# MNRAS LaTeX 2ε template – title goes here

Ivo Saviane,1 Irina Yegorova,2 Dominique Proust3
1 European Southern Observatory, Alonso de Cordova 3107, Vitacura, Casilla
19001, Santiago de Chile 19, Chile.
2 Departamento de Ciencias Fisicas, Universidad Andres Bello, Fernandez Concha
700, Las Condes, Santiago, Chile
3 GEPI, Observatoire de Paris, F92195 Meudon Principal CEDEX, France E-mail:
<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

# The galaxy cluster AC114. II. Stellar populations and the mass-metallicity
relation

Ivo Saviane,1 Irina Yegorova,2 Dominique Proust3
1 European Southern Observatory, Alonso de Cordova 3107, Vitacura, Casilla
19001, Santiago de Chile 19, Chile.
2 Departamento de Ciencias Fisicas, Universidad Andres Bello, Fernandez Concha
700, Las Condes, Santiago, Chile
3 GEPI, Observatoire de Paris, F92195 Meudon Principal CEDEX, France E-mail:
<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

We investigate the mass-metallicity relation for galaxies in the Abell cluster
AC114 from 7 hours of VIMOS/MR data collected at the ESO-VLT telescope in
2009. The dynamical analysis completed in our previous paper allowed us to
select cluster members, whose spectra are here analyzed with stellar
population synthesis models. Active and passive galaxies are identified based
on the presence/absence of the [O ii]$\lambda 3727$, [O iii]$\lambda\lambda
4959,5007$ and/or H$\beta$ emission lines, depending on the galaxy redshift.
We find that active galaxies have lower average masses than passive ones, and
have lower average metallicities. The mass-metallicity relation (MZR) of the
cluster is found to be steeper than that for galaxies in the local universe.

###### keywords:

galaxies: clusters and metallicities - galaxies: distances and redshifts

††pubyear: 2022††pagerange: MNRAS LaTeX 2ε template – title goes here–7

##  1 Introduction

The mass-metallicity relation (MZR) is one of the fundamental constraints to
galaxy evolution, and for this reason it is the subject of a substantial body
of literature (see Maiolino & Mannucci, 2019, for a review), with a section of
that literature looking at its evolution with redshift: indeed, studies of the
MZR of galaxies are now available up to $z\sim 8$ (Jones et al., 2020).

However, metallicities at different redshifts are obtained with different
methods, owing to abundance-sensitive spectral features moving across the
optical-infrared range. This makes a comparison of both absolute and relative
trends across cosmic time rather difficult. Furthermore, spectra of high-
redshift objects cannot be obtained at high resolution, so metallicity indices
calibrated with nearby galaxies must be used. For example in the case of star
forming galaxies, abundance indicators are calibrated based on local H ii
regions, but physical conditions of the interstellar medium (ISM) at high-
redshift might be different than those at present day, thus adding further
uncertainty to the process.

Interestingly, up to $z\sim 0.5$ emission lines that are used to compute
abundances and physical conditions of the ISM, still fall in the optical
range, and sufficiently good spectra can be obtained with 10m-class
telescopes, with reasonable exposure times. Thus a few years ago we identified
the galaxy cluster AC114 as an interesting target to obtain emission-line
spectra of galaxies spanning a range of masses: the cluster sits at $z\sim
0.32$, and as explained in Saviane et al. (2014), we expect that star-forming
galaxies at that epoch are a factor $\sim~{}1.4$ more metal-poor than those in
the local universe. Therefore, obtaining gas-phase metallicities of AC114
cluster members gives us the possibility to check the evolution of the MZR in
the last $\sim~{}4$ Gyr, with abundances obtained with the exact same method
used for local star-forming regions.

The substantial body of work on the cluster has been reviewed in our previous
paper (Proust et al., 2015), nevertheless we summarize here the key features
that make it relevant for the current study. The cluster is classified as
Bautz-Morgan type II-III (Abell et al., 1989) with a galaxy distribution that
tends to be diffuse. Indeed, X-ray emission has an irregular morphology, with
two components dominating the soft part of the spectrum (below 0.5 keV): a
tail extends about 400 kpc from the central emission to the southeast (see De
Filippis et al., 2004). Its compact core is dominated by a cD galaxy and has
strong lensing power with several bright arcs and multiple image sources
(Smail et al., 1995; Natarajan et al., 1998; Campusano et al., 2001). Note
however that it is also classified as a “non-cool core” cluster by Zhang et
al. (2016), meaning that a central supermassive black hole shall not be
present. AC114 has a higher fraction of blue, late-type galaxies compared with
lower redshift clusters (up to 60% outside the core region; Couch et al. 1998;
Sereno et al. 2010, making it a prototypical Butcher-Oemler cluster. Having
many active members is a crucial property when studying the MZR based on H ii
regions.

In the course of data reduction, we realised that abundance-sensitive
absorption features can also be reliably measured in our VIMOS spectra,
therefore a study of the stellar component of AC114 galaxies is presented in
this work, while ISM abundances will be the subject of a forthcoming paper. It
shall be noted that a recent study of stellar populations can also be found in
Rodríguez Del Pino et al. (2014), who targeted thirteen disk galaxies with
integral field spectroscopy. While most galaxies in their sample to not
display emission lines, they still host a young stellar population; this is
not centrally concentrated therefore it is suggested that star formation was
recently truncated by gradual processes such as ram-pressure stripping or weak
galaxy-galaxy interactions.

The present study builds on a dynamical analysis of the velocities of galaxies
in our sample, which was performed in Proust et al. (2015) hereafter referred
to as Paper I. As explained in that paper, thanks to a sample that reaches 1
magnitude fainter ($R\simeq 21$), and which is about 30% larger than those in
previous studies, we could revisit the structure and dynamics of the cluster,
which helped us select cluster members.

###  1.1 Main results from Paper I

The mean redshift obtained in Paper I is $z=0.31665\pm 0.0008$ and the
velocity dispersion is $\sigma=1893_{-82}^{+73}$ km s-1 based on a catalogue
of 524 velocities. The distribution in redshift of all galaxies is shown in
Fig.1 with bins $z=0.01$ widths (magenta histogram). The cluster has a very
elongated main radial filament spanning 12000 km s-1 in redshift space. In
addition, a radial foreground structure was detected within the central 0.5/h
Mpc radius, recognizable as a redshift group at the same central redshift
value. A background structure could also be identified (see also Fig. 2). Two
hundred and sixty five galaxy members were identified, which yield a dynamical
mass $M_{200}=(4.3\pm 0.7)\times 10^{15}$M⊙/h for AC114 and $M_{v}=(5.4\pm
0.7\pm 0.6)\times 10^{15}~{}$M⊙/h from the intrinsic velocity dispersion out
to a radius of 3.98/h Mpc.

###  1.2 Analysis of the stellar populations based on Paper I and new data

Once the redshift is known, the spectrum can be corrected to the rest-frame:
this exercise led to the identification of a number of galaxies whose redshift
value needed revision, mainly because of their low S/N values, in particular
when the R value of the correlation peak is lower or close to 3, except for
emission-line objects (see Sect. 2). A few redshifts could be determined from
the position of the most prominent emission and absorption lines: O ii and CaI
H and K .

The new distribution in redshift of all galaxies in shown in Fig.1 (black
lines) where objects are clearly clustered in a few redshift intervals whose
limits are identified by the vertical lines and horizontal segments. The
limits are listed in Table 1. Galaxies belonging to the cluster are defined as
those within the lower and upper limits of rows 3 and 4 of Table 1: they span
a $\sim$5$\sigma$ range of velocities.

In this paper, we reanalyse the data obtained in Paper I. Section 2 presents
the observations and reduction of the data and section 3 analyses the stellar
populations properties; these properties are used in section 4 to construct
and interpret the mass-metallicity relation, and finally section 5 presents
the conclusions based on the global results obtained for AC114.

Table 1: Redshift limits used to group galaxies. Galaxies belonging to AC114 are defined as those within the lower and upper limits of rows 3 and 4\.  # | $z$ | $c\,z$ | Notes
---|---|---|---
|  | km s-1 |
1 | 0. | 1503 | 45,059 |
2 | 0. | 1900 | 56,961 |
3 | 0. | 2840 | 85,141 | lower limit
4 | 0. | 3310 | 99,231 | upper limit
5 | 0. | 4558 | 136,645 |
6 | 0. | 5487 | 164,496 |
7 | 0. | 6048 | 181,314 |
8 | 0. | 6678 | 200,201 |
Figure 1: Distribution in redshift of all galaxies. The vertical and
horizontal segments mark the limits of redshift ranges that are used later in
the analysis. These limits are consistent with redshift clustering presented
in Paper I, which are based on a larger sample including literature data
(magenta histogram).  Figure 2: Three-dimensional representation of objects in
the sky area covered by our VIMOS observations, with blue dots identifying
active galaxies, and red dots identifying passive ones. AC 114 is the clump of
objects near $cz=100,000\,{\rm km\,s^{-1}}$. The transversal scale has been
greatly amplified to better show structures surrounding the cluster: at the
redshift of the cluster, the 16 arcmin of RA span of VIMOs, convert into 4.8
Mpc linear distance, while the linear distance spanned by the redshift limits
of the cluster is $\sim 300$ Mpc (see Wright, 2006, for the calculations).
Figure 3: A zoom into the three-dimensional representation of AC114. The three
segments show velocity values from 90,000 km s-1 to 100,000 km s-1 in steps of
5,000 km s-1. The transversal scale has been amplified by a factor $\sim 10$
compared to the radial scale, but it is still clear that the cluster is a very
elongated structure. Note that gaps in the object density are artefacts of the
partial sky coverage of the VIMOS footprint.

##  2 Observations and reductions

Table 2: Observations log (OB = Observing Block).

Date | OB starta | Exp. timeb | Filter | Grism
---|---|---|---|---
2009-09-17 | 02:37 | 2250 | GG475 | HR red
2009-09-17 | 03:32 | 2250 | GG475 | HR red
2009-09-17 | 04:11 | 2250 | GG475 | HR red
2009-09-17 | 04:50 | 2250 | GG475 | HR red
2009-09-17 | 05:28 | 2250 | GG475 | HR red
2009-09-17 | 06:07 | 2250 | GG475 | HR red
2009-09-21 | 03:54 | 2250 | GG475 | HR red
2009-08-16 | 05:25 | 2250 | GG475 | MR
2009-08-16 | 06:15 | 2250 | GG475 | MR
2009-08-16 | 06:55 | 2250 | GG475 | MR
2009-08-21 | 07:24 | 2250 | GG475 | MR
2009-08-21 | 06:11 | 2250 | GG475 | MR
2009-09-25 | 04:02 | 2250 | GG475 | MR
2009-09-25 | 04:53 | 2250 | GG475 | MR

a beginning of the observation block, UT time

b exposure time of the science part, in seconds

A complete account of the observations and data reduction is given in paper I,
so just a summary is given here. VIMOS (Le Fèvre et al., 2003) was used to
carry out observations in service mode at the Very Large Telescope, under
program 083.A-0566, between August 16 and September 25, 2009 (see Table 2).
Seven exposures for each grism HR-red and MR were obtained, for a total
shutter time of $\sim 4.4$h for each setup. The MR grism has a spectral
resolution of 580 for a $1^{\prime\prime}$ slit, over the 500 - 1000 nm
spectral range, while the HR-red grism has a spectral resolution of 2500 for a
$1^{\prime\prime}$ slit, over the 630 - 870 nm spectral range.

The galaxy selection was made from the pre-imaging frames of the cluster. To
construct the masks, initially, the SIMBAD catalogue was used without imposing
any restriction criteria. As a first step we identified already known galaxies
of the cluster. Then we selected the non-stellar objects in this region by eye
in order to punch a maximum number of slits in each of the four quadrants.
Such a visual inspection allows to discriminate between extended objects and
stars.

###  2.1 Data reduction

The VIMOS pipeline was used to reduce the data. We reduced separately the data
taken with the grism HR red and MR. Each scientific frame was bias subtracted
and flat field corrected, the cosmic rays were removed, and sky emission lines
subtracted (the same procedure was done for the standard stars). The spectra
were wavelength calibrated and seven frames for each VIMOS quadrant were
combined. This gave us S/N = 20 per pixel on the average at the continuum
level on the final spectra. Since the continuum of most of the galaxies is too
weak we did not use the spectra extracted by the pipeline. Instead, we
extracted one-dimensional spectra with MIDAS. Flux calibration was obtained
with spectrophotometric standards from (Hamuy et al., 1992, 1994).

In this work, only MR data are used, while HR spectroscopy is used in the
forthcoming paper on the nebular abundances, to resolve N ii from H$\alpha$
emission lines.

Apparent magnitudes in the R-band were obtained by computing the integrated
flux of our spectra in the range 4897.5 Å to 9785.5 Å and calibrating it using
objects in common with the superCOSMOS database (Maddox et al., 1990a, b).
Photometry could also have been obtained from pre-imaging frames, but these
are observed under variable sky conditions, so magnitudes would have to be
calibrated on the superCOSMOS system anyway. The precision of our calibration
was estimated in paper I as $\simeq 0.5$ magnitudes, which is comparable with
that of superCOSMOS. Furthermore, synthetic photometry can be obtained with
vastly more efficient means compared to reducing pre-imaging frames from
scratch, and we show in the sections below that luminosities and masses
obtained this way yield well defined mass-metallicity relations. As shown in
Paper I, our galaxy sample reaches luminosities as faint as $R\sim 21$, but it
starts to suffer from incompleteness at $R\sim 20$.

We obtained a total of 163 redshifts, combining and checking those already
obtained in Paper I. The complete redshift values with their individual error
measurements are published in Table 5 and correspond to the highest R-value
obtained from the cross-correlations (Tonry & Davis, 1979). UK-J $B_{j}$ and
UK-J $R_{j}$ magnitudes are from superCOSMOS.

The content of this table is as follows:

  * •

number of the object in each quadrant from slit position;

  * •

right ascension (J2000);

  * •

declination (J2000);

  * •

UK-J $B_{j}$ magnitude from superCOSMOS;

  * •

UK-J $R_{f}$ magnitude from superCOSMOS;

  * •

$R_{comp}$ computed magnitude from spectra;

  * •

redshift;

  * •

redshift error;

  * •

R value from the cross-correlation (Tonry & Davis 1979);

  * •

notes

In the course of data analyses, we identified two objects that were removed
from the table, as their spectra tricked the cross-correlation algorithm,
which returned wrong redshift values; indeed objects Q4/39 and Q4/59 were
identified as M-dwarf stars, after finding that their $D_{n}(4000)$ values
were too high for their supposed mass (see below).

Using data from Table 5, Fig. 2 shows the spatial distribution of all galaxies
in our VIMOS sample with a measured spectrum, with blue and red dots marking
active and passive galaxies respectively. The cluster is clearly visible as
the dense agglomeration of objects near $cz=100,000$ km s-1, and it appears
that most of its galaxies are not actively forming stars, although in the
cluster periphery there are some that do. Couch et al. (2001) note that in
AC114 the star formation rate from the H$\alpha$ emission is not exceeding 4
M⊙ yr-1 and that the H$\alpha$ luminosity function is an order of magnitude
below that observed for field galaxies at the same redshift. As anticipated in
Sec. 1.1, figure 2 also suggests the existence of two other structures, which
appear as filaments extending for tens thousands km s-1, one of them seemingly
connected to the cluster core. Figure 3 shows a zoom into the three-
dimensional representation of the cluster. It is clear that the cluster has a
very elongated structure spanning $\simeq 10,000$ km s-1.

###  2.2 Classification of galaxies

Figure 4: For each redshift range from table 1, this figure shows the co-added
spectra of active galaxies. The vertical blue bars mark the position of the
following emission lines: [S ii] $\lambda\lambda$6716,6731, H$\alpha$, [O iii]
$\lambda\lambda$4959,5007, H$\beta$, H$\gamma$, H$\delta$, [O ii] $\lambda$;
and Ca H and K absorption lines (3968Å, 3934Å). Figure 5: Same as Fig. 4 for
co-added spectra of passive galaxies.

After correcting spectra to rest-frame the 163 galaxies were classified as
‘active’ or ‘passive’ based on the presence/absence of [O ii]$\lambda 3727$,
[O iii]$\lambda\lambda 4959,5007$ and/or H$\beta$ emission lines, depending on
the galaxy redshift, with the following procedure.

The first step consisted in blindly scanning all spectra for the presence of
[O iii]$\lambda 5007$ and/or H$\beta$: a linear fit to the continuum was
subtracted, and a maximum in flux was searched in the window
$4800\leq\lambda\leq 4841$ for H$\beta$ and in the window $4980\leq\lambda\leq
5090$ for [O iii]$\lambda 5007$. A Gaussian fit was then attempted, provided
that the peak in flux is three times larger than the continuum noise (assumed
to be the semi-quartile of the flux variation in the spectral window). To
estimate errors, the fit was repeated for the nominal continuum level, and by
assuming a higher or lower continuum by the amount of noise. The result of
this procedure was that, for each spectrum where a fit is possible, and for
each of the two lines, the central wavelength, its FWHM and error, and the
area of the Gaussian and its error are obtained. After this first screening, a
quality parameter $q$ is assigned to each spectrum, defined as the sum for the
two lines, of their Gaussian area divided by its error. Spectra are then
sorted by fit quality, and plotted one by one to visually check the quality of
the fit. We find that spectra with $q\geq 4$ show unambiguous presence of
emission lines, so galaxies that satisfy this criterion were classified as
‘active’, while the rest were classified as ‘passive’. While performing the
visual checks, a few galaxies preliminarily classified as ‘passive’ were
noticed having a clear [O ii]$\lambda 3727$ emission line, and were added to
the list of ‘active’ objects: their spectra are shown in Fig. 24 of Appendix
A. While some passive galaxies might have very weak emission lines that are
not detected by our methods, the sections below show that the two galaxy
classes have very different properties, thus our classification must be
broadly correct.

Table 7 summarises the results, and adds further remarks on some galaxies,
such as those actually identified as local M-dwarf stars, and those having
low-quality spectra that are not further considered in this paper.

In order to show the redshift evolution of their spectral energy distribution
(SED), Figures 4 and 5 show the co-added spectra of the two galaxy classes, in
the redshift ranges defined in Table 1.

It can be seen that, in the redshift range of AC114, the majority of galaxies
are passive: there are 30 galaxies classified as such (59%), versus 21
classified as active (41%). This confirms the visual impression given by Fig.
2.

##  3 Stellar populations

###  3.1 Representative age and metallicity of active and passive galaxies in
AC114

Figure 6: To select the best simple stellar population (SSP) model that fits
the co-added spectrum of active galaxies in AC114, the dispersion of the
model-to-data ratio is plotted as a function of age and metallicity: this
figure shows that the minimum value is reached for a model age of 2 Gyr, and a
metallicity Z= 0.008.  Figure 7: Same as Fig. 6, for the co-added spectrum of
passive galaxies in AC114. In this case, the best model fit is reached for an
age of 10 Gyr, and a metallicity Z= 0.02.  Figure 8: The co-added spectrum of
all passive galaxies in AC114 redshift range (black) fitted with the best
model of Bruzual & Charlot (2003) in red. This gives the typical age and
metallicity of the underlying stellar population.  Figure 9: Same as figure 8
for active galaxies in the cluster.

As a first step towards the characterisation of galaxies in AC114, we fitted
the co-added spectra obtained above in the redshift range of the cluster, with
stellar population synthesis models from Bruzual & Charlot (2003), following
the prescriptions of Saviane & Jerjen (2007). Briefly, model spectra were
degraded to the resolution of our VIMOS data, and then for each model
spectrum, the steps were the following:

  * •

select the spectral range from 3800 to 6000Å;

  * •

remove emission lines from co-added spectra, before fitting: i.e., remove
regions near H$\beta$ and [O iii] ;

  * •

normalise both model and co-added spectra to maximum flux within the fit
range;

  * •

compute ratio of model to observed spectrum;

  * •

compute average and dispersion of spectral ratio;

  * •

select the best fit model as the one that gives the lowest dispersion of the
spectral ratio.

The results of this procedure are shown in Figures 6 and 7, for active and
passive galaxies, respectively: the typical age of active galaxies in the
cluster is 2 Gyr, and their metallicity is Z = 0.008; while in the case of
passive galaxies, the best fit is reached for an age of 10 Gyr, and a
metallicity Z= 0.02. Galaxies cannot be older than the universe, so it is
reassuring to see that the age of the universe at the redshift of AC114 is
indeed no less than 10 Gyr (see Figures 15 and 16). The best-fit models are
shown overlaid to the co-added spectra in figures 8 and 9. Figures 6 and 7
show that adopting model ages and metallicities closest to the best fits,
would also give good results: this uncertainty is included in the calculation
of M/L values in Sect. 3.3.

This exercise reveals that passive galaxies must have formed the bulk of their
population early on, in a strong burst that quickly enriched their ISM to
solar metallicities: indeed, the remarkably good fit of a simple stellar
population model (SSP) indicates that those stars were formed in a relatively
short time span. On the contrary, the low stellar metallicity of active
galaxies is an indication that star formation (SF) proceeded at a low pace in
these systems, which are indeed still forming stars. Still, the good fit of a
SSP to the observed spectrum is telling us that most stars in active galaxies
were also formed rapidly, although later than those in passive galaxies.

Passive galaxies are in general more massive than active ones (see below), but
physical sizes do no vary too much between the two classes; therefore one
expects a higher gas density in passive galaxies, which would explain their
higher SFR, according to the Schmidt-Kennicutt law (Kennicutt, 1998). Figures
2 and 3 show that passive galaxies are concentrated near the cluster centre,
where mass must have been accumulating in the early formation epochs, thus
easing the formation of the largest objects (see also Fig. 6 in Proust et al.
2015). In addition, SF in the densest regions of galaxy clusters is quenched
more effectively than in their outskirts (see, e.g., Deshev et al., 2017):
therefore the active/passive dichotomy can also be explained, at least
partially, by an effect of the environment.

###  3.2 Comparison with Couch & Sharples (1987)

Figure 10: Colour-magnitude diagram for AC114 galaxies with photometry in
superCOSMOS database (Maddox et al., 1990b). The solid line is the red
sequence as defined in Paper I, and the horizontal dotted line represents the
separation between red and blue galaxies as in CS87.

To confirm the results above, it is instructive to compare them to Couch &
Sharples (1987, hereafter CS87). They classify their galaxy sample into ‘red’
and ‘blue’ objects, based on their position in the colour-magnitude plot,
which is shown in Fig. 10 for the subset of galaxies having published colours
in superCOSMOS. In CS87, blue galaxies are defined has having colours smaller
than $B_{{\rm J}}-R_{{\rm F}}=2$, which is represented by the dotted line in
the figure; the solid line represents the ‘red sequence’ as defined in paper
I, which is similar to the one found by CS87. Passive and active galaxies are
plotted with red or blue symbols, so the figure demonstrates that most active
galaxies would be classified as ‘blue’ by CS87, while passive galaxies gather
around the red sequence.

Most red galaxies in CS87 have spectra comparable to nearby E/S0 galaxies, and
indeed their comparison to models of population synthesis yields a typical age
of 10 Gyr. They also find 15% of red galaxies having strong H$\delta$
absorption, which are interpreted as hosting a short 0.5 Gyr burst with an age
of 1.5 Gyr. These results are consistent with our findings above, where the
typical age of passive galaxies was indeed found to be 10 Gyr, with a few
exceptions identified in Sec. 3.5.3. This age is used in Sect. 3.3 to compute
M/L ratios, therefore we can expect that masses of passive galaxies are on the
average correct: indeed the tight sequence found in Sect. 3.5.2 and Fig. 17
validate this conclusion.

As expected from their spread in colour, blue galaxies in CS87 have a more
diverse SF history and interestingly, almost half of them are pure absorption-
line objects. By comparison to local galaxies and models, their emission line
objects are split between having spectra comparable to nearby spiral galaxies
(their type 2) and hosting a current SF burst (their type 1). Finally, blue
galaxies with no emission lines are interpreted as having experienced a SF
burst of varying strength, up to 1.5 Gyr prior to the epoch of observation.
The stacked spectrum of our active galaxies, which by definition are emission-
line objects, must be a mix of type 1 and type 2 of CS87, therefore the age of
2 Gyr found above is a reasonable guess. Once more, this is reinforced by Fig.
17, which shows a well-defined trend of metallicity index versus masses. Note
also that some objects do require a small correction to their ages, as
discussed in Sect. 3.5.3.

Only four galaxies are in common between this work and CS87, and they all
belong to the red sequence, with normal H$\delta$ absorption. A larger overlap
in the two sets would be needed for a thorough comparison, nevertheless in
Appendix C we show that our spectra match the classification of CS87.

In the context of this discussion, it is also interesting to examine ‘postage
stamps’ of our targets as shown in figures 25 and 26, for passive and active
galaxies, respectively. As expected for an E/SO type, most passive galaxies
have regular shapes, but there are some galaxies that appear to be interacting
(at least in projection) with nearby smaller objects (see image caption). It
is interesting to note that two of these potentially interacting galaxies (44
and 49) present an anomalous $D_{n}(4000)$ index, as discussed in Section
3.5.3, which means that their ages deviate from the typical one of their
class. Conversely, galaxies 64 and 21 also have a deviant $D_{n}(4000)$ index,
but they appear to be isolated in Fig. 25. Active galaxies appear more diffuse
than passive ones, and contrary to passive ones, objects with anomalous
$D_{n}(4000)$ index do not have close companions but rather appear irregular
in shape (numbers 36, 56, 48, 40, and 26). The implications of galaxy
morphology will be examined in more detail in a forthcoming paper where
nebular abundances will be obtained (Andrade et al., in preparation), and
contrasted with the stellar abundances obtained in this work.

###  3.3 Mass-to-Light ratios

Figure 11: V-band mass-to-light ratio as a function of the age of the
population, from BC2003: the band represents uncertainties introduced by the
choice of different initial IMF, spectral library, theoretical evolutionary
tracks, and metallicities. Based on the age returned by the best-fit models,
representative M/L ratios can be estimated for passive and active galaxies,
which are shown by the dots and error bars. Arrows in the lower right corner
show that, for ages greater than $\sim 0.1$ Gyr, a change in age of 0.5 dex
translates into a similar 0.5 dex change in M/L ratio.

Having in hand the age of the stellar populations, the next step is the
computation of their mass-to-light ratios. From the plots in BC03 we obtained
Fig. 11, which shows M/LV vs. age: as discussed in BC03, the mass-to-light
ratio is affected by the choice of the initial mass function (IMF), spectral
library, theoretical evolutionary tracks, and metallicity, so the shaded band
in the figure represents those uncertainties, estimated as $1/4$ of the full
variation interval.

As shown in Table 2, the most complete photometric catalogue for our galaxies
is in the R-band, therefore mass-to-light ratios were converted from V- to
R-band using the expression: $M/L_{R}=M/L_{V}\times 10^{-0.4\,(V-R)}$, where
the $V-R$ colour of an old population was taken as the average of globular
clusters in the Harris (1996) catalogue: $V-R=0.64\pm 0.17$. The resulting
ratios are then $M/L_{R}=0.64\pm 0.26$ and $M/L_{R}=2.41\pm 1.12$ for active
and passive galaxies, respectively.

###  3.4 Galaxy distances, luminosities, and masses

Figure 12: Luminosity distance, expressed in Mpc, vs. redshift: the curve is a
quadratic interpolation of selected values calculated using the cosmology
calculator of Wright (2006).  Figure 13: Comparing absolute $R$ magnitudes vs.
integrated flux, for passive galaxies in AC114 redshift range. As expected,
the two quantities are proportional to each other, but the correlation breaks
down at the lower end of total fluxes. An uncertainty of 0.5 magnitudes has
been assumed both for the total flux and the R-band photometry.  Figure 14:
Same as Fig. 13, for active galaxies in AC114 redshift range.  Figure 15: Mass
of passive galaxies as a function of redshift and age of the universe. The
dotted line represents a linear fit through the data. Errors in mass are taken
from Table 3, while errors in redshift are smaller than the symbols.  Figure
16: Mass of active galaxies as a function of redshift. The dotted line
represents a linear fit through the data, while the contours represent the
general distribution of passive galaxies in the same plot of Fig. 15. Errors
in mass are taken from Table 3, while errors in redshift are smaller than the
symbols.

To compute galaxy luminosities, apparent magnitudes were converted to absolute
ones with the usual $M_{R}=R-\mu$, where the distance modulus is $\mu=5\,\log
D_{L}-5$. The luminosity distance $D_{L}$ is a function of redshift, and it
was evaluated following Wright (2006), with input parameters $H_{0}=69.6$,
$\Omega_{{\rm matter}}=0.286$, and $\Omega_{{\rm vac}}=0.714$ (flat universe).
After computing the relation for a few selected values, it was interpolated by
a quadratic function, as shown in Fig. 12. We find:

$D_{L}=2245.7\,z^{2}+4628.4\,z-20.838$

which is valid for $0.05\leq z\leq 0.7$. At the mean redshift quoted in
section 1.1, the distance is $D_{L}=~{}1670\pm 17$ Mpc

Absolute luminosities were then computed as
$L_{R}=10^{-0.4\,(M_{R}-M_{R_{\odot}})}$, with $M_{R_{\odot}}=4.5$. Finally,
luminosities were converted into masses with the mass-to-light ratios computed
in Sect. 3.3. Uncertainties on mass values were computed by adding in
quadrature uncertainties from photometry and from mass-to-light ratios.

As a consistency check between photometric and spectral properties, the
integrated spectral flux is compared to absolute $R$ magnitudes in Figures 13
and 14 for the two classes of galaxies. It is reassuring to see that there is
a good correlation between the two quantities: the correlation coefficient is
$-0.94$, with a dispersion of $\sim 0.24$ magnitudes for passive galaxies, and
the same quantities are $-0.85$ and $\sim 0.4$ magnitudes for active ones.
However the correlation breaks down for galaxies at the faint end of the
luminosity range: therefore for the next analysis we retained only objects
with $\Sigma({\rm flux)>10^{-15}\,{\rm erg\,cm^{-2}sec^{-1}}}$.

Figure 15 shows that passive galaxies in AC114 span an order of magnitude in
mass; interestingly, looser cluster members at higher and lower redshift tend
to define a mild tendency of having masses increasing while moving to more
recent times: from 9.8 to 10.3 Gyr ago, their mass increases from $0.4\times
10^{11}\,M_{\odot}$ to $10^{11}M_{\odot}$. It could be inferred that some
galaxies continued forming stars both before and after a cluster-wide main SF
episode.

On the contrary, Fig. 16 shows that active galaxies closer to us have lower
masses than objects located in the main body of the cluster, which also span a
smaller mass range of $\sim 0.5$ dex, and are generally less massive than
passive galaxies in the same region. This confirms the idea that active
galaxies formed later than passive ones, and did not participate to the
initial strong galaxy formation epoch.

###  3.5 Metallicities

Table 3: Main characteristics of passive galaxies entering the mass-metallicity relation. Luminosities and masses are given in $10^{10}$ solar units. Galaxies are sorted by luminosity, from fainter to brighter. Typical error on $R$ magnitudes is 0.5, which yields $\sim 50\%$ errors on luminosities and masses, after accounting for the uncertainty in mass-to-light ratios.  Q#-sl# | RA | DEC | $z$ | $R$ | $M_{R}$ | $L_{R}$ |  $\mathcal{M}/\mathcal{M_{\odot}}$ $\pm$ |  $D_{n}(4000)$ $\pm$ | [m/H] $\pm$
---|---|---|---|---|---|---|---|---|---
Q4-45 | 344.652 | -34.781 | 0.31390 | 20.600 | -20.490 | 0.990 | 2.390 | 1.110 | 1.589 | 0.099 | -0.668 | 0.308
Q4-23 | 344.646 | -34.743 | 0.32270 | 20.600 | -20.560 | 1.060 | 2.550 | 1.180 | 1.387 | 0.095 | -1.355 | 0.402
Q4-24 | 344.638 | -34.744 | 0.31450 | 20.500 | -20.600 | 1.090 | 2.630 | 1.220 | 1.438 | 0.076 | -1.160 | 0.296
Q4-16 | 344.623 | -34.732 | 0.31850 | 20.400 | -20.730 | 1.230 | 2.970 | 1.380 | 1.660 | 0.121 | -0.481 | 0.336
Q4-21 | 344.711 | -34.738 | 0.30610 | 19.500 | -21.530 | 2.580 | 3.117 | 1.448 | 1.402 | 0.086 | -1.022 | 0.445
Q3-26 | 344.618 | -34.614 | 0.31480 | 20.200 | -20.900 | 1.440 | 3.480 | 1.620 | 1.494 | 0.096 | -0.962 | 0.349
Q4-43 | 344.620 | -34.778 | 0.31720 | 20.200 | -20.920 | 1.470 | 3.540 | 1.650 | 1.394 | 0.084 | -1.327 | 0.350
Q4-63 | 344.719 | -34.811 | 0.31870 | 20.200 | -20.930 | 1.490 | 3.580 | 1.660 | 1.386 | 0.082 | -1.359 | 0.345
Q4-28 | 344.612 | -34.751 | 0.31470 | 20.000 | -21.100 | 1.730 | 4.180 | 1.940 | 1.592 | 0.090 | -0.659 | 0.276
Q2-35 | 344.919 | -34.651 | 0.31290 | 19.700 | -21.380 | 2.260 | 5.440 | 2.530 | 1.642 | 0.092 | -0.525 | 0.257
Q4-41 | 344.616 | -34.774 | 0.31820 | 19.700 | -21.430 | 2.350 | 5.650 | 2.630 | 1.626 | 0.081 | -0.567 | 0.231
Q4-20 | 344.664 | -34.739 | 0.30860 | 19.600 | -21.450 | 2.400 | 5.780 | 2.690 | 1.704 | 0.099 | -0.379 | 0.244
Q4-2 | 344.646 | -34.707 | 0.31980 | 19.500 | -21.640 | 2.850 | 6.880 | 3.200 | 1.581 | 0.074 | -0.690 | 0.228
Q4-38 | 344.618 | -34.767 | 0.30100 | 19.200 | -21.790 | 3.270 | 7.890 | 3.670 | 1.576 | 0.136 | -0.705 | 0.447
Q4-47 | 344.645 | -34.785 | 0.31670 | 19.300 | -21.810 | 3.350 | 8.080 | 3.760 | 1.624 | 0.092 | -0.572 | 0.266
Q3-40 | 344.642 | -34.648 | 0.31830 | 19.300 | -21.830 | 3.390 | 8.180 | 3.800 | 1.753 | 0.090 | -0.278 | 0.195
Q4-58 | 344.674 | -34.802 | 0.31080 | 19.200 | -21.870 | 3.520 | 8.490 | 3.940 | 1.694 | 0.096 | -0.401 | 0.241
Q4-49 | 344.665 | -34.788 | 0.31820 | 20.000 | -21.130 | 1.780 | 8.560 | 3.971 | 1.895 | 0.164 | -0.185 | 0.255
Q4-51 | 344.693 | -34.792 | 0.30950 | 19.100 | -21.960 | 3.830 | 9.220 | 4.280 | 1.748 | 0.151 | -0.288 | 0.357
Q4-31 | 344.652 | -34.755 | 0.31780 | 19.100 | -22.020 | 4.060 | 9.800 | 4.550 | 1.825 | 0.094 | -0.155 | 0.166
Q4-32 | 344.658 | -34.757 | 0.31710 | 18.200 | -22.920 | 9.260 | 22.330 | 10.380 | 1.947 | 0.103 | -0.014 | 0.114
Q4-52 | 344.609 | -34.794 | 0.31790 | 18.200 | -22.920 | 9.320 | 22.460 | 10.440 | 1.925 | 0.104 | -0.033 | 0.129
Q4-64 | 344.631 | -34.815 | 0.31450 | 18.300 | -22.800 | 8.290 | 39.865 | 18.536 | 2.077 | 0.136 | -0.072 | 0.067
Q4-44 | 344.671 | -34.779 | 0.31520 | 18.300 | -22.800 | 8.330 | 40.065 | 18.616 | 2.063 | 0.133 | -0.075 | 0.075
Table 4: Main characteristics of active galaxies entering the mass-metallicity relation. Luminosities and masses are given in $10^{10}$ solar units. Galaxie are sorted by luminosity, from fainter to brighter. Typical error on $R$ magnitudes is 0.5, which yields $\sim 50\%$ errors on luminosities and masses, after accounting for the uncertainty in mass-to-light ratios. Q#-sl# | RA | DEC | $z$ | $R$ | $M_{R}$ | $L_{R}$ |  $\mathcal{M}/\mathcal{M_{\odot}}$ $\pm$ |  $D_{n}(4000)$ $\pm$ | [m/H] $\pm$
---|---|---|---|---|---|---|---|---|---
Q4-15 | 344.608 | -34.731 | 0.30800 | 20.900 | -20.140 | 0.720 | 0.460 | 0.190 | 1.117 | 0.063 | -2.497 | 0.399
Q4-14 | 344.673 | -34.730 | 0.30460 | 20.700 | -20.320 | 0.840 | 0.540 | 0.220 | 0.931 | 0.071 | -3.674 | 0.449
Q2-38 | 344.846 | -34.663 | 0.29870 | 20.100 | -20.870 | 1.400 | 0.900 | 0.370 | 1.013 | 0.053 | -3.155 | 0.335
Q4-19 | 344.609 | -34.737 | 0.31070 | 20.200 | -20.870 | 1.400 | 0.900 | 0.360 | 0.996 | 0.048 | -3.262 | 0.304
Q4-46 | 344.712 | -34.784 | 0.32950 | 20.100 | -21.110 | 1.760 | 1.130 | 0.460 | 1.155 | 0.056 | -2.256 | 0.354
Q1-18 | 344.846 | -34.751 | 0.31300 | 19.700 | -21.380 | 2.260 | 1.450 | 0.590 | 1.137 | 0.073 | -2.370 | 0.462
Q3-1 | 344.673 | -34.555 | 0.32950 | 19.800 | -21.410 | 2.320 | 1.480 | 0.600 | 1.175 | 0.063 | -2.129 | 0.399
Q1-37 | 344.849 | -34.791 | 0.31600 | 19.600 | -21.510 | 2.530 | 1.620 | 0.660 | 1.290 | 0.056 | -1.402 | 0.354
Q4-26 | 344.675 | -34.748 | 0.32460 | 18.900 | -22.280 | 5.130 | 1.644 | 1.330 | 1.071 | 0.031 | -2.410 | 0.230
Q1-30 | 344.901 | -34.779 | 0.31860 | 19.600 | -21.530 | 2.580 | 1.650 | 0.670 | 1.251 | 0.066 | -1.648 | 0.418
Q2-32 | 344.854 | -34.643 | 0.30910 | 19.500 | -21.550 | 2.640 | 1.690 | 0.690 | 1.174 | 0.055 | -2.136 | 0.348
Q4-22 | 344.666 | -34.741 | 0.31450 | 19.400 | -21.700 | 3.010 | 1.930 | 0.780 | 1.137 | 0.059 | -2.370 | 0.373
Q4-48 | 344.621 | -34.787 | 0.31460 | 20.000 | -21.100 | 1.730 | 2.215 | 0.450 | 1.470 | 0.088 | -0.726 | 0.392
Q3-12 | 344.654 | -34.583 | 0.30940 | 19.000 | -22.060 | 4.190 | 2.680 | 1.090 | 1.227 | 0.051 | -1.800 | 0.323
Q4-40 | 344.604 | -34.772 | 0.31780 | 19.000 | -22.120 | 4.460 | 2.850 | 1.160 | 1.225 | 0.037 | -1.813 | 0.234
Q4-60 | 344.685 | -34.806 | 0.31830 | 19.000 | -22.130 | 4.470 | 2.860 | 1.160 | 1.224 | 0.039 | -1.819 | 0.247
Q1-40 | 344.848 | -34.799 | 0.32280 | 18.100 | -23.060 | 10.580 | 3.393 | 2.750 | 1.081 | 0.047 | -2.335 | 0.348
Q4-36 | 344.650 | -34.764 | 0.31310 | 18.400 | -22.690 | 7.480 | 9.557 | 1.950 | 1.757 | 0.094 | -0.025 | 0.132
Q4-56 | 344.684 | -34.800 | 0.31380 | 18.400 | -22.690 | 7.520 | 9.597 | 1.960 | 1.743 | 0.095 | -0.039 | 0.149

####  3.5.1 Introduction of the D4000 index

Having measured the masses of AC114 galaxies in the previous sections,
metallicities were then estimated using the $D_{n}(4000)$ index, which
measures the ratio of spectral continuum at the red and blue side of the 4000Å
discontinuity. As in Balogh et al. (1999), the average of the continuum and
its dispersion were calculated within the bands 3850–3950Å and 4000–4100Å,
both for the galaxy sample and the models. If we call $\overline{F_{{\rm
red}}}$ and $\overline{F_{{\rm blue}}}$ the average continuum level in the two
bands computed over $n_{{\rm pix}}$ spectral elements, and $\sigma_{{\rm
red}}$ and $\sigma_{{\rm blue}}$ their dispersion, then $D_{n}(4000)$ and its
uncertainty are computed as ${\rm D_{n}(4000)}=\overline{{\rm
F_{red}}}/\overline{{\rm F_{blue}}}$ and $\sigma_{{\rm
D_{n}(4000)}}=D_{n}(4000)\times(\sigma_{{\rm F,blue}}/\overline{{\rm
F_{blue}}}+\sigma_{{\rm F,red}}/\overline{{\rm F_{red}}})$, where
$\sigma_{{\rm{\rm F}}}=\sigma_{{\rm px}}/\sqrt{n_{{\rm px}}}$.

####  3.5.2 $D_{n}(4000)$ vs. mass for galaxy spectra

Figure 17: Relationship between $D_{n}(4000)$ and galaxy mass, for passive
(red dots) and active (blue dots) galaxies. The straight lines and shaded
areas represent weighted linear fits to the data, and the dispersion around
the fit. The top panel shows the original measurements, with the most evident
outliers marked with large circles. To these objects, a correction in mass has
been applied, as explained in the text, and their new position can be seen in
the bottom panel. Their spectra are plotted in Fig. 18.

Figure 17 shows the results of our measurements, with $D_{n}(4000)$ plotted as
a function of galaxy mass. It appears that passive and active galaxies occupy
distinct places in such diagram, with active galaxies generally with lower
masses and smaller $D_{n}(4000)$ values, compared with passive galaxies. In
the hypotheses that galaxies within a class have a common age (Sec. 3.1), the
variation of the index must be mainly due to metallicity differences, thus the
graph can be interpreted as passive galaxies having higher metallicities than
active ones.

Figure 17 also shows that some objects lay outside the general trend defined
by galaxies of comparable mass, which can be interpreted by using Fig. 19: it
is likely that galaxies with larger/smaller values of $D_{n}(4000)$ have
typical ages that are larger/smaller than the typical age inferred in Sec.
3.1.

####  3.5.3 Special cases of of $D_{n}(4000)$

Figure 18: In the left two panels, spectra of passive galaxies having
$D_{n}(4000)$ values that are larger or smaller than those of galaxies of
comparable mass, are plotted in the left-hand panels, with index values
increasing from bottom to top. The right-hand panels plot spectra of ‘deviant’
active galaxies. Within each panel, the $D_{n}(4000)$ value is printed in the
top left corner, following the galaxy identification. In each spectrum, the
blue and red bands highlight the wavelength range that is used to compute the
ratio of continuum near the 4000 Å break. In the right two panels, images of
the same galaxies, extracted from the VIMOS preimaging are displayed.

An inspection of Fig. 18 shows that such interpretation must be correct:
active galaxies with large $D_{n}(4000)$ values have red SEDs which resemble
those of passive galaxies of similar index values, thus having ages greater
than the typical 2 Gyr. Furthermore, their spectra have very weak [O
ii]$\lambda 3727$ and H$\alpha$ emission lines, which indicates a low SFR. At
the opposite side, active galaxies with relatively low index values, have blue
SEDs and very prominent emission lines, indicative of an age younger than 2
Gyr. For these galaxies, a second possibility is a metallicity which is
intrinsically lower than galaxies of comparable mass, for example due to
accretion of fresh gas: indeed the same Fig. 18 shows a low-luminosity tail
emerging from object Q4/26, so this option cannot be excluded.

As Fig. 11 demonstrates, an age increase/decrease translates into a
larger/smaller M/L ratio, with a change of 0.1 dex in M/L for every change of
0.1 dex in age. Thus, if galaxies with large $D_{n}(4000)$ are older than the
others, their luminosities must be converted into larger masses: from Fig. 17,
it can be evinced that corrections of $\sim\pm 0.3$ dex in mass, i.e. in M/L,
would bring deviant galaxies back into the general trend. Such corrections are
within the uncertainties of the measurements.

Such a change would bring the age of active galaxies with large $D_{n}(4000)$
close to that of passive galaxies, in agreement with the conclusion inferred
from their spectra. At the other side, active galaxies with a low
discontinuity index would re-enter the general trend, if their ages were $\sim
1$ Gyr.

Looking at passive galaxies, object Q4/49 has a $D_{n}(4000)$ value and SED
close to that of objects Q4/44 and Q4/64: again, a correction of $\sim+0.3$
dex in M/L, corresponding to a correction of $\sim+0.3$ dex in age, would
bring its mass in agreement with the general trend. Conversely, galaxies with
relatively low index values have bluer SEDs than galaxies with masses
initially estimated to be the same: a reduction in age of a few Gyr would thus
move them to the left in the figure, and reconcile $D_{n}(4000)$ measurements
with estimated masses.

The lower panel of Fig. 17 shows how the mass and hence age corrections
proposed above improve the definition of the loci occupied by the two classes
of galaxies.

####  3.5.4 $D_{n}(4000)$ for model spectra

Figure 19: Relationship between $D_{n}(4000)$ and [m/H] from the BC03 model
spectra, for ages 1, 2, 5, 10, and 20 Gyr, going from left to right. The
shaded areas represent linear (for ages $\leq 2$ Gyr) and quadratic fits to
the points, and the dispersion around the fits. The vertical blue dotted
segment marks the smaller value of $D_{n}(4000)$ for active galaxies, while
the vertical red dotted segments mark the range of $D_{n}(4000)$ for passive
galaxies.

To convert $D_{n}(4000)$ values to metallicities, a calibration of the index
was obtained from BC03 models, as shown in Fig. 19. It can be seen that a
given value of $D_{n}(4000)$ can correspond to a range of metallicities,
depending on the age of the population: for most active and passive galaxies,
the calibration for ages 2 Gyr and 10 Gyr were used. Calibrations for 5 Gyr
and 20 Gyr (in a relative extrapolation) were used for passive galaxies
needing correction to their masses as explained in Sect. 3.5.3, while for
active galaxies calibrations for 1 Gyr and 5 Gyr were used.

The figure also shows that $D_{n}(4000)$ measured for active galaxies reaches
values lower than those spanned by theoretical models. Although the linear
extrapolation looks relatively safe, metallicities lower than [m/H]$\sim-2.5$
must be regarded as uncertain.

Metallicities can now be added to the main parameters of passive and active
galaxies in AC114, which are summarised in Tables 3 and 4. It is also useful
to be able to inspect the morphology of each object, therefore Figs. 25 and 26
in Appendix B show images of the same galaxies, extracted from the VIMOS pre-
imaging.

##  4 Discussion

Taking the key quantities from the aforementioned tables, we can now use them
to define the evolutionary status of galaxies in AC 114. In particular, we can
construct the mass-metallicity relation, and interpret it by way of a simple
model of chemical evolution: it will tell us that both the mass of a galaxy,
and its location within the cluster, are key drivers of evolution.

###  4.1 Mass-metallicity relation

Figure 20: Illustration of different ways to obtain an MZR, with boxes
representing two galaxies of different mass, and arrows connecting the start
and end points of their evolution, assumed to happen in the same time
interval. Light blue is pristine gas, darker blue is enriched gas, and red
represents metals. In case (A), two galaxies of different mass produce the
same fraction of metals over pristine gas, so no MZR is generated. In case
(B), the more massive galaxy generates a higher fraction of metals. In case
(C), both galaxies produce the same fraction of metals, but the low-mass
objects ejects more of them into the IGM. Finally in case (D), gas is lost
before chemical evolution is complete, in such a way that the less massive
galaxy will have generated a smaller fraction of metals, compared with the
more massive one.

As subsequent stellar generations are created within a galaxy, they enrich the
ISM and build the stellar mass at the expense of gas mass. It is therefore
expected that mass and metallicity grow in time, and if galaxies were able to
perform the gas conversion at the same rate, isolated galaxies of the same age
would have converted the same fraction of gas into stars, and thus would have
the same metallicity, irrespective of mass. But galaxies show an ubiquitous
mass-metallicity relation, as testified by a vast literature on the subject
(see, e.g. Maiolino & Mannucci, 2019, for a review). This fact can have
several explanations, which are illustrated in Fig. 20: galaxies of larger
mass might be able to convert larger fractions of gas into stars, in the same
amount of time (higher SF efficiency); their stellar generations might be able
to produce more metals than those in low-mass galaxies (higher yield); metals
produced by low-mass galaxies might be more easily lost into the IGM; or
finally, low-mass galaxies might have been originally higher mass objects,
whose SF was truncated due to loss of residual gas into the IGM. It is then
interesting to investigate the existence of a MZR of AC114 galaxies, and see
what it can tell us about the evolutionary history of the cluster.

####  4.1.1 The MZR and its Interpretation

Figure 21: The top panel show the location of passive galaxies (red symbols)
and active galaxies (blue symbols) in the [m/H] vs. stellar mass plane. The
tracks are close-box models as explained in the text: for passive galaxies the
tracks just reach their representative point, while for active galaxies we
also plot the expected evolution until close to terminal gas fraction
$\mu=0.001$. Tracks are not plotted for active galaxies located in the passive
galaxy region, which are identified by their quadrant-slit combination. The
magenta points on active tracks represent the expected evolution until the
present time for a linear $\mu(t)$ function set by past evolution. The lower
panel shows the calculated $\mu$ as a function of the galaxy mass and
metallicity, assuming the yield values displayed in the plot. The grey area
shows the typical value of $\mu$ for local galaxies, taken from Fig. 22.
Magenta segments represent the expected evolution until the present time. .
Figure 22: Gas mas fraction distribution for local spiral galaxies, with data
from McGaugh & de Blok (1997). The curve is a Gaussian with average and
dispersion as displayed in the plot.

When $D_{n}(4000)$ is converted into a metallicity, Fig. 17 is converted into
Fig. 21, which shows the location of AC114 galaxies in the [m/H] vs. mass
plane. As expected, the trend of $D_{n}(4000)$ with mass converts into a mass-
metallicity relation, for the two galaxy classes. To interpret the MZR, we
make use of the so-called closed-box model (Searle & Sargent, 1972), which,
despite its simplicity, allows to draw interesting conclusions on the
evolutionary status of the galaxy cluster.

In the close box approximation, the evolution of a galaxy is parameterised by
the gas mass fraction $\mu=m_{{\rm gas}}/m_{{\rm tot}}=m_{{\rm gas}}/(m_{{\rm
gas}}+m_{{\rm stars}})$: a galaxy starts as a mass of pristine gas ($\mu=1$),
which is converted into stars thereby producing metals, until $\mu=0$. All
along, the total mass remains constant. In the same framework, according to
Pagel (1997) the average abundance of a stellar population starts at zero and
it evolves with $\mu$ according to the equation:

$<z>=1+\frac{\mu\,\ln\mu}{1-\mu}$ (1)

where the metallicity is expressed as $z=Z/p$, and $p$ is the metal yield of a
stellar generation111The yield of metal $i$ is given by
$p_{i}=[\int_{m(t)}^{m_{U}}m\,q_{i}(m)\,\phi(m)\,dm]/S$, where $q_{i}(m)$ is
the yield per stellar mass, and $\phi(m)$ is the initial mass function; the
global yield is then $p=\Sigma p_{i}$. $S$ is the mass in stars and remnants
of each generation, such that $p\times S$ is the mass of metals returned to
the ISM.: therefore $<Z>\rightarrow p$ when $\mu\rightarrow 0$. At the end of
its SF history a galaxy will then have an average metallicity similar to the
yield, which can then be estimated in this way. Thus, if passive galaxies have
reached the end of their SF history, then Fig. 21 is telling us that their
yield depends on galaxy mass. And because the stellar yield only depends on
the IMF, what must be changing is the ability of a galaxy to return its metals
to the ISM: indeed a common explanation for the MZR is that a less massive
object has a shallower gravitational potential from which metals can escape
more easily (e.g., Dekel & Silk, 1986).

If instead the yield is constant, then $\mu$ must depend on the galaxy mass:
this is shown in the lower panel of Fig. 21, where $\mu$ has been calculated
assuming that the constant yield is close to that of the most metal rich
galaxies, i.e. equal to $Z_{\odot}$. The plot shows that larger galaxies would
have been able to convert a higher fraction of their initial mass into stars.
Because passive galaxies do not have associated large gas masses, the
unprocessed gas must have been lost into the IGM, thus truncating their SF:
again, the gas dispersal would be easier for less massive objects. Of course
this would be a stochastic process which could be also influenced by the
galaxy location in a cluster and a MZR might be not expected, but the close-
box tracks of Fig. 21 run almost parallel to the MZR, so even a random
truncation of the SF would naturally lead to a relationship. Therefore this is
our preferred explanation for the distribution of passive galaxies in Fig. 21.

If active galaxies had the same yield as passive ones, their low metallicities
would be explained by a very slow chemical evolution: the close-box model
predicts that only 2% of the initial gas mass needs to be processed
($\mu=0.98$), to reach a metallicity [m/H]=$-2$. The age of the cluster is
$\sim 10$ Gyr, so if active galaxies continued their evolution at the same
pace, today ($\sim 3.7$ Gyr later) they would have reached [m/H]=$-1.86$, and
consumed only an additional 1% of gas mass.

This means that we should observe a mass of gas associated to active galaxies
$\sim 50$ times the mass of stars, and that their total masses would be more
than an order of magnitude larger than those of passive galaxies. Both
deductions are not in agreement with observations, so we must conclude that
active galaxies have a lower yield compared to passive ones. Indeed this is
also the conclusion of Köppen et al. (2007), based on the idea that stars in
less massive galaxies are generated in less massive stellar clusters: these
are not able to form massive stars, thereby skewing the IMF to lower masses
and reducing yields. More massive galaxies can produce more massive clusters
because they compress their gas more efficiently, as predicted by the Schmidt-
Kennicutt law (Schmidt, 1959; Kennicutt, 1998).

To put some constraints to low-mass galaxy yields, we imposed that the total
masses of active galaxies should not be larger than those of passive ones:
thus we plotted the track that reaches the higher mass at end of evolution,
and changed the yield until its total mass was similar to that of the most
massive passive galaxy. The result is $p=0.03\,Z_{\odot}$, which we adopted
for all active galaxies: this value is well within the yield reduction
computed by Köppen et al. (2007).

As for passive galaxies, the location of active ones along their tracks
constrains their mass fractions, which are plotted in the lower panel of Fig.
21. But because these objects are still evolving, we can expect that nowadays
they will have reached higher masses and metallicities and lower $\mu$ values:
therefore we also computed the expected evolution from an age of the universe
$t_{10}\simeq 10$ Gyr to the present time, assuming a linear dependence of
$\mu(t)=1-(1-\mu_{10})/t_{10}\times t$. This evolution is shown by the magenta
symbols and magenta segments in the upper and lower panels of Fig. 21,
respectively.

The lower panel of the figure also shows the $\pm 1\,\sigma$ range of gas
fractions in local spiral galaxies, taken from McGaugh & de Blok (1997) (see
Fig. 22). The comparison tells us that, at the current pace of evolution, most
active galaxies in AC114 will not be able to reach the same level of gas
processing as galaxies in the local universe.

Some active galaxies in the upper panel of Fig. 21 are located in the region
of passive galaxies, so a comparable yield was adopted to compute their
evolutionary tracks. These are galaxies Q1-30, Q1-37, Q4-36, Q4-48, and Q4-56:
from Fig. 26, it appears that the first three objects are relatively large
galaxies with a disk/bulge morphology, therefore their high $D_{n}(4000)$
values might be due to the fact that our VIMOS slits include a large fraction
of the central, older and more metal-rich regions. To some extent this also
applies to Q4-48, while the case of Q4-36 requires more investigation, because
its classification as a high luminosity galaxy does not seem granted by its
appearance.

####  4.1.2 Relationships for passive and active galaxies

Figure 23: The upper panel shows the location of active and passive galaxies
in the [m/H] vs. mass plane, with the two classes identified by blue and red
colour, respectively. A general trend of metallicity increasing with mass can
be seen. In the lower panel, symbols are inversely proportional to the
measurement errors: the red and blue solid curves are weighted linear fit to
points representing passive and active galaxies. The blue dashed curve is the
relation obtained after excluding galaxies with metallicities lower than -3.
The red and blue shaded areas mark the location of the tracks taken from Fig.
21. The dashed curves on top of the panel represent mass-metallicity relations
from Peng et al. (2015), with red and blue colour distinguishing passive and
active galaxies.

To compare the MZR of AC114 galaxies with other relations from the literature,
they are plotted again in Fig. 23, without model tracks. Weighted linear fits
to the data run nearly parallel to the model tracks discussed in the previous
section, both for passive and active galaxies. Therefore in our scenario,
mass-metallicity relations for AC114 objects are naturally generated by galaxy
evolution.

In the lower panel of the figure, dashed blue and red lines are the
relationships for active and passive galaxies found by Peng et al. (2015), who
analysed the Sloan Digital Sky Survey (SDSS) spectra of local galaxies
($z\approx 0.05$). Clearly, galaxies in the AC114 cluster have much lower
metallicities than field ones, and for the case of passive galaxies, we must
conclude that the dense cluster environment curtailed star formation of its
more massive members in a short time after their formation.

Active galaxies are evolving along different tracks than passive ones,
probably because of lower metal yields. And they are also evolving slower, as
it can be inferred from their younger ages and larger mass fractions. We also
expect that, even if allowed to reach the end of their chemical evolution,
they would still be $\sim 1$ dex more metal poor than local active galaxies.
The likely explanation is that the cluster environment facilitates the loss of
metals from these lower mass objects, thus reducing their effective yields
even further.

An evolutionary scenario would then be that the first objects to form got the
largest masses, formed significant amounts of metals in SF bursts that lasted
only a few Gyr, cut short through their mutual interactions. The stripped gas
stayed in the cluster potential and is now facilitating the loss of metals
from less massive galaxies that formed later via ram pressure stripping.

##  5 Summary

In this paper, we investigate the mass-metallicity relation for galaxies in
the Abell cluster AC114 from 7 hours of VIMOS/MR data collected at the ESO-VLT
telescope in 2009. The dynamical analysis of Proust et al. (2015) allowed us
to select cluster members, whose spectra are here analyzed with models from
Bruzual & Charlot (2003).

Active and passive galaxies are identified based on the presence/absence of
the [O ii]$\lambda 3727$, and/or [O iii]$\lambda\lambda 4959,5007$ and/or
H$\beta$ emission lines, depending on the galaxy redshift, and we conclude
that active galaxies have lower average masses and metallicities than passive
ones.

We establish that the mass-metallicity relation of the cluster is found to be
steeper than that for galaxies in the local universe. In a forthcoming paper,
the MZR of active galaxies, based on the oxygen abundance of their gaseous
component, will be presented and discussed in light of the present results.

Fig. 2 showed that passive and active galaxies tend to occupy central and
external regions respectively. Therefore the results above show the well known
fact that galaxy evolution in the centre of galaxy clusters is faster than
that in their periphery. In the central regions of AC114, and in the course of
the last 9.8 Gyr, galaxies reached masses and metallicities that are
comparable to present-day galaxies, while in the outskirts galaxies are still
forming stars, and in their past they have built substantially less mass than
central objects.

The case of large active galaxies highlights the fact that conclusions about
their global properties are influenced by population gradients within their
bodies. Therefore a conclusive study could only be realised by a two-
dimensional mapping of their SEDs, such as can be afforded by IFU instruments:
in this respect, MUSE at the VLT would be an obvious choice.

## Acknowledgements

DP thanks ESO in the context of the Visiting Scientists program for its
hospitality at Santiago (Chile). We thank the anonymous referee for a careful
reading of the manuscript, which greatly improved the presentation of our
work. Alain Andrade is also thanked for helping in the revision of the paper.

## Data Availability

The complete redshift values of the 163 galaxies of this work are listed in
Table 5 and the description of each column can be found in subsection 2.1.

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Table 5: Positions, photometric data and velocities for galaxies of AC114. The R column is the correlation peak height to the amplitude of the asymmetric noise (see text for explanation). Gal. | R.A. | Dec. | $B_{j}$ | $R_{f}$ | Rcomp | redshift | error | R | notes
---|---|---|---|---|---|---|---|---|---
| (J2000) | (J2000) | mag | mag | mag |  |  |  |
Quadrant 1 |  |  |  |  |  |  |  |  |
1 | 22 59 20.70 | -34 42 24.6 |  |  | 20.0 | 0.39852 | 0.00030 | 3.50 | em:H$\alpha$,S1
2 | 22 59 31.37 | -34 42 28.4 |  |  | 20.5 | 0.35306 | 0.00018 |  | em:OII,H$\beta$,2OIII,H$\alpha$,S1
3 | 22 59 33.46 | -34 42 39.8 |  |  | 18.0 | 0.05717 | 0.00011 |  | em:H$\beta$,2OIII,H$\alpha$,S1
5 | 22 59 32.31 | -34 42 58.2 | 20.2 |  | 20.1 | 0.05842 | 0.00009 |  | em:H$\beta$,2OIII,H$\alpha$,S1
8 | 22 59 36.51 | -34 43 25.4 | 20.2 |  | 20.3 | 0.38058 | 0.00022 |  | em:H$\alpha$,S1
9 | 22 59 25.99 | -34 43 31.9 |  |  |  | 0.21032? |  |  | very uncertain
10 | 22 59 18.60 | -34 43 43.6 | 20.5 | 18.8 | 18.9 | 0.21940 | 0.00024 | 4.76 | em: H$\alpha$,S1 65840 km s-1
11 | 22 59 17.61 | -34 43 57.1 |  |  | 19.3 | 0.62605 | 0.00031 |  | measured on OII,H,K (CaI)
12 | 22 59 20.05 | -34 44 09.4 | 20.8 | 19.5 | 18.2 | 0.22127 | 0.00012 | 3.56 | em:H$\beta$,H$\alpha$ 66509 km s-1
13 | 22 59 20.00 | -34 44 20.3 |  |  | 20.0 | 0.40975 | 0.00022 | 4.79 |
15 | 22 59 28.09 | -34 44 29.0 |  |  | 19.4 | 0.07309 | 0.00045 |  | em:H$\beta$,2OIII,H$\alpha$
16 | 22 59 28.82 | -34 44 44.1 |  |  | 18.6 | 0.40819 | 0.00039 | 3.57 |
17 | 22 59 38.42 | -34 44 52.4 |  |  | 20.5 | 0.63163 | 0.00045 | 3.09 | very weak
18 | 22 59 23.15 | -34 45 03.0 |  |  | 19.7 | 0.31301 | 0.00007 |  | em:H$\alpha$,S1
19 | 22 59 18.51 | -34 45 13.1 | 21.2 | 20.3 | 19.6 | 0.40876 | 0.00022 |  | em:OII,H$\beta$,2OIII,H$\alpha$,S1
20 | 22 59 40.37 | -34 45 21.8 |  |  | 18.1 | 0.23192 | 0.00018 | 5.23 |
22 | 22 59 24.80 | -34 45 36.6 | 22.3 | 19.6 | 18.8 | 0.41015 | 0.00025 | 5.19 | em:H$\beta$,H$\alpha$
24 | 22 59 28.59 | -34 45 49.8 | 21.6 | 19.6 | 18.9 | 0.33369 | 0.00018 | 4.84 | em:OII
25 | 22 59 35.08 | -34 45 57.0 | 22.4 |  | 20.9 | 0.48600 | 0.00055 |  | measured on H,K (CaI)
26 | 22 59 32.09 | -34 46 13.6 |  |  | 19.6 | 0.53351 | 0.00056 |  | measured on OII,H,K (CaI)
27 | 22 59 32.34 | -34 46 23.3 |  |  | 19.9 | 0.46615 | 0.00022 | 3.22 |
28 | 22 59 26.16 | -34 46 32.5 | 20.2 | 19.2 | 19.0 | 0.35118 | 0.00028 | 3.11 | weak confirmed on H,K (CaI)
30 | 22 59 36.18 | -34 46 43.1 | 20.8 | 18.8 | 19.6 | 0.31860 | 0.00033 | 3.15 |
31 | 22 59 33.65 | -34 46 49.2 |  |  | 19.9 | 0.33342 |  |  | measured on H$\alpha$
34 | 22 59 25.25 | -34 47 05.9 |  |  | 20.0 | 0.63182 | 0.00068 |  | measured on H,K (CaI)
36 | 22 59 34.22 | -34 47 19.9 |  |  | 19.6 | 0.49978 |  |  | measured on OII
37 | 22 59 23.76 | -34 47 27.4 | 21.0 | 19.8 | 19.6 | 0.31604 | 0.00004 |  | em:H$\beta$,2OIII,H$\alpha$,S1
38 | 22 59 18.07 | -34 47 39.8 |  |  | 19.9 | 0.49805 | 0.00064 | 3.04 | very weak
40 | 22 59 23.41 | -34 47 57.6 | 20.6 | 18.9 | 18.1 | 0.32283 | 0.00027 | 3.37 | em:H$\beta$,2OIII,H$\alpha$,S1: 96857 km s-1
41 | 22 59 16.77 | -34 48 06.8 |  |  | 20.0 | 0.47645 | 0.00054 |  | measured on H,K (CaI)
44 | 22 59 15.92 | -34 48 43.8 | 22.2 | 21.1 | 19.1 | 0.57779 | 0.00066 |  | measured on OII,H,K (CaI)
45a | 22 59 31.40 | -34 48 50.3 |  |  |  | 0.35159 |  |  | measured on H$\alpha$,S1
45b | 22 59 31.40 | -34 48 50.4 |  |  |  | 0.34461 |  |  | measured on H$\alpha$,S1
Quadrant 2 |  |  |  |  |  |  |  |  |
1a | 22 59 16.30 | -34 33 40.6 | 20.8 | 19.0 | 19.6 | 0.25782 | 0.00023 | 3.43 | em:H$\beta$,2OIII,H$\alpha$
1b | 22 59 16.30 | -34 33 40.6 |  |  |  | 0.33069 | 0.00024 | 3.18 | em:H$\alpha$, 2 galaxies
2 | 22 59 18.23 | -34 33 50.7 |  |  | 18.2 | 0.49578 | 0.00029 |  | em:OII,H$\beta$,2OIII,H$\alpha$,S1
3 | 22 59 18.00 | -34 33 58.5 |  |  | 20.5 | 0.41000 | 0.00033 |  | em:OII,H$\beta$,2OIII,H$\alpha$
4 | 22 59 13.67 | -34 34 11.1 |  |  | 19.3 | 0.40846 | 0.00033 |  | em:H$\beta$,2OIII
5 | 22 59 15.03 | -34 34 23.5 | 22.5 | 20.1 | 18.8 | 0.41440 | 0.00034 | 3.01 | very uncertain
6 | 22 59 23.91 | -34 34 32.7 | 21.3 | 19.2 | 18.9 | 0.42284 | 0.00021 | 6.45 |
7 | 22 59 21.45 | -34 34 41.8 |  |  | 18.6 | 0.56176 | 0.00052 |  | measured on H,K (CaI)
8 | 22 59 12.89 | -34 34 48.4 |  |  | 21.2 | 0.25744 | 0.00018 |  | measured on 2OIII
9 | 22 59 17.73 | -34 34 54.1 |  |  | 20.9 | 0.37079 | 0.00063 |  | measured on 2OIII
10 | 22 59 18.77 | -34 35 02.2 | 21.5 | 19.9 | 19.5 | 0.40848 | 0.00023 | 3.36 |
11 | 22 59 16.88 | -34 35 08.5 |  |  | 17.6 | 0.34546 |  |  | measured on H$\alpha$
12 | 22 59 32.84 | -34 35 24.7 | 21.9 |  | 19.2 | 0.59523 |  |  | measured on OII
13 | 22 59 26.31 | -34 35 32.2 |  |  | 17.8 | 0.27013 | 0.00038 | 3.52 | very weak
14 | 22 59 20.21 | -34 35 43.7 |  |  | 18.7 | 0.37370 | 0.00009 |  | em:OII,H$\alpha$,S1
15 | 22 59 15.61 | -34 35 51.1 |  |  | 20.7 | 0.41045 | 0.00026 | 4.01 |
16 | 22 59 19.18 | -34 35 59.7 | 21.4 | 20.4 | 19.7 | 0.56822 | 0.00074 |  | measured on OII,H,K (CaI)
18 | 22 59 12.16 | -34 36 28.0 | 21.0 |  | 19.4 | 0.38110 | 0.00042 | 4.03 |
19 | 22 59 19.31 | -34 36 38.6 | 22.00 | 20.1 | 18.8 | 0.37706 | 0.00021 | 6.71 |
21 | 22 59 21.69 | -34 36 53.5 |  |  | 19.2 | 0.41304 | 0.00020 | 6.52 |
22 | 22 59 36.79 | -34 36 53.5 |  |  | 21.3 | 0.39104 |  |  | em:H$\alpha$
23 | 22 59 21.57 | -34 37 07.8 | 19.3 | 18.4 | 18.3 | 0.41310 | 0.00022 | 6.81 |
24 | 22 59 36.81 | -34 37 21.2 |  |  | 19.2 | 0.37972 | 0.00052 |  | em:OII,H$\beta$,2OIII
Gal. | R.A. | Dec. | $B_{j}$ | $R_{f}$ | Rcomp | redshift | error | R | notes
---|---|---|---|---|---|---|---|---|---
| (J2000) | (J2000) | mag | mag | mag |  |  |  |
26 | 22 59 12.34 | -34 37 37.5 | 22.3 | 20.2 | 19.4 | 0.48676 | 0.00018 |  | em:OII,H$\beta$,2OIII
27 | 22 59 21.00 | -34 37 46.0 | 20.6 | 20.0 | 19.9 | 0.49642 | 0.00024 | 4.53 |
28 | 22 59 15.56 | -34 37 57.2 | 20.8 | 19.4 | 19.7 | 0.61361 | 0.00047 |  | measured on OII,H$\beta$
29 | 22 59 17.62 | -34 38 03.8 |  |  | 20.9 | 0.52230 |  |  | measured on OII
31 | 22 59 24.67 | -34 38 18.4 | 21.4 | 20.6 | 18.5 | 0.49717 | 0.00016 |  | em:OII,H$\beta$,2OIII,H$\alpha$
32 | 22 59 25.00 | -34 38 35.6 | 20.8 | 19.5 | 19.5 | 0.30913 | 0.00028 |  | em:H$\beta$,2OIII,H$\alpha$
33 | 22 59 16.65 | -34 38 43.7 | 20.1 | 18.1 | 20.4 | 0.46922 | 0.00051 | 3.17 | em:H$\beta$ 140472 km s-1
35 | 22 59 40.57 | -34 39 05.1 |  |  | 19.7 | 0.31290 | 0.00016 | 7.72 |
36 | 22 59 25.79 | -34 39 11.8 | 21.4 | 20.0 | 19.8 | 0.46367 | 0.00020 | 6.71 |
37 | 22 59 13.94 | -34 39 33.0 | 21.8 | 20.7 | 18.7 | 0.13963 | 0.00063 |  | measured on H$\beta$,2OIII
38 | 22 59 23.08 | -34 39 47.1 |  |  | 20.1 | 0.29869 | 0.00017 |  | em:H$\beta$,2OIII,H$\alpha$,S1
39 | 22 59 39.01 | -34 39 53.1 |  |  | 20.6 | 0.37203 | 0.00012 |  | em:OII,H$\beta$,2OIII,H$\alpha$
40 | 22 59 36.39 | -34 39 58.4 |  |  | 20.3 | 0.41913 | 0.00023 | 4.96 |
41 | 22 59 43.80 | -34 40 06.7 |  |  | 20.5 | 0.54027 |  |  | measured on OII
Quadrant 3 |  |  |  |  |  |  |  |  |
1 | 22 58 41.57 | -34 33 18.1 | 20.3 | 18.8 | 19.8 | 0.32951 | 0.00015 |  | em:H$\beta$,2OIII,H$\alpha$,S1
2 | 22 58 41.21 | -34 33 27.8 | 20.5 | 18.8 | 18.7 | 0.33307 | 0.00011 |  | em:H$\beta$,2OIII,H$\alpha$,S1
3 | 22 58 40.11 | -34 33 34.6 | 21.9 | 19.2 | 19.4 | 0.33127 | 0.00024 | 5.48 |
4 | 22 58 40.08 | -34 33 42.2 | 21.3 | 19.9 | 20.0 | 0.21259 | 0.00019 |  | em:H$\beta$,2OIII,H$\alpha$,S1
5 | 22 58 41.10 | -34 33 49.0 |  |  | 20.5 | 0.33189 | 0.00020 |  | em:H$\beta$,2OIII,H$\alpha$,S1
7 | 22 58 48.40 | -34 34 01.6 |  |  |  | 0.31092 | 0.00054 |  | measured on H,K (CaI)
8 | 22 58 42.73 | -34 34 23.0 | 22.0 |  | 19.7 | 0.39969 | 0.00028 | 3.48 | em:H$\alpha$:119724 km s-1
10 | 22 58 22.64 | -34 34 33.3 | 19.8 | 18.7 | 18.8 | 0.31365 | 0.00012 | 4.63 |
11 | 22 58 35.95 | -34 34 45.6 | 22.4 | 20.6 | 20.1 | 0.21245 | 0.00026 |  | em:H$\beta$,2OIII,H$\alpha$,S1
12 | 22 58 36.88 | -34 34 58.2 | 21.5 | 19.6 | 19.0 | 0.30939 | 0.00009 |  | em:H$\beta$,2OIII,H$\alpha$,S1
13 | 22 58 51.29 | -34 35 07.6 |  |  | 19.3 | 0.65153 | 0.00048 | 3.03 | very weak
15 | 22 58 41.00 | -34 35 21.4 |  |  | 20.3 | 0.47562 | 0.00038 | 3.01 | very weak
16 | 22 58 45.84 | -34 35 30.4 | 20.9 | 20.4 | 20.0 | 0.17739 | 0.00012 |  | em:H$\beta$,2OIII,H$\alpha$,S1
18 | 22 58 56.50 | -34 35 48.2 | 20.8 | 19.3 | 19.5 | 0.22176 | 0.00009 |  | em:H$\beta$,2OIII,H$\alpha$,S1
19 | 22 58 23.26 | -34 35 57.0 | 20.0 | 18.6 | 18.9 | 0.15860 | 0.00015 | 4.84 |
20 | 22 58 32.46 | -34 36 07.9 |  |  | 19.2 | 0.35487 | 0.00038 | 3.07 | confirmed on OII
21 | 22 58 32.39 | -34 36 15.5 | 21.1 | 19.2 | 20.8 | 0.35231 | 0.00021 |  | em:H$\beta$,H$\alpha$,S1,S2
23 | 22 58 22.99 | -34 36 26.7 | 22.1 | 20.5 | 20.5 | 0.35371 | 0.00035 |  | measured on H,K (CaI)
24 | 22 58 28.94 | -34 36 34.7 |  |  | 18.2 | 0.30386 | 0.00045 | 3.01 | very weak
25 | 22 58 40.69 | -34 36 41.1 |  |  | 20.1 | 0.59540 | 0.00068 |  | measured on H,K (CaI)
26 | 22 58 28.24 | -34 36 48.7 |  |  | 20.2 | 0.31477 | 0.00024 | 4.07 |
28 | 22 58 24.28 | -34 37 02.1 |  |  | 18.4 | 0.25761 | 0.00018 | 5.51 |
29 | 22 58 32.32 | -34 37 07.5 | 21.6 | 20.1 | 20.6 | 0.47283 | 0.00038 | 3.15 |
30 | 22 58 27.93 | -34 37 16.9 | 21.2 | 20.1 | 21.1 | 0.40221 | 0.00023 |  | em:OII, H$\beta$,2OIII,H$\alpha$
31 | 22 58 44.38 | -34 37 24.2 | 22.2 | 20.2 | 20.1 | 0.39764 | 0.00040 | 3.04 | em:H$\beta$,H$\alpha$
33 | 22 58 39.09 | -34 37 52.3 |  |  | 18.8 | 0.34480 |  |  | measured on H$\alpha$
34 | 22 58 39.66 | -34 38 03.8 |  |  | 20.1 | 0.35179 | 0.00011 |  | em:OII,H$\beta$,2OIII,H$\alpha$,S1
35 | 22 58 33.67 | -34 38 18.9 | 21.4 | 20.1 | 18.9 | 0.43091 | 0.00014 |  | em:H$\beta$,2OIII,H$\alpha$,S1
36 | 22 58 31.51 | -34 38 24.2 | 22.3 | 20.7 | 20.5 | 0.20326 | 0.00031 |  | em:H$\alpha$,S1
38 | 22 58 41.21 | -34 38 40.5 |  |  | 19.3 | 0.43205 | 0.00029 | 3.58 | em: OII
39 | 22 58 46.63 | -34 38 45.7 | 21.9 | 20.6 | 20.2 | 0.39935 | 0.00011 |  | em:OII,H$\beta$,2OIII,H$\alpha$,S1
40 | 22 58 34.07 | -34 38 53.1 | 21.9 | 19.9 | 19.3 | 0.31828 | 0.00031 | 4.12 |
43 | 22 58 39.42 | -34 39 13.5 |  |  | 20.2 | 0.51016 | 0.00028 | 3.32 |
44 | 22 58 39.45 | -34 39 24.1 | 22.2 | 19.4 | 19.5 | 0.22464 | 0.00017 | 3.34 | em: H$\alpha$,S1 67344 km s-1
45 | 22 58 43.50 | -34 39 32.3 | 21.6 | 20.6 | 20.1 | 0.62790 | 0.00061 |  | measured on OII,H,K (CaI)
46 | 22 58 42.70 | -34 39 40.6 |  |  | 20.4 | 0.52153 | 0.00028 | 3.21 |
49 | 22 58 37.09 | -34 40 06.8 | 19.4 | 18.2 | 18.9 | 0.20657 | 0.00019 | 4.82 | em:H$\alpha$,S1 61753 km s-1
Quadrant 4 |  |  |  |  |  |  |  |  |
2 | 22 58 34.99 | -34 42 24.0 |  |  | 19.5 | 0.31982 | 0.00035 | 5.82 | em:H$\alpha$ weak
3 | 22 58 45.17 | -34 42 28.9 | 22.6 | 19.1 | 20.1 | 0.41385 | 0.00042 |  | measured on OII,H,K (CaI)
4 | 22 58 35.53 | -34 42 35.8 |  |  | 19.7 | 0.31108 | 0.00014 | 3.42 |
Gal. | R.A. | Dec. | $B_{j}$ | $R_{f}$ | Rcomp | redshift | error | R | notes
---|---|---|---|---|---|---|---|---|---
| (J2000) | (J2000) | mag | mag | mag |  |  |  |
5 | 22 58 42.57 | -34 42 42.5 |  |  | 20.6 | 0.42558 | 0.00012 |  | em:H$\beta$,2OIII,H$\alpha$,S1
7 | 22 58 32.37 | -34 42 55.1 | 20.4 | 19.7 | 20.0 | 0.16952 | 0.00009 |  | em:H$\beta$,2OIII,H$\alpha$,S1
8 | 22 58 32.75 | -34 43 02.0 |  |  | 19.5 | 0.17061 | 0.00022 |  | em:H$\beta$,2OIII,H$\alpha$,S1
9a | 22 58 31.70 | -34 43 18.8 | 21.0 | 18.9 | 18.4 | 0.29095 | 0.00009 | 10.95 |
9b | 22 58 31.70 | -34 43 18.8 |  |  |  | 0.25836 | 0.00035 |  | measured on H,K (CaI)
10 | 22 58 25.31 | -34 43 25.5 | 20.5 | 18.8 | 19.5 | 0.31228 |  |  | em:H$\alpha$ very weak
11 | 22 58 32.43 | -34 43 30.3 | 19.0 | 17.4 | 20.0 | 0.30473 | 0.00008 | 3.02 | em:H$\alpha$
12 | 22 58 26.84 | -34 43 34.1 | 22.1 | 21.0 | 19.7 | 0.31166 | 0.00026 | 3.05 |
13 | 22 58 26.64 | -34 43 44.0 | 22.4 | 20.5 | 20.9 | 0.47564 | 0.00016 | 3.40 |
14 | 22 58 41.60 | -34 43 46.5 | 20.5 | 18.9 | 20.7 | 0.30460 | 0.00007 |  | em:OII,H$\beta$,2OIII,H$\alpha$,S1
15 | 22 58 25.91 | -34 43 52.0 | 21.2 | 19.8 | 20.9 | 0.30799 | 0.00030 |  | em:H$\beta$,2OIII
16 | 22 58 29.50 | -34 43 55.5 | 22.2 | 19.8 | 20.4 | 0.31854 | 0.00035 | 3.52 |
18 | 22 58 25.81 | -34 44 06.4 |  |  | 20.6 | 0.32247 | 0.00029 | 3.19 |
19 | 22 58 26.24 | -34 44 12.1 |  |  | 20.2 | 0.31070 | 0.00011 |  | em:H$\beta$,2OIII,H$\alpha$,S1
20 | 22 58 39.44 | -34 44 18.6 |  |  | 19.6 | 0.30863 | 0.00011 |  | em:H$\beta$,2OIII,H$\alpha$,S1
21 | 22 58 50.72 | -34 44 16.7 |  |  | 19.5 | 0.30605 | 0.00024 | 4.32 |
22 | 22 58 39.87 | -34 44 28.0 |  |  | 19.4 | 0.31445 | 0.00022 | 5.69 |
23 | 22 58 35.11 | -34 44 34.8 | 21.9 |  | 20.6 | 0.32268 | 0.00030 | 3.09 |
24 | 22 58 33.19 | -34 44 39.7 |  |  | 20.5 | 0.31448 | 0.00028 | 3.12 |
25 | 22 58 30.22 | -34 44 42.4 |  |  | 19.4 | 0.30132 | 0.00035 | 3.20 |
26 | 22 58 42.00 | -34 44 51.2 | 20.6 | 18.9 | 18.9 | 0.32462 | 0.00012 |  | em:OII,H$\beta$,2OIII,H$\alpha$,S1
27 | 22 58 29.86 | -34 44 55.5 |  |  | 21.4 | 0.31265 | 0.00023 | 3.08 |
28 | 22 58 26.83 | -34 45 02.6 |  |  | 20.0 | 0.31467 | 0.00032 | 3.53 |
31 | 22 58 36.47 | -34 45 17.4 | 21.0 | 19.3 | 19.1 | 0.31779 | 0.00024 | 5.33 |
32 | 22 58 38.00 | -34 45 24.1 | 20.2 | 17.7 | 18.2 | 0.31708 | 0.00024 | 5.65 |
33 | 22 58 41.30 | -34 45 30.6 | 21.1 | 19.6 | 19.8 | 0.10193 | 0.00029 |  | em:OII,H$\beta$,2OIII,H$\alpha$,S1
35 | 22 58 47.92 | -34 45 45.0 | 21.7 | 19.5 | 20.2 | 0.17040 | 0.00011 |  | em:H$\beta$,2OIII,H$\alpha$,S1
36 | 22 58 35.91 | -34 45 50.1 | 19.8 | 17.7 | 18.4 | 0.31307 | 0.00019 | 5.65 |
37 | 22 58 29.65 | -34 45 57.5 |  |  | 19.1 | 0.17051 | 0.00018 | 3.01 | uncertain
38 | 22 58 28.42 | -34 46 02.3 | 22.0 | 19.8 | 19.2 | 0.30100 | 0.00026 | 5.63 |
40 | 22 58 24.92 | -34 46 18.7 | 20.8 | 19.0 | 19.0 | 0.31779 | 0.00004 |  | em:H$\alpha$
41 | 22 58 27.85 | -34 46 26.3 |  |  | 19.7 | 0.31816 | 0.00033 | 4.55 |
42 | 22 58 29.75 | -34 46 32.1 | 22.2 | 20.5 | 19.9 | 0.50228 | 0.00035 | 3.04 |
43 | 22 58 28.88 | -34 46 40.3 |  |  | 20.2 | 0.31717 | 0.00033 | 3.68 |
44 | 22 58 40.93 | -34 46 45.7 | 20.6 | 18.6 | 18.3 | 0.31523 | 0.00030 | 4.73 |
45 | 22 58 36.47 | -34 46 51.2 |  |  | 20.6 | 0.31389 | 0.00019 | 4.66 |
46 | 22 58 51.00 | -34 47 00.6 | 20.8 | 19.5 | 20.1 | 0.32953 | 0.00028 | 3.10 |
47 | 22 58 34.80 | -34 47 07.1 | 22.2 |  | 19.3 | 0.31674 | 0.00019 | 6.24 |
48 | 22 58 29.14 | -34 47 13.1 | 21.6 | 20.0 | 20.0 | 0.31462 | 0.00022 | 5.29 |
49 | 22 58 39.56 | -34 47 17.4 | 21.8 | 19.1 | 20.0 | 0.31818 | 0.00024 | 4.89 |
50 | 22 58 48.11 | -34 47 21.3 |  |  | 20.2 | 0.33032 | 0.00039 | 3.02 | very weak
51 | 22 58 46.31 | -34 47 31.5 | 20.8 | 18.8 | 19.1 | 0.30954 | 0.00019 | 6.32 |
52 | 22 58 26.11 | -34 47 37.5 | 22.6 | 20.9 | 18.2 | 0.31789 | 0.00042 | 3.04 | confirmed on H,K (CaI)
54 | 22 58 33.80 | -34 47 46.6 | 22.3 | 19.5 | 18.9 | 0.34779 | 0.00030 | 3.06 | weak
55a | 22 58 33.96 | -34 47 53.3 |  |  | 20.2 | 0.31776 | 0.00031 | 3.26 |
55b | 22 58 33.96 | -34 47 53.3 |  |  |  | 0.17199 |  |  | em:H$\beta$,2OIII,H$\alpha$,S1
56 | 22 58 44.05 | -34 47 59.4 |  |  | 18.4 | 0.31380 | 0.00015 |  | measured on H,K (CaI)
57 | 22 58 49.74 | -34 48 01.9 | 21.6 | 19.7 | 19.1 | 0.31451 | 0.00028 | 3.35 |
58 | 22 58 41.65 | -34 48 06.7 | 21.9 | 19.2 | 19.2 | 0.31082 | 0.00028 | 6.10 |
60 | 22 58 44.45 | -34 48 21.3 |  |  | 19.0 | 0.32473 | 0.00031 | 4.42 |
61 | 22 58 38.22 | -34 48 27.8 | 20.7 | 18.6 | 19.7 | 0.72072 | 0.00014 | 3.05 |
62 | 22 58 42.81 | -34 48 33.4 | 21.0 | 18.9 | 29.9 | 0.31090 | 0.00047 | 5.43 |
63 | 22 58 52.67 | -34 48 39.1 |  |  | 20.2 | 0.31874 | 0.00024 | 4.98 |
64 | 22 58 31.46 | -34 48 53.3 |  |  | 18.3 | 0.31447 | 0.00030 | 5.28 |
65 | 22 58 34.69 | -34 49 01.3 |  |  | 20.6 | 0.47788 | 0.00023 | 5.33 |

##  Appendix A Spectra of galaxies identified as active based on [O
ii]$\lambda 3727$

Figure 24: Spectra of galaxies identified as active based on [O ii]$\lambda
3727$.

##  Appendix B Postage stamps

Figure 25: Images of galaxies classified as passive, from the VIMOS
preimaging, sorted in luminosity from top to bottom. The number identifies the
slit, while the quadrant can be evinced from Table 3. Each frame is
10.5$\arcsec\times$ 10.5$\arcsec$, and the image quality is $\sim
1^{\prime\prime}$ FWHM. Most galaxies have regular shapes, as expected from
E/SO types, but notable exceptions are galaxies 44, 51, 49, 53, and 58, which
appear to be interacting with smaller neighbours.  Figure 26: Same as Fig. 25
for galaxies classified as active.

##  Appendix C Galaxies in common with CS87

Figure 27: Representation on sky of galaxies from this work (black symbols), and from Couch and Sharples (1987; open circles). After transforming coordinates from B1950.0 to J2000.0, objects were matched with a radius of $10^{\prime\prime}$. Matching galaxies are marked with a double circle.  Table 6: Parameters of galaxies in common between this work and Couch & Sharples (1987). Columns 1 and 2 are the identifications in the respective works; columns 3, 4 and 5, are coordinates in this work and the separation in arcsec; columns 6 and 7 are the redshift measured in this work and the difference with CS87; columns 8, 9, and 10 are the colours and magnitude from CS87 and the difference in magnitude with the present work; column 11 is the cross-correlation quality from CS87; columns 12 and 13 are the [Oii]$\lambda 3727$ and H$\delta$ lines measurements from CS87; columns 14 and 15 are the metallicity and error from this work; and finally the last column is our galaxy classification, as (p)assive or (a)ctive. An empty line separates possible matches from unlikely ones.  S23 | CS87 | RA[2000] | DEC[2000] | d′′ | z | $\Delta z$ | B${}_{{\rm J}}-$RF | RF | $\Delta R$ | r | [Oii] | H$\delta$ | [m/H] |  $\epsilon$[m/H] | flag
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
1 | 2 | 3 | 4 | 5 | 6 | 7 | 8 | 9 | 10 | 11 | 12 | 13 | 14 | 15 | 16
Q4-51 | C5 | 344.693 | -34.792 | 0.0 | 0.310 | -0. | 0011 | 2.20 | 19.19 | -0. | 09 | 6.0 |  |  | -0.29 | 0.36 | p
Q4-58 | C145 | 344.674 | -34.802 | 0.0 | 0.311 | -0. | 0015 | 2.37 | 19.18 | 0. | 02 | 7.1 |  | 1.3 | -0.40 | 0.24 | p
Q4-49 | C247 | 344.665 | -34.788 | 0.0 | 0.318 | -0. | 0015 | 2.46 | 19.12 | 0. | 88 | 7.6 |  | 2.8 | -0.18 | 0.26 | p
Q4-47 | C832 | 344.645 | -34.784 | 3.6 | 0.317 | -0. | 0011 | 2.25 | 19.98 | -0. | 68 | 5.3 |  | 3.6 | -0.57 | 0.27 | p
Q4-44 | C243 | 344.670 | -34.779 | 2.9 | 0.315 | -0. | 0100 | 1.58 | 19.57 | -1. | 27 | 2.2 | 19.4 | 5.2 | -0.08 | 0.08 | p
Q4-63 | C191 | 344.720 | -34.813 | 7.8 | 0.319 | 0. | 0144 | 1.63 | 19.69 | 0. | 51 | 4.7 |  | 5.6 | -1.36 | 0.35 | p
Q4-56 | C134 | 344.681 | -34.801 | 9.5 | 0.314 | 0. | 2172 | 0.86 | 19.66 | -1. | 26 |  |  |  | -0.04 | 0.15 | a
Figure 28: Spectra of galaxies in common with CS87. The separation on sky is
displayed in each panel, with separation increasing from top to bottom.
Vertical dotted lines mark the position of [Oii]$\lambda 3727$ and H$\delta$.

After transforming from the B1950.0 to the J2000.0222Using the NED coordinate
calculator https://ned.ipac.caltech.edu/ system, coordinates were used
together with a matching radius of 10 arcsec, to find common objects between
this work and Couch & Sharples (1987). The result is shown in Fig. 27 and is
summarised in Table 6, where galaxies are listed from top to bottom in order
of increasing separation on sky. There are only a few galaxies in common, and
in particular only three have null separation on sky; galaxy Q4-47 might also
be a possible match, because the difference in redshift is comparable to the
first three. The other three galaxies are likely not the same object in the
two works, because of the large difference in redshift. Besides, [Oii]$\lambda
3727$ is detected in galaxy C243 but it is absent in our Q4-44; and
conversely, it is not present in C134 but we detect it in Q4-56, as Fig. 28
shows. Indeed, Q4-56 is classified as ‘active’ as we have shown above.

Based on their colours, all four common galaxies belong to the ‘red sequence’
in CS87 and have normal H$\delta$ absorption, with the exception of C832 which
has a marginally stronger absorption. Figure 28 shows the spectra of the seven
galaxies, and it confirms the absence of emission lines in the four best
matches.

##  Appendix D Summary of classification of all galaxies

The results of Section 2.2 are summarised in Table 7. In the table, and from
left to right, columns contain an ordinal number, the VIMOS quadrant and slit,
the redshift, the classification flag (a)ctive or (p)assive, and notes: this
scheme is repeated in three blocks. Notes are the following: [Oii] = galaxy
identified by the presence of [O ii]$\lambda 3727$; OI = bad sky subtraction
of [O i]$\lambda 5579$ oxygen line, which creates a spurious [O ii]$\lambda
3727$ emission line for objects at redshift $\simeq 0.4967$; M-dwarf = object
is actually a local star; :: = spectrum is uncertain and is not used in this
paper.

Table 7: Classification of galaxies as described in Section 2.2.  n | q | sl | z | a/p | notes | n | q | sl | z | a/p | notes | n | q | sl | z | a/p | notes
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
1 | 1 | 1 | 0.3985 | a | [Oii] | 56 | 2 | 29 | 0.523 | a |  | 110 | 4 | 7 | 0.1695 | a |
2 | 1 | 2 | 0.3531 | a |  | 57 | 2 | 31 | 0.4972 | a |  | 111 | 4 | 8 | 0.1706 | a |
3 | 1 | 3 | 0.05717 | a |  | 58 | 2 | 32 | 0.3091 | a |  | 112 | 4 | 9 | 0.2584 | p |
4 | 1 | 5 | 0.05842 | a |  | 59 | 2 | 33 | 0.5797 | a |  | 113 | 4 | 10 | 0.3123 | p |
5 | 1 | 8 | 0.3806 | a |  | 60 | 2 | 35 | 0.3129 | p |  | 114 | 4 | 11 | 0.3047 | p |
6 | 1 | 10 | 0.2194 | p |  | 61 | 2 | 36 | 0.4637 | a |  | 115 | 4 | 12 | 0.3117 | p |
7 | 1 | 11 | 0.626 | a | [Oii] | 62 | 2 | 37 | 0.1396 | a |  | 116 | 4 | 13 | 0.4756 | p |
8 | 1 | 12 | 0.2213 | a |  | 63 | 2 | 38 | 0.2987 | a |  | 117 | 4 | 14 | 0.3046 | a |
9 | 1 | 13 | 0.4098 | p |  | 64 | 2 | 39 | 0.372 | a |  | 118 | 4 | 15 | 0.308 | a |
10 | 1 | 15 | 0.07309 | a |  | 65 | 2 | 40 | 0.4191 | p |  | 119 | 4 | 16 | 0.3185 | a | ::
11 | 1 | 16 | 0.4082 | a | [Oii] | 66 | 2 | 41 | 0.5403 | a | [Oii] | 120 | 4 | 18 | 0.527 | p |
12 | 1 | 17 | 0.6316 | p |  | 67 | 3 | 1 | 0.3295 | a |  | 121 | 4 | 19 | 0.3107 | a |
13 | 1 | 18 | 0.313 | a |  | 68 | 3 | 2 | 0.3331 | a |  | 122 | 4 | 20 | 0.3086 | p |
14 | 1 | 19 | 0.4088 | a |  | 69 | 3 | 3 | 0.3313 | p |  | 123 | 4 | 21 | 0.3061 | p |
15 | 1 | 20 | 0.2319 | p |  | 70 | 3 | 4 | 0.2126 | a |  | 124 | 4 | 22 | 0.3145 | a |
16 | 1 | 22 | 0.4102 | a |  | 71 | 3 | 5 | 0.3319 | a |  | 125 | 4 | 23 | 0.3227 | p |
17 | 1 | 24 | 0.3337 | p |  | 72 | 3 | 6 | 0.4968 | p | OI | 126 | 4 | 24 | 0.3145 | p |
18 | 1 | 25 | 0.486 | p |  | 73 | 3 | 8 | 0.3997 | p | :: | 127 | 4 | 25 | 0.3013 | p |
19 | 1 | 26 | 0.5335 | a | [Oii] | 74 | 3 | 10 | 0.3137 | p | :: | 128 | 4 | 26 | 0.3246 | a |
20 | 1 | 27 | 0.4662 | p |  | 75 | 3 | 11 | 0.2125 | a |  | 129 | 4 | 27 | 0.3127 | p |
21 | 1 | 28 | 0.3512 | p |  | 76 | 3 | 12 | 0.3094 | a |  | 130 | 4 | 28 | 0.3147 | p |
22 | 1 | 30 | 0.3186 | a | [Oii] | 77 | 3 | 13 | 0.6515 | a |  | 131 | 4 | 31 | 0.3178 | p |
23 | 1 | 31 | 0.3334 | p |  | 78 | 3 | 15 | 0.4756 | p |  | 132 | 4 | 32 | 0.3171 | p |
24 | 1 | 34 | 0.6319 | p |  | 79 | 3 | 16 | 0.1774 | a |  | 133 | 4 | 33 | 0.1019 | a |
25 | 1 | 36 | 0.4998 | a | [Oii] | 80 | 3 | 17 | 0.4967 | p | OI | 134 | 4 | 35 | 0.1704 | a |
26 | 1 | 37 | 0.316 | a |  | 81 | 3 | 18 | 0.2218 | a |  | 135 | 4 | 36 | 0.3131 | a | [Oii]
27 | 1 | 38 | 0.498 | p |  | 82 | 3 | 19 | 0.1586 | p |  | 136 | 4 | 37 | 0.1705 | p |
28 | 1 | 40 | 0.3228 | a |  | 83 | 3 | 20 | 0.3549 | a |  | 137 | 4 | 38 | 0.301 | p |
29 | 1 | 41 | 0.4764 | p |  | 84 | 3 | 21 | 0.3523 | p | :: | 138 | 4 | 39 | 0.3259 | p | M-dwarf
30 | 1 | 44 | 0.5778 | a |  | 85 | 3 | 23 | 0.3537 | p |  | 139 | 4 | 40 | 0.3178 | a |
31 | 2 | 1 | 0.2578 | a |  | 86 | 3 | 24 | 0.3039 | p | :: | 140 | 4 | 41 | 0.3182 | p |
32 | 2 | 2 | 0.4958 | a |  | 87 | 3 | 25 | 0.5954 | p |  | 141 | 4 | 42 | 0.5023 | p |
33 | 2 | 3 | 0.41 | a |  | 88 | 3 | 26 | 0.3148 | p |  | 142 | 4 | 43 | 0.3172 | p |
34 | 2 | 4 | 0.4085 | a |  | 89 | 3 | 27 | 0.4969 | p | OI | 143 | 4 | 44 | 0.3152 | p |
35 | 2 | 5 | 0.4144 | p |  | 90 | 3 | 28 | 0.2576 | p |  | 144 | 4 | 45 | 0.3139 | p |
36 | 2 | 6 | 0.4228 | p |  | 91 | 3 | 29 | 0.4728 | a | :: | 145 | 4 | 46 | 0.3295 | a | [Oii]
37 | 2 | 7 | 0.5618 | p |  | 92 | 3 | 30 | 0.4022 | a |  | 146 | 4 | 47 | 0.3167 | p |
38 | 2 | 8 | 0.2574 | a |  | 93 | 3 | 31 | 0.332 | a |  | 147 | 4 | 48 | 0.3146 | a | [Oii]
39 | 2 | 9 | 0.3708 | a |  | 94 | 3 | 33 | 0.3448 | p |  | 148 | 4 | 49 | 0.3182 | p |
40 | 2 | 10 | 0.4085 | p | :: | 95 | 3 | 34 | 0.3518 | a |  | 149 | 4 | 50 | 0.3303 | p |
41 | 2 | 11 | 0.3455 | p | :: | 96 | 3 | 35 | 0.4309 | a |  | 150 | 4 | 51 | 0.3095 | p |
42 | 2 | 12 | 0.5952 | p |  | 97 | 3 | 36 | 0.2033 | a |  | 151 | 4 | 52 | 0.3179 | p |
43 | 2 | 13 | 0.2701 | p |  | 98 | 3 | 38 | 0.4321 | p | :: | 152 | 4 | 54 | 0.3478 | p |
44 | 2 | 14 | 0.3737 | a | [Oii] | 99 | 3 | 39 | 0.3994 | a |  | 153 | 4 | 55 | 0.172 | a |
45 | 2 | 15 | 0.4105 | p |  | 100 | 3 | 40 | 0.3183 | p |  | 154 | 4 | 56 | 0.3138 | a | [Oii]
46 | 2 | 16 | 0.5682 | a | [Oii] | 101 | 3 | 43 | 0.5101 | p | :: | 155 | 4 | 57 | 0.3145 | a | ::
47 | 2 | 18 | 0.6152 | a | [Oii] | 102 | 3 | 44 | 0.2246 | p | :: | 156 | 4 | 58 | 0.3108 | p |
48 | 2 | 19 | 0.3771 | a | [Oii] | 103 | 3 | 45 | 0.6279 | p |  | 157 | 4 | 59 | 0.3183 | p | M-dwarf
49 | 2 | 21 | 0.413 | a | [Oii] | 104 | 3 | 46 | 0.5215 | a | [Oii] | 158 | 4 | 60 | 0.3183 | a |
50 | 2 | 22 | 0.391 | a |  | 105 | 3 | 49 | 0.2066 | p |  | 159 | 4 | 61 | 0.7207 | p |
51 | 2 | 23 | 0.4131 | p |  | 106 | 4 | 2 | 0.3198 | p |  | 160 | 4 | 62 | 0.2838 | p |
52 | 2 | 24 | 0.3797 | a |  | 107 | 4 | 3 | 0.4139 | a | [Oii] | 161 | 4 | 63 | 0.3187 | p |
53 | 2 | 26 | 0.4868 | a |  | 108 | 4 | 4 | 0.3111 | p |  | 162 | 4 | 64 | 0.3145 | p |
54 | 2 | 27 | 0.4964 | a | [Oii] | 109 | 4 | 5 | 0.4135 | a |  | 163 | 4 | 65 | 0.4779 | p |
55 | 2 | 28 | 0.6136 | p | :: |  |  |  |  |  |  |  |  |  |  |  |