Abstract
We present initial results of a study of low surface brightness
dwarf galaxies within eight different clusters
(A262, A400, A426, A569, A1367, A1656, A2635, A3526)
as part of our program to determine the clustering properties, luminosity functions, and
morphologies of dwarf galaxies in a wider range of cluster environments. Deep
V-band images covering up to 1 square degree in each cluster were obtained
in addition to control fields which are used to statistically correct for
background galaxy
contamination.
In addition, fiber spectroscopy data has been obtained for over 500 galaxies
in A3526.
This data will be used to study the mass contribution of
and the effect of the cluster environment on these low mass galaxies.
1. Introduction
Dwarf galaxies are the most common type of galaxy thus far discovered and may
be substantial contributers of mass to the universe; with a high enough M/L ratio, dwarfs
and their associated dark matter halos may even compose a significant portion
of the overall cluster dark matter fraction. Our work addresses three issues related
to dwarf galaxies:
Dwarf Galaxy Luminosity Functions - First, the data will be used in
determining the cluster luminosity functions. The faint-end of the luminosity
function in clusters has been difficult to establish due to complex selection effects related to the detection of low luminosity, low SB systems.
Different cluster environments may play a significant role in determining the numbers
and morphologies of member galaxies causing cluster-to-cluster variations
in the luminosity function.
A thorough understanding of the faint-end slope will determine if
dwarfs (and their associated dark matter halos) can account for a substantial fraction of the cluster dark matter content.
Environmental Influences on Dwarf Properties - Because
the dwarfs are of low mass, they are more fragile than giant galaxies and
thus are a better indicator of environmental effects. Morphology, luminosity, structure (such as existence of a nucleus), and surface brightness may all be
affected by environmental factors. Specifically, we shall look at the dwarf properties with respect to their radial cluster distributions (and proximity
to larger galaxies) and between
different clusters.
For example, previous work has revealed that
the dwarf/giant ratio may increase with
increasing cluster richness (Ferguson & Sandage 1991) and Thompson and Gregory(1993) find a drop in the density of dwarf spheroidals toward the Coma cluster
core. Preliminary data analysis of A426 appears to support this latter
trend while A2634 does not.
Dwarfs at the Current Epoch - Finally, this data set will be used to determine the cluster-to-cluster variations in dwarf properties expected at the current epoch, which
is essential to interpret observations at higher redshift.
2. Observational Overview
Observations of A3526 were acquired on 13-20 May 1993 on the LCO 1m with a
TEK 3 2048 x 2048 CCD camera at prime focus. The seven northern clusters
were observed between Oct. 1994 and Dec. 1996. These were all taken at the Michigan-Dartmouth-MIT (MDM) 1.3m telescope
with a Loral 2048 x 2048 CCD. All data were taken in the V-band
for maximum light transmission and clusters were observed in a mosaic
pattern with each field exposed 4 X 15 minutes for a total exposure of
1 hour.
Control fields were also taken in the same manner 3 - 4 degrees away
from the cluster center. All data were calibrated directly
from observations of standards taken on each observing run. For fields
in which same night standards were not available, stars in overlap
regions of neighboring fields were used to obtain the calibration.
Total field coverage for each cluster
is about 1 square degree. The grey scale image (Fig. 2) displays the spatial
coverage
of A3526. While all data is currently in the V band,
B-band data for at least A1367 will be obtained in a future observing run.
In May 1995, a spectroscopic followup was made of the
brightest ~ 500 galaxies in A3526 using the LCO 2.5m with a
2D Frutti + fibers. A 600 line/mm grating was used to obtain a spectral resolution of
8.5 angstroms over the region 3800-6500A. Spectra were taken in 5 different setups and include cluster galaxies as faint
as Mv = -14.
The table below provides a summary of the CCD data obtained.
3. Data Reduction
Great attention was paid to flat fielding and cosmic ray removal in order to remove intrinsic variations
from the chip and to allow for the detection of very low surface brightness
galaxies. Each image was first flattened with twilight sky flats. A second
flat was then made by combining 10 flattened and normalized control fields.
To avoid undesirable artifacts, we used a percentile clipping algorithm after
removing the brightest pixels in each set of 10. All images were then
divided by this flat and the 4 images in each field were summed after cosmic ray removal.
The final sky fluctuations in the summed images were less than 0.2% of the
mean sky level.
We used the software package, FOCAS (Valdez 1989), to locate the dwarf galaxies.
FOCAS locates the objects in an image by searching for a specified number of
contiguous pixels with values greater than a specified threshold above the
local sky. Best parameter values were determined to be 10 contiguous pixels
and 3.5sigma for the detection of the lowest surface brightness galaxies
without the detection of noise. The objects were classified by defining a
PSF
and all objects not classified as galaxies were
removed from the catalog. The final list was then verified by eye to remove
any spurious detections. Control fields were reduced in exactly the same manner
and were used to determine the contribution of background galaxies
in each magnitude bin.
In both the cluster and control images, an assessment of completeness must
be made. For this, we are using false galaxy simulations. These are
added to cluster and control fields in a systematic manner by running simulations of galaxies of
successive surface brightness and magnitude ranges.
Fiber data was reduced using the HYDRA package in IRAF and radial velocities
were determined via
Fourier cross correlation using the IRAF task FXCOR on absorption line
spectra. Velocity errors
were typically in the range 50-100 km/s although they increased at fainter
magnitudes. Most of the velocities obtained for low surface brightness dwarfs,
however, were through emission line spectra. Each of the five setups
contains at least 100 spectra and at fainter magnitudes the
success rate at determining velocities drops to ~ 20%. We
expect to obtain a greater percentage of reliable velocities with
a higher S/N template in the future.
A summary of the spectral results is given in the table below.
4. Discussion
1. Luminosity function
Figure 3 shows the raw log (number / mag) vs mag counts in
Centaurus along with the background (control field) counts while in figure 4,
background subtracted cluster counts scaled to a square degree in area
are shown.
For most bins, a
significant excess over background counts is observed. A best fit
straight line to the faint magnitudes (Mv > -15) yields a slope of
-0.469.
We have
yet to correct the counts for completeness and it is readily apparent
that incompleteness sets in by m = 22 (Mv = -11). We have also not taken into
account extinction by cluster dust, excess counts due to globular clusters,
differences in large scale structure between the control and cluster
fields, or crowding in the fields. Completeness
is currently being assessed by systematically adding false galaxies in successive
magnitude and surface brightness bins to cluster fields and determining
the recovery efficiency of FOCAS in each magnitude bin.
Three half-magnitude
bins show an excess of galaxies in the controls. This may be due to
FOCAS splitting up large or diffuse galaxies, missing diffuse haloes
altogether, or due to some physical cause such as inauspiciously observing
control fields in a region of higher background density than the
clusters. Of our five control fields, two were thrown out because they
did land on a background cluster.
2.
Spectra
In figures 7 and 8, we show spectra of 5 galaxies (In fig. 8, the spectra correspond to the image 6
numbers:7,20,16). In the images(fig. 5 and
fig. 6),
members are shown in green
rings while higher redshift galaxies are shown in red. Centaurus is at a redshift
of .01. In several cases we find galaxies that would normally, through
morphology, be classified as cluster dwarfs but which prove to be
background galaxies instead. Some of these galaxies are low surface
brightness while others are high surface brightness but M32-like.
Most LSB dwarfs that we were able to obtain velocities for are from
emission line spectra as can be seen here. In figure 6, there is at
least one obvious LSB dwarf that we were unable to obtain a reliable
velocity for. As can be seen in Table 2, at fainter magnitudes and
surface brightnesses there is a sharp drop-off in the number of galaxies
for which were are able to obtain velocities for and we appear to
have reached the limits of the detector by m = 19.
Captions
A field in Centaurus in which several different
galaxy morphologies are visible. Note the large LSB galaxy in the bottom
left.
shows the number of raw counts in the Centaurus cluster fields and
the counts in the control fields (scaled to the total cluster area).
Three bins show an excess of objects in the controls which may be
an error due to FOCAS splitting large and diffuse objects
in this magnitude range. shows the cluster counts backgroundcorrected and scaled to 1 square degree. The best fit slope of the
faint magnitudes (m > 18) is .469 excluding the discrepant bins andthe region m > 22 where incompleteness sets in. We do reach a limitingapparent V magnitude of ~24 and Mv of -8.
Figure 5-8
show several examples of our spectral results. In Figs. 5
and 6, green rings around galaxies represent cluster members while
red rings represent higher redshift galaxies. In Figs. 7 and 8 we show the spectra of 5 of these galaxies. Small bright galaxies often were determined
to be at higher redshifts (ie 20) although this was not always the caseas can be seen in 17 and 13. The larger LSB galaxies that we were ableto obtain velocities for are predominantly from emission lines as is the
case with 23 and 16. For these, the most obvious lines are the [OII]lambda3727
and the triplet of Hbeta and [OIII]\lambda\lambda4959,5007.
