray-bright (left panel) and the X-ray faint (right panel) combined group. Only galaxies down to MR = −15 + 5 log h100 in the
X-ray-bright and MR = −17+5 log h100 in the X-ray faint groups
were considered. The expected 1σ deviation of each OLF is
shaded, horizontally for the group OLF (Zabludoff & Mulchaey,
2000, ZM00 hereafter), and diagonally for the field (Lin et al.,
1996) and fossil group (FG) (Mendes de Oliveira et al., 2006)
OLFs (which are quite similar).
The combined OLF of the two X-ray bright groups shows evidence of a non-Schechter form with an excess of bright galaxies
or a pronounced dip at MR ∼ −20 + 5 log h100 . The same dip is
present in the X-ray faint group OLF (red squares), although it
disappears at larger radii, outside the WFI field of view (black
triangles). To test the significance of this dip in the X-ray bright
groups with respect to the X-ray faint groups, we compared the
counting errors of the X-ray bright and X-ray faint groups in
each magnitude bin (black triangles). The difference in counts
between the two OLFs exceeds the sum of their errors in the
bins from MR = -20.5 to -18.5 +5 log h100 , which suggests that
they indeed differ. For the sake of clarity, we do not show the
counting errors for each bin in Fig. 19.
The X-ray bright OLF is inconsistent with the OLF of X-ray
detected poor groups found by ZM00 (at more than 1σ). The
22
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
Table 6. Dwarf-to-giant (D/G) ratios obtained from the OLFs.
group
RR 143
RR 210
RR 216
RR 242
Combined X-ray bright
Combined X-ray faint
D/G0.5
1.3±0.7
4.2±1.4
3.0±1.0
3.4±1.1
2.2±0.6
3.5±0.8
D/G0.5,17
0.0±0
1.8±0.7
1.0±0.4
0.9±0.4
0.4±0.2
1.3±0.4
D/GWFI
1.9±1.2
3.9±2.1
8.0±3.8
4.3±1.6
3.0±0.9
5.8±2.0
D/GWFI,17
0.0±0
2.0±1.2
2.5±1.4
0.9±0.5
0.4±0.2
2.2±0.9
Column 2: D/G0.5 : dwarf-to-giant ratio for galaxies within 0.5 h−1
100 Mpc.
Column 3: D/G0.5,17 : dwarf-to-giant ratio for galaxies within 0.5 h−1
100 Mpc using
only dwarfs brighter than MR = −17 + 5 log h100 .
Column 4: D/GWFI : dwarf-to-giant ratio for galaxies within the WFI field of
view.
Column 5: D/GWFI,17 : dwarf-to-giant ratio for galaxies within the WFI field of
view using only dwarfs brighter than MR = −17 + 5 log h100 .
LCRS field OLF is also unable to fit our observations because
there are not enough giants of intermediate luminosity (below
MR ∼ −20 + 5 log h100 ). The X-ray faint group OLF on the other
hand is consistent with the OLF of the ZM00 groups, although
there remains an excess of bright giant galaxies. However, within
the rms expected from the Monte Carlo simulations, the X-ray
faint group OLF is consistent with the OLF of the ZM00 groups.
Interestingly, the OLF of the simulated and observed fossil
groups (FGs) are quite different, although this difference is based
on only one observed group. The OLF of the sample of simulated FGs shows a prominent gap between -20 ≤ MR ≤-19 + 5
log h100 mag, whereas the observed FG reaches its maximum in
this range. A similar gap in the OLF of X-ray faint groups was
also found by Miles et al. (2004), however, their group membership is based on galaxy colours and not on spectroscopy as in this
study. The same lack of galaxies is present in our groups, both Xray bright and faint. In the X-ray faint groups however, this gap
vanishes if a larger area is considered (black triangles in Fig. 19),
similar to the findings of Miles et al. (2006). They found that the
gap in the infrared LF of X-ray faint groups is vanishing at larger
radii. However, the gap in our combined X-ray bright group OLF
remains and is not an effect of the small survey area. The OLF of
the observed FG seems to drop at fainter magnitudes, although
the authors indicate that this may be due to their limiting magnitude of MR ∼ -18 + 5 log h100 mag and may actually be rather
a dip at this magnitude than a continuing drop towards fainter
magnitude. This makes the FG OLF consistent with our X-ray
faint group OLF, although it cannot be distinguished from the
other two OLFs with our low number statistics, as shown by the
results of the Monte Carlo simulations (shaded areas).
Incompleteness effects are certainly an issue in our sample, although they are not expected to play a major role down
to MR ∼ −17 + 5 log h100 . Apart from 3 galaxies in the range
−19 ≤ MR ≤ −17 that were missed in RR 143, all candidates
without measured redshift above MR = −17 are accounted for
by the completeness correction. If these 3 missed candidates
were members of RR 143, the OLF of the combined X-ray
bright groups would be slightly higher in this range and closer to
the ZM00 OLF. However, this would influence neither that gap
above -19, nor the excess of bright galaxies relative to the other
OLFs, and it would not reach into the 1σ area of the field and
fossil group OLFs. So our combined X-ray bright group OLF
would still significantly differ from the OLFs in the literature.
Dwarf-to-giant (D/G) ratios are determined from the
completeness-corrected counts for both individual groups and
for the two combined samples. The limit between giants
and dwarfs is again MR = −19 + 5 log h100 following
Zabludoff & Mulchaey (2000). The values are computed for the
full dwarf sample and for galaxies brighter than MR = −17 +
5 log h100 , so as to compare with literature values. We also computed D/G for all galaxies within the WFI field of view since
outside this field information about faint member galaxies is very
poor. The errors are the combined counting errors of each dwarf
and giant sample. All values are given in Table 6. An average
D/G value for the Zabludoff & Mulchaey (2000) X-ray luminous poor groups (down to MR = −17 + 5 log h100 ) is 1.9 ±
0.4 compared to a value of less than 1.0 for local group galaxies in this magnitude range (Pritchet & van den Berg, 1999).
Our X-ray faint groups have higher D/G ratios than our X-ray
bright groups. The value of the combined X-ray faint E+S systems is close to the value found for the Zabludoff & Mulchaey
(2000) X-ray luminous poor groups (D/GW FI, 17 = 2.2 ± 0.9),
whereas the D/G of our X-ray bright systems is significantly
lower (D/GW FI, 17 = 0.4 ± 0.2). By considering the two X-ray
bright E+S systems to be the groups with more complete coverage, the result is unlikely to be caused mainly by incompleteness
effects. However, in the field of RR 143, three candidates in the
range -19 +5 log h100 < MR < −17 + 5 log h100 (see Fig. 3), were
not observed with VIMOS. Assuming that these 3 galaxies were
all members of the E+S system, the D/G would reach a value of
∼ 1, similar to that of RR 242. This value is still lower than the
D/G of 1.9 from Zabludoff & Mulchaey (2000), as well as the
D/G ratio found for our X-ray faint E+S systems.
4.3. What maps the hot inter-galactic medium?
Groups with an extended hot IGM usually have a giant elliptical that is typically the brightest group member and located
near or at the peak of the “smooth, symmetric” X-ray emission (see e.g., the review in Mulchaey, 2000). Another distinctive characteristic of groups with a hot IGM component was
found by Zabludoff & Mulchaey (1998). By means of multiobject spectroscopy, they found a significantly higher number of
faint galaxies (∼ 20 - 50 members down to magnitudes as faint
as MB ∼ −14 + 5 log h100 to -16 +5 log h100 ) in groups with a
significant amount of hot IGM.
Analysis of the faint population in our E+S systems suggests that the presence of extended diffuse X-ray emission is not
necessarily connected to the presence of such a numerous faint
galaxy population. At the same time, our E+S systems appear
to be gravitationally bound structures. The velocity dispersion
profile suggests that the dispersion is not constant with radius. It
shows a maximum at ∼ 0.2 h−1
100 Mpc and decreases until ∼ 0.5
−1
h100 Mpc before increasing again at higher projected radii from
the optical group centre. This suggests that the groups have a
dynamical boundary at ∼ 0.5 h−1
100 Mpc. Furthermore, dynamical analysis of the four E+S systems indicates that the pair is
displaced from the optical group centre, suggesting that the hot
IGM is a “local” phenomenon. The X-ray emission is centred on
the E member of the pair indicating that it may be the principal
reason for the presence of the hot X-ray emitting gas.
In Paper III, we showed that RR 143 and RR 242 have a luminous, extended, hot IGM, while both RR 210 and RR 216,
although of similar “optical” and kinematical characteristics, are
X-ray underluminous with respect to other loose groups or mature Es, if we consider the emission connected with the elliptical galaxy. The diffuse X–ray emission from the hot intragroup medium (IGM) detected in compact but also in loose,
poor groups, has often been taken as a direct measure of the
group potential (see e.g., Mulchaey, 2000). At the same time,
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
Sansom et al. (2000) offered another interpretation of the large
spread in X-ray luminosity among ETGs. They suggest that ellipticals, and in general early-type galaxies, showing fine structure – such as e.g., shells and dust-lanes – tend to have a fainter
X-ray luminosity, although the dispersion is very large. These
early-type galaxies are considered dynamically young i.e., showing evidence of recent accretion/merging events.
Several young systems, although in small groups, do not
show extended emission and lack a substantial group component
(O’Sullivan et al., 2001). XMM-Newton and Chandra observations found three underluminous elliptical galaxies (NGC 3585,
NGC 4494, and NGC 5322) all of which show evidence of recent
dynamical disturbances, including kinematically distinct cores
as in the case of NGC 474 (Hau et al., 1996). Rampazzo et al.
(2006) found that the X-ray luminosity of NGC 474, the earlytype member of another E+S system (Arp 227) in the northern
hemisphere, lies about two orders of magnitude below that of
dominant group members and is located in the area of the log LB
- log LX plane where the X-ray emission could be explained by
the superposition of discrete X-ray sources.
The position of interacting or post-interacting galaxies, such
as those exhibiting fine structures or kinematical perturbations,
in the LX - LB plane is consistent with the hypothesis that their
X-ray emission comes from discrete sources only, although their
LX /LB ratios are not as low as that of NGC 474. Sansom et al.
(2000) and O’Sullivan et al. (2001) interpreted the negative trend
observed between LX /LB (linked to the gas content in early-type
galaxies ) and the morphological disturbance quantified by the
fine-structure parameter Σ (linked to the age/dynamical stage) as
evidence that several gigayears are required to accumulate hot,
gaseous halos, so that recent mergers/young systems are deficient in hot gas. O’Sullivan et al. (2001) also attributed some
of the scatter seen in the global LX versus LB relation to the
evolutionary stage and past merger history of early-type galaxies. Brassington et al. (2007) studied the X-ray emission of nine
merging systems believed to represent different phases of the
merging process. They suggested that (1) the X-ray luminosity
peaks ∼300 Myr before nuclear coalescence; (2) at a time ∼1 Gyr
after coalescence, the merger remnants are fainter compared to
mature ellipticals; while (3) at a greater dynamical age (≥3 Gyr)
remnants start to resemble typical ellipticals in their hot gas content. On these grounds, the above authors support the idea that a
halo regeneration takes place within low LX merger remnants.
A possible explanation of the diverse X-ray properties of our
E+S systems could then be connected with the dynamical age of
the dominant E galaxy. Groups with similar environments will be
in different evolutionary phases that can be traced by the giant
central elliptical galaxy: interaction (accretion/merging) disrupts
the hot gas halo, which is then built up again during ongoing
evolution on a timescale of a few gigayears. This implies that
the majority of the X-ray emission is unlikely to be associated
with the group potential, as also suggested by the kinematics
and dynamics of these systems. In this scenario, RR 242 would
be the most evolved system, and RR 216, which is expected to
have a significant number of faint companions and faint X-ray
emission might be experiencing an active phase of dynamical
evolution. The same might be valid for RR 210, which, despite
its apparent lack of faint companions, is embedded in a compact
larger-scale structure (in projection and in redshift space). The
apparent lack of faint companions concentrated around the E+S
pair might be due to incompleteness. Another possible explanation of the relative X-ray faintness of these two systems is that
they are still in the process of collapse, so the IGM has yet to be
compressed and heated to X-ray temperatures (Rasmussen et al.,
23
2006). However, this seems unlikely regarding the short crossing times of the pairs, arguing for a virialization at least of the
group centres. The clear signatures of interaction found in the
two pairs, including the presence of diffuse light suggests a longlasting coevolution of the pair galaxies. The case of RR 143 is
quite different with a real lack of faint galaxies. It is interesting
that the scarcity of faint members continues on larger scales. We
find only two other galaxies of similar redshift within ∼ 1 h−1
100
Mpc.
5. Summary and conclusions
We have presented VLT-VIMOS observations in a search for
faint galaxy members of four E+S systems. Candidate members were identified by applying photometric criteria to WFI images covering a field of view of about 0.2 h−1
100 Mpc radius (see
section 2). We investigated the morphological and photometric
characteristics of the new group members as well as their spectral properties. We used the new data to determine the group
dynamics as well as the combined group luminosity functions.
We found the following results:
1. Two and ten new members are confirmed for RR 143 and
RR 242, respectively. We found two and three new members
associated with RR 210 and RR 216, respectively, which are
both only partially covered by our VIMOS observations. The
new members increase the galaxy populations to 4, 7, 6, and
16 members in a 0.5◦ × 0.5◦ field around the pairs RR 143,
RR 210, RR 216, and RR 242, respectively, down to MR ∼
-12 +5 log h100 . We applied incompleteness corrections that
were necessary for the subsequent investigations.
2. The morphological study of the new members based on
our detailed surface photometry indicated a high fraction of S0 galaxies (40%), most with low B/T ratios. A
morphology-radius relation is apparent for the combined
group sample. Galaxies with higher B/T ratios appear to be
more concentrated towards the field centres (i.e., towards
the E member of the pair), although there is a large spread
of B/T ratios. However, the relation between morphology
and local projected number density seems to be more
significant (at the 95% confidence level) arguing for the
presence of a morphology-density relation. This suggests
that the very local environment has a strong influence on
galaxy morphologies and is responsible for shaping our faint
galaxies. Signatures of interaction and merging are found
in the group sample. Asymmetries, filaments, and shells are
detected in several galaxies.
3. Spectra of the new members are indicative of an old stellar
population in the vast majority of galaxies. No blue objects,
dIrr or tidal dwarfs are present in our sample. This is
also supported by the colour-magnitude relation for group
galaxies. Although we reach into the domain of dwarf
irregular (dIrr) and spheroidal (dSph) galaxies that we find
in the Local Group, many of them would likely be below
our detection limit.
4. Our dynamical analysis indicates short crossing times for
all systems suggesting that at least the centres of the groups
are virialised. The E pair members dominate the groups: in
RR 143 and RR 210, they represent ∼ 1/2 and in RR 216
and RR 242 ∼ 1/3 of the total group light. The dynamical
quantities appear uncorrelated with group X-ray luminosity.
24
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
RR 242 has the highest velocity dispersion and virial mass,
while RR 143 has the lowest (both are X-ray bright). The harmonic and virial radius of the groups are similar for RR 143,
RR 210, and RR 242. In RR 216, member galaxies are less
concentrated towards the group centre. This may be due to
the different large-scale environment of this pair, which is in
the outskirts of the Hydra-Centaurus cluster region.
The velocity dispersion seems to vary with distance from
the optical group centre. The pair ellipticals being the
brightest group galaxies are more centrally concentrated
than giants and dwarfs, but still about 0.1 h−1
100 Mpc from the
group centre (apart from RR 210, which is located precisely
at the optical group centre). This also means that the hot
IGM, which is centred on the elliptical, is shifted from the
optical group centre. The X-ray emission then seems to
be connected with the bright elliptical of the pair and its
evolutionary phase rather than with the group environment.
5. The OLF of X-ray bright E+S systems differs from the OLF
of X-ray faint systems, in line with the findings of previous
authors (Miles et al., 2004). The X-ray bright E+S system
OLF also differs from both that of the sample of X-ray luminous poor groups in Zabludoff & Mulchaey (2000), and
the OLF of the local field and isolated galaxies (Lin et al.,
1996). This comparison suggests that the OLF of poor Xray detected galaxy systems is not universal, in contrast to
the results of Zabludoff & Mulchaey (2000). Despite the environmental differences between the two X-ray luminous
groups in our sample, their normalised OLFs are quite similar, showing a lower D/G ratio (compared to the groups
in Zabludoff & Mulchaey (2000)) with equal numbers of
dwarfs and giants. Within the giant regime (MR brighter than
−19 + 5 log h100 ), our X-ray bright groups also show an interesting behaviour: they have a higher number of bright galaxies or, to put it in a different way, they lack galaxies between 20 ≤ MR ≤ -19 + 5 log h100 mag. Their OLF is reminiscent of
the luminosity function of NGC 5846, which is a group considered to be dynamically evolved (Mahdavi et al., 2005).
The X-ray bright OLF is also comparable to the OLF of a
sample of simulated fossil groups (D’Onghia et al., 2005)
showing a similar gap around MR ∼ -19 + 5 log h100 mag.
This could indicate that our X-ray bright E+S systems are
more dynamically evolved than the Zabludoff & Mulchaey
(2000) groups and that the E pair members are the remnants
of this evolution.
On the other hand, the OLF of the X-ray faint E+S systems
agrees with the OLF of the Zabludoff & Mulchaey (2000)
groups and these systems, furthermore, have a similar D/G
ratio (∼ 2). The X-ray faint group OLF also agrees very well
with the spectroscopically confirmed OLF of an observed
fossil group (Mendes de Oliveira et al., 2006). The X-ray
faint groups may thus be a phase in the dynamical evolution
of the Zabludoff & Mulchaey (2000) groups where the
recent or ongoing interaction, in which the E member is
involved, could have destroyed or at least decreased the
luminosity of the IGM. The X-ray halo could then be
built up again during the subsequent passive evolution of
the elliptical (e.g., Sansom et al., 2000; Brassington et al.,
2007). Their OLF would be consistent with these systems
evolving into fossil groups similar to that observed by
Mendes de Oliveira et al. (2006).
Interaction-induced rejuvenation episodes may be present in
a small fraction of our sample as suggested by the presence of
fine-structures. The E members in both RR 210 and RR 216 are
also good candidates for showing rejuvenation signatures in their
stellar population. This finding would reinforce our hypothesis
that their faint X-ray emission is connected with the phase of
their dynamical evolution. The presence of a young stellar population in both the giant and faint galaxy members can be ascertained by the study of absorption line-strength indices that will
be carried out in a future paper.
Acknowledgements. We thank the anonymous referee for many useful suggestions that certainly improved the content of this paper. RG wants to thank Jesús
Varela for many helpful discussions and his assistance in preparing the OLFs. RG
and RR thank the INAF-Osservatorio Astronomico di Padova and the Institut für
Astronomie Universität Wien, respectively, for the kind hospitality during the
paper preparation. RG, RR and WWZ acknowledge the support of the Austrian
and Italian Foreign Offices in the framework science and technology bilateral
collaboration (project number 25/2004). RG and WWZ acknowledge the support of the Austrian Science Fund (project P14783). RR and GT acknowledge the partial support of the Agenzia Spaziale Italiana under contract ASIINAF I/023/05/0. 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. The Digitized Sky Survey (DSS) was produced at the Space
Telescope Science Institute under U.S. Government grant NAG W-2166. The images of these surveys are based on photographic data obtained using the Oschin
Schmidt Telescope on Palomar Mountain and the UK Schmidt Telescope. The
plates were processed into the present compressed digital form with the permission of these institutions. This research has made use of SAOImage DS9,
developed by Smithsonian Astrophysical Observatory.
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Appendix A: Properties of group members used for
dynamical calculations
Since our WFI observations probably do not cover the complete extension of each group, we searched for additional group
member galaxies in the environment of each pair outside the
WFI field (see also Paper III). The NASA/IPAC Extragalactic
25
Database (NED) was used for this search. The area investigated was 90′ corresponding to ∼ 1 h−1
100 Mpc for the farthest
pair. We also chose this radius to mimic the observed area of
Zabludoff & Mulchaey (1998), the group sample to which we
compared our results. The maximum velocity difference within
which a galaxy was considered to be a group member was chosen to be ∆v ≤ 1000 km s−1 .
Previous authors have proposed a luminosity-weighted
formulation to calculate the group dynamics (e.g.,
Fergusson & Sandage, 1990; Firth et al., 2006). To investigate the differences introduced by weighting, we calculated
the group dynamics by using both luminosity weighting and
uniform weights (wi = 1).
In the luminosity-weighted approach, each galaxy was
weighted by its relative luminosity calculated from the galaxy’s
R-band magnitude
wi = 10−0.4 mRi
(A.1)
The coordinates of the optical group centre αgroup and δgroup are
given by averaging the luminosity-weighted coordinates αi and
δi of all group members
P
P
δi wi
i αi wi
αgroup = P
and δgroup = Pi
(A.2)
w
i i
i wi
The luminosity-weighted mean velocity, i.e., the velocity of the
optical group centre was calculated similarly
P
vi wi
(A.3)
vgroup = Pi
i wi
The luminosity-weighted line-of-sight velocity dispersion
was calculated by summing over the weighted squared-deviation
of each group member from the group velocity
σr =
"P
i
wi (vi − vgroup )2
P
i wi
# 12
(A.4)
The mean harmonic radius is a measure of the compactness
of the group and was calculated from the projected separations
Ri j between the i-th and j-th group member
#−1
"P P
i
j<i (wi w j )/Ri j
P P
(A.5)
RH =
i
j<i wi w j
The virial radius was connected to RH with
Rvir =
π RH
2
(A.6)
The crossing time tc was computed from the mean distance
of group members from the optical group centre, hri, and the
mean velocity relative to the group centre, hvi, following the definition of Rood & Dickel (1978)
tc =
hri
hvi
(A.7)
Multiplying tc with H0 gives the crossing time in units of the
age of the universe, independent of the choice of H0 .
Different mass estimators can be found in the literature.
Following the discussion in Heisler et al. (1985), we computed
the virial mass Mvir and the projected mass MP
Mvir =
3 2
σ Rvir
G r
(A.8)
26
MP =
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
32
πG
P
i
wi (vi − vgroup )2 r⊥i
P
i wi
(A.9)
The mass-to-light ratio M⊙ /L⊙ was computed from the virial
mass and the group luminosity obtained by summing up the luminosities of the individual galaxies. These were calculated from
the absolute magnitude of the galaxies obtained by the distance
modulus given in Table 5 and the absolute magnitude of the Sun
MR⊙ = 4.42 mag (taken from Binney & Merrifield, 1998)
Li = 100.4(MR⊙ −MRi )
(A.10)
Table A.1 provides the properties of the new group members
within the WFI field (ID starting with “RR”) as well as galaxies
found in the NED. All galaxies as well as only galaxies within
0.5 h−1
100 Mpc from the optical group centre (r⊥ , Cols. 3 and 8)
were used to calculate the group dynamics given in Table 5. The
peculiar velocity v pec (v pec = vi − vgroup ) and the projected radius
from the optical group centre r⊥ were used in the analysis of the
dependence of velocity dispersion on radius (Fig. 15).
The last column in Table A.1 gives the weights normalised
to the sum over all weights, i.e., their contribution to the total
luminous mass. The absolute magnitudes MR were used to calculate the luminosity LR (relative to the Sun), which was then
converted into the normalised weights wni by normalising each
LR by the sum of all luminosities (i.e., the group luminosity).
Columns 2-6 give the luminosity-weighted values, whereas
uniform weights were used for values in cols. 7-11.
R. Grützbauch et al.: Small-scale systems of galaxies. IV.
27
Table A.1. Properties of group members used for dynamical calculations.
galaxy
RR143 09192
RR143 24246
NGC 2305 (RR 143a)
NGC 2307 (RR 143b)
NGC 2297
ESO 087- G 050
-25.
403.
-80.
194.
-79.
217.
r⊥
MR
LR
[Mpc]
[mag]
[L⊙ ]
luminosity-weighted
0.013
-16.23
1.81E+08
0.187
-16.53
2.39E+08
0.063
-22.50
5.84E+10
0.101
-21.72
2.85E+10
0.355
-21.17
1.72E+10
0.671
-18.78
1.90E+09
RR210 11372
RR210 13493
NGC 4105 (RR 210a)
NGC 4106 (RR 210b)
2MASX J12063106-2951336
IC 2996
IC 3005
2MASX J12052132-3002465
6dF J1207305-301156
IC 0760
2MASX J12083457-3008549
ESO 441- G 004
MCG -05-29-004
IC 3010
2MASX J12041324-2918486
6dF J1206496-302745
IC 0764
ESO 441- G 014
AM 1207-294 NED02
ESO 440- G 044
ESO 440- G 039
ESO 441- G 011
6dF J1203467-284015
-91.
-10.
-82.
132.
180.
238.
-294.
-2.
516.
208.
139.
300.
120.
-36.
-15.
277.
114.
128.
146.
180.
27.
118.
-111.
0.064
0.061
0.020
0.014
0.042
0.109
0.089
0.152
0.154
0.189
0.180
0.213
0.236
0.212
0.263
0.240
0.253
0.273
0.290
0.398
0.411
0.393
0.458
-14.80
-18.13
-22.86
-21.53
-17.79
-19.11
-19.68
-17.49
-16.39
-20.34
-17.89
-17.19
-18.39
-20.71
-17.69
-17.59
-20.49
-18.50
-18.89
-19.19
-17.42
-17.61
-17.39
4.89E+07
1.05E+09
8.19E+10
2.41E+10
7.68E+08
2.59E+09
4.38E+09
5.83E+08
2.12E+08
8.05E+09
8.42E+08
4.42E+08
1.34E+09
1.13E+10
7.01E+08
6.39E+08
9.24E+09
1.48E+09
2.12E+09
2.79E+09
5.46E+08
6.51E+08
5.32E+08
0.0003
0.0067
0.5242
0.1540
0.0049
0.0166
0.0280
0.0037
0.0014
0.0515
0.0054
0.0028
0.0085
0.0724
0.0045
0.0041
0.0591
0.0095
0.0135
0.0178
0.0035
0.0042
0.0034
-186.
-105.
-177.
37.
85.
143.
-389.
-97.
421.
113.
44.
205.
25.
-131.
-110.
182.
19.
33.
51.
85.
-68.
23.
-206.
0.080
0.050
0.007
0.010
0.046
0.110
0.111
0.154
0.178
0.176
0.209
0.212
0.224
0.239
0.250
0.262
0.278
0.304
0.321
0.390
0.416
0.403
0.453
-14.90
-18.23
-22.96
-21.63
-17.89
-19.21
-19.78
-17.59
-16.49
-20.44
-17.99
-17.29
-18.49
-20.81
-17.79
-17.69
-20.59
-18.60
-18.99
-19.29
-17.52
-17.71
-17.49
5.37E+07
1.15E+09
8.99E+10
2.64E+10
8.43E+08
2.84E+09
4.80E+09
6.39E+08
2.32E+08
8.82E+09
9.24E+08
4.85E+08
1.46E+09
1.24E+10
7.68E+08
7.01E+08
1.01E+10
1.62E+09
2.32E+09
3.06E+09
5.99E+08
7.14E+08
5.83E+08
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
RR216 03519
RR216 04052
RR216 12209
NGC 4373 (RR 216b)
IC 3290 (RR 216a)
ESO 322-IG 002
NGC 4373A
IC 3370
ESO 322- G 009
ESO 321-IG 028
ESO 322- G 011
2MASX J12294019-4007220
ESO 322- G 007
MCG -06-27-023
ESO 321- G 021
ESO 322- G 019
ESO 321- G?026
ESO 322- G 020
2MASX J12185570-4005358
NGC 4499
2MASX J12301164-3845537
ESO 322- G 017
ESO 322- G 024
493.
-30.
-526.
173.
119.
41.
-288.
-293.
353.
-151.
19.
-111.
-26.
176.
-34.
-123.
-108.
200.
-28.
506.
260.
175.
-66.
0.097
0.037
0.075
0.107
0.125
0.234
0.188
0.237
0.290
0.395
0.480
0.468
0.606
0.563
0.752
0.660
0.658
0.674
0.810
0.671
0.659
0.692
0.715
-16.15
-18.84
-17.83
-23.32
-21.98
-19.18
-21.99
-23.02
-20.53
-20.15
-19.83
-17.20
-18.79
-19.91
-21.34
-19.81
-19.26
-20.29
–
-21.21
–
-19.03
-18.21
1.69E+08
2.02E+09
7.95E+08
1.25E+11
3.63E+10
2.76E+09
3.67E+10
9.47E+10
9.56E+09
6.73E+09
5.02E+09
4.45E+08
1.92E+09
5.40E+09
2.02E+10
4.92E+09
2.97E+09
7.66E+09
–
1.79E+10
–
2.40E+09
1.13E+09
0.0004
0.0052
0.0021
0.3247
0.0945
0.0072
0.0954
0.2463
0.0249
0.0175
0.0130
0.0012
0.0050
0.0140
0.0524
0.0128
0.0077
0.0199
0.0000
0.0465
0.0000
0.0062
0.0029
461.
-62.
-558.
141.
87.
9.
-320.
-325.
321.
-183.
-13.
-143.
-58.
144.
-66.
-155.
-140.
168.
-60.
474.
228.
143.
-98.
0.124
0.053
0.083
0.118
0.136
0.258
0.211
0.238
0.309
0.425
0.464
0.446
0.598
0.589
0.768
0.642
0.690
0.656
0.832
0.652
0.664
0.706
0.701
-16.17
-18.86
-17.85
-23.34
-22.00
-19.20
-22.01
-23.04
-20.55
-20.17
-19.85
-17.22
-18.81
-19.93
-21.36
-19.83
-19.28
-20.31
–
-21.23
–
-19.05
-18.23
1.73E+08
2.06E+09
8.11E+08
1.27E+11
3.71E+10
2.81E+09
3.74E+10
9.66E+10
9.75E+09
6.87E+09
5.12E+09
4.54E+08
1.96E+09
5.51E+09
2.06E+10
5.02E+09
3.03E+09
7.81E+09
–
1.82E+10
–
2.45E+09
1.15E+09
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
0.0435
RR242 08064
RR242 13326
RR242 15689
RR242 20075
RR242 22327
RR242 23187
RR242 24352
RR242 25575
RR242 28727
RR242 36267
NGC 5090 (RR 242a)
NGC 5091 (RR 242b)
NGC 5082
NGC 5090B
2MASX J13201668-4327195
NGC 5090A
ESO 270- G 007
2MASX J13181305-4330182
AM 1317-425
2MASX J13195606-4306498
2MASX J13192359-4417358
2MASX J13212941-4248564
2MASX J13161705-4307595
ESO 269- G 076
2MASX J13270026-4359472
ESO 270- G 014
ESO 269- G 072
ESO 269- G 069
-79.
-348.
211.
-475.
-283.
-120.
-823.
-865.
-503.
-523.
-99.
9.
376.
748.
-453.
-93.
230.
6.
-197.
-156.
-162.
-57.
390.
-537.
-342.
328.
-220.
-555.
0.222
0.095
0.108
0.012
0.068
0.135
0.106
0.080
0.115
0.138
0.096
0.105
0.036
0.100
0.150
0.114
0.336
0.263
0.371
0.361
0.376
0.555
0.566
0.560
0.756
0.939
0.755
0.781
-15.75
-17.45
-15.87
-17.27
-18.54
-16.23
-19.89
-17.18
-16.55
-14.93
-22.76
-21.17
-21.55
-20.89
-19.58
-21.34
-20.63
-18.50
–
–
-18.70
-17.60
-18.80
-19.10
–
-20.73
-21.26
-20.13
1.17E+08
5.61E+08
1.31E+08
4.75E+08
1.53E+09
1.82E+08
5.31E+09
4.37E+08
2.45E+08
5.51E+07
7.46E+10
1.73E+10
2.45E+10
1.33E+10
3.99E+09
2.02E+10
1.05E+10
1.48E+09
–
–
1.77E+09
6.44E+08
1.94E+09
2.56E+09
–
1.15E+10
1.87E+10
6.62E+09
0.0005
0.0026
0.0006
0.0022
0.0070
0.0008
0.0243
0.0020
0.0011
0.0003
0.3413
0.0789
0.1120
0.0610
0.0182
0.0923
0.0480
0.0067
0.0000
0.0000
0.0081
0.0029
0.0089
0.0117
0.0000
0.0526
0.0857
0.0303
85.
-184.
375.
-311.
-119.
44.
-659.
-701.
-339.
-359.
65.
173.
540.
912.
-289.
71.
394.
170.
-33.
8.
2.
107.
554.
-373.
-178.
492.
-56.
-391.
0.196
0.075
0.096
0.043
0.084
0.126
0.118
0.104
0.118
0.169
0.098
0.109
0.049
0.133
0.104
0.106
0.308
0.237
0.314
0.305
0.394
0.489
0.516
0.539
0.726
0.903
0.714
0.741
-15.65
-17.35
-15.77
-17.17
-18.44
-16.13
-19.79
-17.08
-16.45
-14.83
-22.66
-21.07
-21.45
-20.79
-19.48
-21.24
-20.53
-18.40
-3.50
-3.50
-18.60
-17.50
-18.70
-19.00
-3.50
-20.63
-21.16
-20.03
1.07E+08
5.10E+08
1.19E+08
4.32E+08
1.39E+09
1.66E+08
4.82E+09
3.98E+08
2.23E+08
5.01E+07
6.78E+10
1.57E+10
2.23E+10
1.21E+10
3.63E+09
1.83E+10
9.54E+09
1.34E+09
1.47E+03
1.47E+03
1.61E+09
5.85E+08
1.77E+09
2.33E+09
1.47E+03
1.05E+10
1.70E+10
6.02E+09
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
0.0357
View publication stats
(v pec )
[km s−1 ]
wni
(v pec )
[km s−1 ]
0.0017
0.0022
0.5491
0.2677
0.1613
0.0179
r⊥
[Mpc]
-130.
298.
-185.
89.
-184.
112.
MR
LR
[mag]
[L⊙ ]
uniform weights
0.095
-16.29
1.93E+08
0.283
-16.59
2.54E+08
0.141
-22.56
6.21E+10
0.180
-21.78
3.03E+10
0.326
-21.23
1.83E+10
0.604
-18.84
2.02E+09
wni
0.1667
0.1667
0.1667
0.1667
0.1667
0.1667