A Photometric and Spectroscopic Study of Dwarf and Giant Galaxies in the Coma Cluster. I....

A PHOTOMETRIC AND SPECTROSCOPIC STUDY OF DWARF AND GIANT GALAXIES IN THE COMACLUSTER. I. WIDE-AREA PHOTOMETRIC SURVEY: OBSERVATION AND DATA ANALYSIS1

Y. Komiyama,2M. Sekiguchi,

3N. Kashikawa,

4M. Yagi,

4M. Doi,

5M. Iye,

4S. Okamura,

6K. Shimasaku,

6N. Yasuda,

4

B.Mobasher,7D. Carter,

8T. J. Bridges,

9and B. M. Poggianti

10

Received 2001 July 30; accepted 2001 September 28

ABSTRACT

We carried out a deep photometric and spectroscopic survey of wide areas in the Coma cluster, aiming toinvestigate the properties of galaxy population in different environments within the cluster. We present theresults in a series of papers. This paper, the first of the series, describes the imaging observations and photo-metric data reduction. Imaging data were taken with the wide-field mosaic CCD camera, which was attachedto the prime focus of the 4.2 m William Herschel Telescope at La Palma, Canary Islands. Our observationscovered a large field of view (2.22 deg2) from the cluster center to the outskirts, and our photometry is com-plete to a limiting magnitude of R ’ 23 mag. The limit of secure star-galaxy discrimination is, however,brighter at R ¼ 20 mag. We identified 3147 galaxies down to this limit in the part (1.32 deg2) of the surveyarea, together with 662 galaxies identified in the control field SA 57. We measured surface photometricparameters and compiled a photometric catalog for these galaxies. Statistical properties of the catalog areshown in this paper, while the catalog itself is given in a forthcoming paper in the series.

Subject headings: galaxies: clusters: general — galaxies: clusters: individual (Coma) —galaxies: photometry

1. INTRODUCTION

The Coma cluster (Abell 1656) is one of the best studiednearby clusters of galaxies. Among nearby clusters, it is therichest and the most compact one showing a reasonabledegree of spherical symmetry. It is classified as type I–II inBautz & Morgan classification and as type B (‘‘ Binary ’’) inRood & Sastry classification because two cD galaxies (NGC4874 and NGC 4889) are located at the center. There arenumerous early works including dynamical studies (e.g.,Zwicky 1937; Rood 1970; Gregory 1975; Kent & Gunn1982), photometric studies (e.g., Strom & Strom 1978; God-win & Peach 1977; Visvanathan & Sandage 1977), and mor-phological and other studies (e.g., Dressler 1980; Butcher &Oemler 1984, 1985). In 1997, a whole conference was dedi-cated to the Coma Cluster (Mazure et al. 1998).

It was in 1983 that a deep and homogeneous photometriccatalog of galaxies in the Coma cluster was published. God-win, Metcalfe, & Peach (1983, hereafter GMP) carried outphotographic photometry and compiled a photometric cat-

alog of 6724 galaxies in a 2=63 square area centered on theComa cluster. The limiting magnitude of the catalog wasbJ ¼ 21, and it was thought to be complete down tobJ ¼ 20. The GMP catalog was the most homogeneous cat-alog with the widest sky coverage at that time, and it becamethe basic source catalog of many studies (e.g., Mazure et al.1988; Mellier et al. 1988; Biviano et al. 1995; Colless &Dunn 1996; Castander et al. 2001).

Since the Coma cluster is more than 5 times as distant asthe Virgo cluster, its dwarf galaxy population has not beenstudied well. The pioneering work on Coma dwarf galaxieswas published by Thompson & Gregory (1993). They usedthe GMP catalog and classified dwarf galaxies morphologi-cally using photographic plates taken with the 4 m telescopeat Kitt Peak in mid-1970s. They presented the spatial distri-bution for each of their morphological types and derivedthe number ratio of dwarf galaxies to giant galaxies. Theyalso presented the luminosity function of each type in thesame manner as Sandage, Binggeli, & Tammann (1985) didfor the Virgo cluster.

CCD imaging of the Coma cluster have been publishedsince 1994 (e.g., Jørgensen & Franx 1994; Kashikawa et al.1995b, 1998; Karachentsev et al. 1995; Bernstein et al. 1996;Ulmer et al. 1997; Lobo et al. 1997a, 1997b; Secker &Harris1997; Trentham 1998). Many of these observations aimed atinvestigating the dwarf galaxy population, which was notstudied in detail by photographic observations. Table 1summarizes the parameters of these photometric observa-tions. They are limited to the central region of the clusterexcept for Kashikawa et al. (1995b, 1998). This is becausesufficient angular resolution is necessary in addition to highsensitivity of CCD to the study of Coma dwarfs. If weassume that Coma dwarfs are similar to Virgo and Fornaxdwarfs, they have scale length of about 100 and are fainterthan R � 16. The image area of a single CCD, even a largeformat CCD (such as 2k � 2k format), is quite small andinefficient to cover the whole extent of the Coma cluster withsufficient angular resolution.

1 Based on observations made with the William Herschel Telescopeoperated on the island of La Palma by the Isaac Newton Group in theSpanish Observatorio del Roque de los Muchachos of the Instituto deAstrofisica de Canarias.

2 Subaru Telescope, 650North A‘ohoku Place, Hilo, HI 96720.3 Institute for Cosmic Ray Research, University of Tokyo, Kashiwa

277-8582, Japan.4 National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka,

Tokyo 181-8588, Japan.5 Institute of Astronomy, University of Tokyo, 2-21-1 Osawa, Mitaka,

Tokyo 181-8588, Japan.6 Department of Astronomy, University of Tokyo, 7-3-1 Hongo, Bun-

kyo, Tokyo 113-0033, Japan.7 Space Telescope Science Institute, 3700 San Martin Drive, Baltimore,

MD 21218.8 Astrophysics Research Institute, Liverpool John Moores University,

TwelveQuaysHouse,EgertonWharf,Birkenhead,Wirral,CH411LD,UK.9 Anglo-Australian Observatory, P.O. Box 296, Epping, NSW 2121,

Australia.10 Osservatorio Astronomico di Padova, vicolo dell’Osservatorio 5,

35122 Padova, Italy.

The Astrophysical Journal Supplement Series, 138:265–278, 2002 February

# 2002. The American Astronomical Society. All rights reserved. Printed in U.S.A.

265

The technique of CCD mosaicking was a breakthroughto this dilemma. Mosaic CCD cameras made observationsof dwarf galaxies feasible beyond a few very nearby clusterssuch as the Virgo and the Fornax clusters. We planned tocarry out a new observational program using the 4.2 mWil-liam Herschel Telescope (WHT) in order to study the prop-erties of galaxies in regions of different density (from thecore to the outskirts) in the Coma cluster. We used a mosaicCCD camera for photometric survey. Spectroscopy wasdone with the multiobject spectrograph, WYFFOS, toincrease spectra of both giant and dwarf galaxies available

to the study. Our photometric survey covered a wide areadown to a considerably deep limiting magnitude comparedto previous observations (see Table 1). In particular, the sur-veyed area is the largest among the recent CCD surveys ofthe Coma cluster. It covers the area for which only photo-graphic data had been available before.

We present the results in a series of papers. This paper,the first of the series, describes the observation and dataanalysis of the wide-area photometric survey. Spectroscopicobservation is presented in Mobasher et al. (2001, hereafterPaper II). Spectral ages and metallicities of non–star-form-ing galaxies are investigated in Poggianti et al. (2001a,2001b, hereafter Papers III and IV). Radial dependence ofgalaxy properties is investigated using spectroscopic data(Carter et al. 2002, hereafter Paper V) and photometric data(Komiyama et al. in preparation). Later papers in the serieswill discuss dynamics, morphology, color-magnitude rela-tion, etc.

The plan of this paper is as follows. In x 2, we introducethe mosaic CCD camera we used for the observation anddescribe the details of the observations. Data reduction pro-cedures and measured photometric parameters are pre-sented in x 3, and a resultant catalog is discussed in detail inx 4.We present statistical properties of the catalog in x 5 andgive a summary in x 6.

2. OBSERVATIONS

We used the Mosaic CCD Camera II (MCCDII), whichwas developed in a collaborative program between the Uni-versity of Tokyo and National Astronomical Observatoryof Japan (Fig. 1; Kashikawa et al. 1995a; Okamura et al.

TABLE 1

Comparison of Recent CCD Observations

Authors

Obs. Area

(deg2) LimitingMagnitude

GMP, Thompson&Gregory (1993)... 6.92 bJ ¼ 21

Kashikawa et al. (1995b, 1998)........... 5.25 R ¼ 18:2a

Karachentsev et al. (1995) .................. 0.0015 R ¼ 25

Bernstein et al. (1996)......................... 0.015 R ¼ 25:5b

Lobo et al. (1997a, 1997b) .................. 0.44 V ¼ 22:5

Secker &Harris (1997) ....................... 0.19 R ¼ 22:5c

Trentham (1998) ................................ 0.19 R ¼ 24

This study .......................................... 1.33 R ¼ 20d

(+Coma 4,5) ..................................... (2.22) (R ’ 23)

Notes.—Summary of recent CCD observations. The photographicobservation byGMP is also listed (top).

a Limiting magnitude of their morphological classification.b At faint magnitude, completeness correction is applied.c Limiting magnitude of their analysis.d Limiting magnitude of secure star-galaxy discrimination.

Fig. 1.—Mosaic CCD Camera II (MCCDII). Forty CCDs are seen aligned in 5� 8 array. A liquid nitrogen tank is attached backside of the Dewar to coolthe CCDs down to appropriate temperature.

266 KOMIYAMA ET AL. Vol. 138

1996). MCCDII has 40 1k � 1k CCDs produced by TI/Japan. Since the CCDs are not buttable, they are aligned inthe 5� 8 array with gaps. The camera is designed so thatfour different exposures provide one contiguous image with-out any gap. The gap between each CCD is about 900 pixels(see Fig. 2), which is optimized to give both as large field ofview as possible and an appropriate overlap (’100 pixels)between adjacent images. The physical pixel size of CCDs is12 lm. CCDs are cooled down to about �80�C by liquidnitrogen so that the dark noise can be almost neglected.Read out noise is measured to be 2–3 ADU, and typicalADC gain is set to be about 3.2 e�/ADU. CCDs and theshutter are controlled by VME-based control system whichis called MESSIA2 (Modularized Expandable System forImage Acquisition; Sekiguchi et al. 1998). A Sparc Work-station communicates withMESSIA2 and controls the cam-era and data acquisitions remotely. The specifications ofMCCDII are summarized in Table 2. Hereafter, we refer tothe data from a single CCD as ‘‘ frame ’’ and the data fromone exposure which consists of 40 1k � 1k frames as‘‘ shot.’’ Accordingly, one contiguous image consists of 4shots, and hence 160 frames.

MCCDII was attached to the prime focus of the 4.2 mWilliam Herschel Telescope (WHT) at Roque de losMuchachos, La Palma, Canary Islands, Spain, in the springof 1996. The scale at the prime focus of WHT is 17>75mm�1, which results in 0>21 per pixel sampling and 50� 32arcmin2 field of view is attained with four exposures whenMCCDII is attached there (see also Fig. 2).

Fig. 2.—CCDgeometry ofMCCDII. Angular scales are calculated for the prime focus ofWHT.

TABLE 2

Specifications of MCCDII

Property Value

CCD chip (pixels)........................ 1018� 1000

Pixel size (lm) ............................. 12

Pixel size (arcsec)......................... 0.21 (WHT-PF)

CCD array .................................. 5� 8= 40

Physical image area (cm2)............ 17.4� 11.1

Physical image area (arcmin2) ..... 50� 32 (WHT-PF)

Read-out noise (ADU)................ 2–3

Dark current ............................... Negligible

CCD gain (e�ADU�1)................ 3.2

Linear range (ADU) ................... <15000

No. 2, 2002 COMA CLUSTER WIDE-AREA PHOTOMETRIC SURVEY 267

Observations were carried out on 1996 April 10 and 14.During the observation, seeing size (FWHM of the point-spread function) varied between 0>8 and 1>5. We plannedto expose five different fields for the Coma cluster from thecenter to the outskirts in two color bands, Johnson B andCousins R bands. However, we could not complete the plansince the weather condition was not perfectly friendly. Werefer to the five fields as Coma 1, which is located at the clus-ter center, and Coma 2 to Coma 5 from the center to the out-skirts (see Fig. 3). The pointing coordinate for each field isdetermined so that there are small overlap regions betweenadjacent fields. All five fields were observed in the R band,however, only Coma 1 and Coma 3 fields were observed inthe B band. We also took exposures of the control field SA57 which is about 3� away from the cluster center. The see-ing size was comparable to the expected scale length ofdwarf ellipticals in the Coma cluster (�100). This does notmean, however, that we can obtain no information on theirluminosity structure. Yagi et al. (2002) performed the simu-lation in which images of some Virgo dwarfs taken with ourmosaic CCD camera in the same run as ours were degraded

to mimic dwarf galaxies located at z ¼ 0:01 0:04. Theyshowed that such degraded galaxies were correctly classifiedas ‘‘ exponential type ’’ if they are brighter thanR ¼ 20 mag.(see x 4.5 of Yagi et al. 2002).

Additional observations were carried out in spring of1998 aiming to extend the survey fields. The B-band imagefor the Coma 2 field was taken under similar seeing size tothat in the 1996 run. Furthermore, external checks on thephotometric zero points were done on 2000 March 16 bytaking calibration data using the Jacobs Kapteyn Telescope(JKT) at La Palma. These data were taken in photometriccondition and cover 11� 11 arcmin2 in the Coma 1 andComa 3 fields in both B and R bands. SA 104 and PG1323�086 compiled by Landolt (1992) were chosen as pho-tometric standard star fields. Details of imaging observationare summarized in Table 3.

In this paper, we describe the method of data reductionand analysis for the Coma 1, Coma 2, Coma 3, and SA 57fields, which we observed in both B andR bands. The result-ing photometric catalog will be published in a separatepaper.

Fig. 3.—Observed fields of our survey. Coma 1 is located at the cluster center and Coma 2 to Coma 5 from center to outskirts. The center of the figure is(13h02m48s, +28�25057>5)J2000.0 and width and height are both 4�. Two cD galaxies, NGC 4889 and 4874, are located at the center of Coma 1 field, and NGC4839 is located at the upper part of the Coma 3 field.


3. DATA REDUCTION PROCEDURES

We have developed a dedicated software package forMCCDII which reduces large amount of data semiauto-matically. This software package is based on AIMS (Auto-mated Image Measuring System; Doi, Fukugita, &Okamura 1993, 1995), which was originally applied to digi-tized data of photographic plates and was improved forCCD data (Yagi et al. 2002).

The software package consists of several componentswhich are specialized for each of the data reduction proce-dures such as bias subtraction, dark subtraction, etc. Thereduction process can be divided into roughly three stages.At the first stage, each frame is reduced individually. Thereduction software used here is almost the same as conven-tional one for single-chip CCD data. Then at the secondstage, image properties of individual frames are measuredand rectified to make up one contiguous image. Finally,objects are detected and photometric and astrometricparameters are measured.We describe below the proceduresused in each of the three stages.

3.1. Reduction of Individual Frames

3.1.1. Bias Subtraction, Corrections, and Dark Subtraction

The variation of the mean counts of bias frames over anight is less than the readout noise. We make the medianbias frame for each chip, which we call ‘‘ mbias.’’ The mbiasframe is subtracted from object frames.

The shutter of MCCDII is of one-wing type. This designwas inevitable because of the limitation of available space atthe prime focus of the WHT. The shutter causes nonuni-formity of exposure time among individual frames in a shot,which has a non-negligible effect in case of short exposure.We correct the nonuniformity by measuring the movingspeed of the shutter wing. The speed was measured and thecorrection factor was calibrated before the observationruns. The photometric error caused by this correction isnegligible for long exposure images.

The prime focus of WHT has optical distortion due to thewide field corrector. The distortion mapping is described asa function of distance from the optical axis,

rtrue ¼ 17:75rfp � 4:1185� 10�5r3fp ; ð1Þ

where rtrue (arcsec) is the angle from the optical axis on thesky and rfp (mm) is the distance from the optical axis at thefocal plane. Hence, 17.75 (arcsec mm�1) corresponds to theapproximate image scale on the focal plane. Object framesare rebinned according to equation (1), i.e., virtual pixel griddescribed by equation (1) is overlaid on the object frame,and photons received at a pixel in the object frame are div-ided into overlaid virtual pixels according to their crosssections.

Dark current due to thermal electrons is negligible forMCCDII since CCDs are sufficiently cooled down. How-ever, it was revealed that limit switches of the shutter hademitted light at 850 nm in the 1996 run. Though a filter wasplaced between these switches and CCDs, small fraction ofemitted light barely penetrated the filter was observed asstray light. The intensity of this light leak was not negligiblecompared with the sky background, about 1/3 of the skylevel in the worst case. The stray light showed characteristicpatterns over the field of view. We took several dark frames,which we call ‘‘ dark exposure,’’ with the shutter opened andthe dome slit closed. The mbias is subtracted from the darkexposures and the median frame of bias-subtracted darkexposures, which we call ‘‘ mdark,’’ is created for each chip.We subtract from all object frames mdark multiplied by aconstant according to the exposure time. Note that theseswitches were replaced before the 1998 run and data takenat 1998 run are free from the stray light.

3.1.2. Flat-Fielding and Sky Subtraction

We assume that the sky is uniform over a single chip andmake a median frame for each chip by taking pixel-to-pixelmedian of many object frames taken with the chip. Thisframe is called ‘‘ mflat.’’ We use not only object frames buttwilight frames which are used instead of dome flat frames.

The local sky is determined and subtracted in each objectframe as follows. First, a frame consisting of 1k � 1k pixelsis divided into meshes. We set the mesh size to be 15>8, i.e.,75 pixels and determine the sky value for each of the meshes.To determine the sky value in a mesh, the histogram of pixelvalues in the mesh is made and this histogram is fitted withGaussian. We regard the center of the Gaussian as the skyvalue for the mesh and the standard deviation as sigma(�SKY) of the sky value. The sky value is assigned to the cen-tral pixel of the mesh. The sky frame is created by bilinearly

TABLE 3

Summary of Imaging Observation

Field Pointing Coordinate (J2000.0) Band

Exposure Time

(s) ObservationDate Comments

Coma 1 ..... 12 59 23.7,+28 01 12.5 R 900 1996 Apr 10 Cluster center

B 1800 1996 Apr 14

Coma 2 ..... 12 57 07.5,+28 01 12.5 R 900 1996 Apr 10

B 1800 1998May 27

Coma 3 ..... 12 57 07.5,+27 11 13.0 R 900 1996 Apr 10 NGC 4839 group

B 1800 1996 Apr 14

Coma 4 ..... 12 57 07.5,+28 50 42.5 R 900 1996 Apr 10/14

Coma 5 ..... 12 57 07.5,+29 40 12.5 R 900 1996 Apr 14

SA 57 ........ 13 09 46.6,+29 23 02 R 900 1996 Apr 10 Control field

B 1200 1996 Apr 14

Note.—Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, andarcseconds.


interpolating these sky values assigned to the central pixelsof all the meshes. Finally, the sky frame is subtracted fromthe object frame.

3.1.3. Coarse Detection, PSFMeasurement, and CosmicRay Removal

Coarse detection of objects is carried out to select objectswhich are used for the measurement of the point-spreadfunction (PSF). The essence of the detection method is socalled ‘‘ connected pixels ’’ algorithm. We regard a lump ofconnected pixels as an object when the lump satisfies the fol-lowing criteria; (1) all the connected pixels have pixel valueslarger than the detection threshold and (2) the number ofconnected pixels is bigger than the threshold. The detectionthreshold is set at 1:5� �SKY. The minimum number of pix-els is set to be 30 pixels which is almost equal to the seeingsize. If we take smaller values, the contamination from cos-mic rays and defects of CCDs becomes significant. If wetake larger values, on the other hand, we miss small faintobjects.

Starlike objects are selected from the detected objects onthe assumption that stellar profile is sharp and circular. Toobtain accurate shape of the PSF, starlike objects are co-added to reduce the Poisson noise. However, the centers ofstarlike objects are not necessarily coincident with the centerof a pixel, the profile would be smeared if the objects aresimply co-added without taking this offset into account.How to obtain an accurate PSF is described in Komiyama,Doi, & Okamura (2001).

Once the PSF is determined, objects whose profiles aremuch sharper than the PSF are regarded as cosmic rays ordefects, and subtracted from each frame. Pixels regarded ascosmic rays or defects are replaced by the weighted meanvalue of 5� 5 pixels around them.

3.2. Making up a Contiguous Image

3.2.1. Smoothing and Object Detection

To make a contiguous image, four different shots are nec-essary. The seeing size may vary from shot to shot accordingto the weather condition. Even if the weather condition isstable during the observation, PSFs may be different fromframe to frame within a shot due to a small tilt between thecamera and the focal plane of the telescope. It is necessaryto unify PSFs of all the frames to secure homogeneous sur-face photometry over the contiguous field.

The PSF at the focus of large telescopes is known to berepresented well by the double Moffat function or approxi-mately double Gaussian (Racine 1996). The main compo-nent, which is caused by an atmospheric turbulence, has anarrower wing than the minor component due to instru-ments. The narrow component contributes to about 86% ofthe central intensity, and it is this narrow component thataffects the homogeneity of the photometry very much. Thewide and weak component has little effect because it isalready much wider than the seeing size. Accordingly, weassume the PSF to be a single Gaussian for simplicity andconvolve each frame with a Gaussian whose sigma isð�2

target � �2frameÞ

1=2 to have a smoothed image, where �frame

is the sigma of the PSF measured for each frame. We set�target to be 3.1 (pixels), which is the largest sigma of the PSFamong all frames in both B and R bands. Degrading all theframes to the worst seeing apparently looks absurd. This is,however, essential to homogeneous surface photometry.

Then objects are detected again. At this time, detectionthreshold is set to be the same level for all the frames. Wemeasure fixed aperture flux in addition to the position forthe detected objects. They are used in the followingmosaicking procedure.

3.2.2. Mosaicking and Additional Flux Correction

Mosaicking is among the most important steps in thepresent data reduction, where a consistent coordinate sys-tem and a consistent flux scale should be established overthe contiguous image.

Usually the objects detected in the overlap region betweenadjacent chips/shots are used to calculate geometricalparameters (relative shift and rotation) and relative sensitiv-ities. However, there are difficulties here, due to the overlapregion being quite narrow (about 66 arcmin2) and the highresolution, 0>21 pixel�1, of MCCDII at the prime focus ofWHT. The region is further reduced due to the optical dis-tortion of the focal plane, especially for those chips locatedaway from the optical axis. In some cases, there are only afew or no objects in the overlap region. Geometrical param-eters in such cases are computed as follows. The relativelocation and rotation of the 40 chips remain unchanged dur-ing the exposure because the CCD chips are glued on thecold plate of MCCDII. These geometry data are knownfrom both the laboratory measurements and previousobservations. The average shift between different shots canbe computed by using all the overlap regions where someobjects are always detected. Accordingly, the geometricalparameters can be computed by combining the two sets ofdata, i.e., relative values in a shot and the average shiftbetween different shots.

Relative sensitivities are more difficult to estimate thangeometrical parameters. At the faint levels, one couldalways find objects in the overlap region. However, it is dan-gerous to rely on such faint objects for computing relativesensitivities. Instead, we use the sky brightness to computerelative sensitivities as explained below.

Basic assumption is that relative sensitivity of CCD chipsis constant. Let ri be the relative sensitivity of CCD chip iand 1� dj be the fraction of flux absorbed by the atmo-sphere for shot j. We refer to dj as the transparency of shot j.If the weather condition is good, dj depends only on the airmass of shot j [dj ¼ djðsecZÞ / secZj]. Then an objectwhose real flux, i.e., flux outside of the atmosphere, is F willbe observed by chip i at shot j as

fi; j ¼ ridjF : ð2Þ

Here ri is calculated accurately in advance using sky frames(object sky frames and twilight sky frames) on the assump-tion that the sky is uniform over the field of view when aver-aged in the exposure time.11

Note that ri is normalized by the average of ri and can bedetermined with 1% and 3% accuracies for R and B bands,respectively. Relative transparency, dj=dk, is independent of

11 The sky is both spatially and temporally variable due to airglow onscales of 1 degree and minutes, in addition to the slow variation on scales ofyears due to the solar cycle. Variation is larger for longer wavelengths. Thespatial variation of short timescales introduces some error in our zeropoint, especially in the R band. However, we consider this extra error quitesmall because the variation was averaged over 15 minutes, which is longerthan the timescale of variation.


F and calculated using common objects detected in the over-lap region as follows:

djdk

¼fðlÞi; j =ri

fðlÞi;k =ri

; ð3Þ

where fðlÞi; j and f

ðlÞi;k are measured with common object l

detected in the overlap region. Averaging over l objectsgives more accurate results. Since dj=dk is, in practice, con-stant for all the chips, dj=dk can be estimated even if thereare no common object in the overlap region between shots jand k of chip i. The ratio dj=dk is averaged over differentchips i to reduce the uncertainty as follows:

djdk

¼ 1

Nchip

XNchip

i

1

li

Xlil

fðlÞi; j =ri

fðlÞi;k =ri

!: ð4Þ

Absolute values of dj are normalized by the average of dj.Accordingly, the relative sensitivity of chip i in shot j isdetermined as ridj.

3.3. Constructing a Contiguous Image

Now that we have geometrical parameters and relativesensitivities for all the chips and shots, we can produce acontiguous image. Care is taken so that the errors are dis-tributed almost equally over the contiguous image. A pre-liminary photometric catalog is made followed by thedetection of objects and the measurement of parameters.This preliminary catalog is used in the absolute calibrationsdescribed below.

3.4. Absolute Calibrations

3.4.1. Astrometric Calibration

We adopt the APM (Automatic Plate Measuring)machine Catalog12 (Irwin &McMahon 1992; Maddox et al.1990) compiled by Royal Greenwich Observatory to selectastrometric standard objects to translate the positions ofobjects into celestial coordinates. It covers almost all overthe sky down to BJ � 22 mag with the objects classified intofour categories (stellar, nonstellar, blended, and noiselike).Its astrometry is based on PPM catalog (Rosner & Bastian1991) and attains 0>5 accuracy for PPM stars.

Stars which are not saturated and are not too faint in thepreliminary catalog are selected and compared with those inthe APM catalog. We matched typically 700–1000 stars percontiguous field. The accuracy of astrometry is estimated tobe ’0>5 rms based on these matched stars (Fig. 4), and itdoes not depend onmagnitude.

3.4.2. Photometric Calibration

Photometric zero points for the Coma 1 and Coma 3fields are calibrated using the data taken in the year 2000, asdescribed in x 2. There are 3–5 photometric standard stars ineach frame. These calibration data yield the photometriczero points for the Coma 1 field accurate to 0.02 mag in bothR and B bands. The corresponding values for the Coma 3field are 0.03 and 0.04 mag inR and B bands, respectively.

Coma 2 field is calibrated relative to Coma 1 field usingcommon objects in the overlap region between these two

fields. We carry out 1000 aperture photometry for theseobjects and determine the photometric zero point for theComa 2 field so that the average magnitude difference ofthese object is equal to zero. Figure 5 shows the resultingmagnitude difference of common objects which are detectedin the overlap region between Coma 1 and Coma 2 fields.

This figure is also useful to estimate the quality of thedata. The deviation around the average represents the errorscaused by data reduction procedures such as flat-fielding,determination of sky, determination of relative sensitivity ofthe CCD chip, and flux correction. We refer to these errorscollectively as the ‘‘ internal error ’’ (�int). The internal erroris thought to represent the quality of the data since it cannotbe reduced even if zero point magnitude is determined withno error.

The standard deviation of magnitude difference forobjects in the magnitude range of 16 < R < 20 is’0.04 mag12 See http://www.ast.cam.ac.uk/�apmcat/.

-1

0

1

2

-1

0

1

2

-1

0

1

2

-2 -1 0 1 2

-1

0

1

2

-1 0 1 2

Fig. 4.—Accuracies of astrometric calibrations for each field. The leftpanels are for theR band, and the right panels are for theB band.


in the R band. Accordingly, �int of R-band images is calcu-lated to be 0.04/

ffiffiffi2

p¼ 0:03 mag, assuming that the internal

errors for both Coma 1 and Coma 2 fields follows the nor-mal distribution. The corresponding number for the B bandis calculated to be 0.04 mag for the magnitude range of18 < B < 21.

3.5. Surface Photometry

Objects are detected, and photometric and geometricalparameters are measured using the contiguous image takenin theR band. We set the detection threshold to be 25.5 magarcsec�2 in the R band, which corresponds to about 1.5�SKY. We set the minimum number of pixels explained inx 3.1.3 to be 30 pixels.

The software module called ‘‘ measure ’’ performs theobject detection and parameter measurements. It also car-ries out deblending when necessary. The performance of thedeblending algorithm is studied by Yagi et al. (2002). Thebasic method of deblending is the multipeak finding, and itis based on the assumption that objects have 180� symmetry(see Yagi et al. 2002 for detailed description). We make thefinal catalog based on R-band images, and B-band imagesare used when wemeasure parameters concerned with color.This is because matching the same galaxies in different cata-logs independently constructed from the data in differentbands is quite tricky. Issues which often cause troublesinclude different deblending, different positions due to dif-ferent appearance of the galaxy, and different limitingmagnitude.

The module ‘‘ measure ’’ often uses the ‘‘ equivalent pro-file.’’ Here we write down briefly the definition of the equiv-alent profile. Let SðIÞ be the area which has the intensityabove I including ‘‘ islands,’’ if any, the equivalent radius r�

is defined as r� ¼ SðIÞ=�½ �1=2. Then, the equivalent profile isrepresented as Iðr�Þ.

The ‘‘measure ’’ module measures the following parame-ters for deblended images. Note that isolated images are leftas they are by the deblender and the parameters are mea-sured for the original images. (i) Intensity-weighted imagecenter (xc, yc). (ii) Isophotal magnitude (miso) defined astotal magnitude within detection threshold. (iii) Number ofpixels within the detection threshold (Npix). (iv) Peak inten-sity defined by the value of the most luminous pixel value(Ipeak). (v) Ellipse parameters, i.e., major axis (a), minor axis(b), axis ratio (q ¼ b=a), and position angle (P.A.). Theseparameters are derived using intensity-weighted secondorder moments. (vi) Petrosian radius (rP). It is defined bythe ratio of the local surface brightness at the radius, IðrPÞ,to the mean surface brightness within the radius, LðrPÞ=�r2P,as follows:

� ¼ IðrPÞLðrPÞ=�r2P

; ð5Þ

where IðrÞ is the intensity at radius r and LðrÞ is the inte-grated luminosity within the circle of radius r. We set � to be0.25. Petrosian radius is calculated based on the equivalentprofile. (vii) Kron radius (rK). It is calculated as

rK ¼

XiriIðriÞXiIðriÞ

: ð6Þ

Summation is taken over all the pixels within the detectionthreshold (Kron 1980). (viii) Kron magnitude (mK) definedby

mK ¼ �2:5 logLð< 3rKÞ ; ð7Þ

where Lð< 3rKÞ is the luminosity within the circle of radius3rK. The fraction of the luminosity enclosed in a circularaperture of radius 3rK is expected to be more than�95% forboth de Vaucouleurs profile and exponential profile (Bertin& Arnouts 1996). We consider mK as the substitute for thetotal magnitude. (ix) Effective parameters (re and �e). Theseare calculated on the assumption that Kron magnitude isequal to the total magnitude. Effective radius is defined asthe radius within which half of the total luminosity isincluded and effective surface brightness is defined by theaverage surface brightness within re. (x) The profile shapeparameter (N) derived by fitting the Sersic law (Sersic 1968),

IðrÞ ¼ I0 exp � r

r0

� �N" #

; ð8Þ

to the radial profile. We use the equivalent profile to fitequation (8). To avoid the effect of seeing, central pixels(r < 3�PSF) are excluded from the fitting. The outer part(r > 2:5re) is also excluded to avoid the effect of excessivenoise. Hereafter, N is used as the indicator of the profileshape. Note that the de Vaucouleurs law profiles haveN ¼ 0:25, and the exponential law profiles haveN ¼ 1:0.

3.5.1. Measurement of Color Parameters

The colors are estimated using aperture photometry. Theaperture center is first determined by using the R-bandimage. Then, the coordinate of the center in the R-bandimage is translated to the coordinate in the B-band image

Fig. 5.—Magnitude differences between Coma 1 and Coma 2 fields in B(top) andR (bottom) bands.


via celestial coordinate, and aperture photometry is carriedout on the B-band image.

The aperture fluxes are measured with the original (non-deblended) images. This is because deblending may givedifferent results between R and B bands. Accordingly,uncertainty is large for deblended objects compared withisolated objects. Furthermore, colors measured with smallerapertures may suffer from larger error which comes fromthe uncertainty that the center of aperture is not exactlycoincident betweenR and B bands due to astrometric uncer-tainties. Circular apertures are used for all objects.

3.5.2. Distance from the Cluster Center

The structure of the Coma cluster may not be simple;there are two cD galaxies at the central region, and the X-ray peak position is coincident with one of the cD galaxies(e.g., White, Briel, & Henry 1993). In this study, we adopt(12h59m42 98, +27�58014>00)J2000.0 as the center of the clus-ter following the definition of GMP. This center is locatedbetween two central cD galaxies, and distance from this cen-ter is calculated for each galaxy.

4. CONSTRUCTION OF THE CATALOG

4.1. Galaxy Extraction

Objects detected by ‘‘measure ’’ consist of galaxies, stars,and false objects such as ghosts, defects, and even satellitetrails. The number of stars is not negligible compared withthe number of galaxies, and discrimination of galaxies fromother objects is necessary. In this study, we adopt the follow-ing discrimination method, which is similar to that used byYagi et al. (2002).

First, the Petrosian radius rP is used as the discriminatorof stellar objects. The radius rP is calculated to be 1.9 � forthe Gaussian PSF, and rP of stellar objects are expected todistribute around this value. In the upper panel of Figure 6,the logarithm of rP=�PSF is plotted against isophotal magni-tude in theR band (Riso).We see a stellar locus running hori-zontally at logðrP=�PSFÞ � 0:35. The locus merges at faintmagnitude (Riso � 21:5 mag) with galaxies and is bent atbright magnitude (Riso � 16 mag) where effect of saturationof CCD begins to change the stellar profile shape.We regardthose objects that have smaller rP than 3 � of the distribu-tion as stars. In the figure, this boundary line is drawn atlogðrP=�PSFÞ ¼ 0:39. For the brighter part, we regard thosehaving smaller rP compared with galaxies as saturated stars.This criterion is represented by the diagonal dashed line inthe figure.

Objects regarded as galaxies here are contaminated byclose double stars which have larger rP. Those stars are iden-tified by another parameter, Ipeak=Npix. In the lower panel ofFigure 6, the logarithm of Ipeak=Npix is plotted against iso-photal magnitude in the R band (Riso). Objects regarded asgalaxies by the rP criterion above are plotted as filled circles.We see a stellar locus running diagonally in this diagram.Close double stars are embedded in the locus. We set criticalIpeak=Npix, and those having higher Ipeak=Npix are excluded.The critical value depends on the seeing size and is deter-mined by inspection of images.

Highly elongated objects such as large cosmic rays, satel-lite images, etc., are rejected by examination of the axisratio, q ¼ b=a. We set the criterion to be q ¼ 0:1 and rejectthose objects that have q less than this value. The number of

objects rejected by this criterion is small compared to thoseby previous criteria.

In this study, we set limiting magnitude of star-galaxy dis-crimination to be Riso ¼ 21 mag because the primary aim ofthe present photometric sample is to supplement the spec-troscopic sample whose limiting magnitude is B ¼ 21 mag.The threshold Riso ¼ 21 mag is really bright compared tothe limiting magnitude of detection (Riso ’ 24 mag), and weare confident about our star-galaxy discrimination down tothis limit. Future studies of fainter objects, however, shouldbe carried out by taking care not to miss galaxies fainterthan Riso ¼ 21 by pursuing more sophisticated star-galaxydiscrimination.

Figure 7 shows RK (Kron magnitude in the R band)plotted as a function of Riso. As expected, RK is system-atically brighter than Riso by �0.3 mag. We adopt a con-servative limit RK ¼ 20 mag for galaxies to be includedin the present photometric catalog. Numbers of galaxiesbrighter than RK ¼ 20 in each field are 1396, 1085, 966and 662 for Coma 1, Coma 2, Coma 3, and SA 57,respectively.

4.2. Notes on the Catalog

4.2.1. Effect of Saturation

We regard a pixel whose count is larger than 15000 ADU(nonlinearity 3%) as a saturated pixel. As far as stellarobjects are concerned, saturation affects very much forR < 17 and B < 17 objects. In Figure 6, it is clearly seen thatstellar locus is bent at R ’ 17 by the effect of saturation.However, it is difficult to estimate how badly a galaxy is

0.5

1

1.5

10 15 20-2

-1

0

1

2

Star Galaxy Discrimination Diagram : coma1r

Fig. 6.—Galaxies are discriminated from other objects using these dia-grams.Top: logarithmic Petrosian radius over PSF sigma of all the detectedobjects plotted against isophotal magnitude. Stellar locus is clearly seenrunning horizontally. Bottom: Peak value divided by the number of pixelsabove the detection threshold (Ipeak=Npix) of all the detected objects plottedagainst isophotal magnitude. Objects regarded as galaxies by the rP crite-rion are plotted as filled circles.


affected by saturation since galaxies have various radial pro-files and various sizes. In fact, most galaxies brighter thanR ’ 14 have peak values larger than 15000 ADU. It is betterto regard those galaxies brighter than R � 14:5 as galaxiesaffected by saturation.

How largely the photometric parameters are affected bysaturation is also difficult to assess. If only one central pixelis saturated, it will not change parameters significantly. Onthe contrary, if there are many saturated pixels at the centralregion of galaxies, parameters for such galaxies are badlyaffected by saturation. Parameters derived from profile fit-ting are thought to be least affected by saturation since cen-tral pixels are excluded from the fitting. Effective radius andeffective surface brightness are thought to be most affectedsince they are calculated using all the pixels above the detec-tion threshold including central pixels. We should be carefulabout parameters of galaxies whose magnitudes are brighterthanR � 14:5.

4.2.2. Effect of Bright Stars

There are several bright stars (typically brighter than 12mag) in the images which make a saturation tail and diffusehalo around them. Hence, in the vicinity of these stars,photometry for a galaxy is difficult and not so accurate. In

Fig. 7.—Kron magnitude in the R band (RK) is plotted as a function ofisophotal magnitude in theR band (Riso).

Fig. 8.—Magnitude distribution of cataloged galaxies in each field


the most difficult case, galaxies which are superposed onboth the saturation tail and diffuse halo may not bedetected. Therefore, the catalog is not complete aroundthese very bright stars. Except for this ‘‘ vicinity effect,’’ thepresent catalog is supposed to be complete down toRK ¼ 20 mag.

5. STATISTICAL PROPERTIES OF THE CATALOG

5.1. Magnitude and Spatial Distributions

Figure 8 shows the magnitude distributions of catalogedgalaxies in each field (Coma 1, Coma 2, Coma 3, and SA57). The detailed numbers are listed in Table 4. As is seen,the number of galaxies decreases with the distance from thecluster center (from Coma 1 to Coma 2, and to Coma 3) andbecomes the smallest at SA 57, which is supposed to be ablank field.

In Figures 9 and 10, spatial distributions of cataloged gal-axies in four different magnitude ranges (RK < 16:3,16:3 < RK < 18, 18 < RK < 19, and 19 < RK < 20) areplotted. Open circles represent that they are isolated gal-axies and filled circles represent that they are deblended gal-axies. Three cD galaxies, NGC 4874, NGC 4889, and NGC4839 are plotted as large crosses. It is clear that brighter gal-axies tend to concentrate at the cluster center, while noobvious concentration can be seen for fainter galaxies.

5.2. Photometric Parameters as a Function ofMagnitude

Figure 11 shows the effective surface brightness as a func-tion of magnitude for galaxies in the Coma field (Coma 1,Coma 2, and Coma 3 fields altogether) and the SA 57 field.The distribution for galaxies in the Coma field at a givenmagnitude is wider than that in SA 57 field, especially it isextended toward lower surface brightness. This implies thatthese low surface brightness galaxies in the Coma fieldwhich are not seen in SA 57 field are cluster member gal-

TABLE 4

Numbers of Cataloged Galaxies in Each Field

RK Coma 1 Coma 2 Coma 3 SA 57

12.0–12.5..... 1 0 0 0

12.5–13.0..... 3 0 2 1

13.0–13.5..... 11 3 2 0

13.5–14.0..... 15 4 2 1

14.0–14.5..... 20 7 10 1

14.5–15.0..... 22 12 14 3

15.0–15.5..... 34 10 10 3

15.5–16.0..... 25 16 12 3

16.0–16.5..... 34 20 14 5

16.5–17.0..... 41 35 18 4

17.0–17.5..... 69 43 29 16

17.5–18.0..... 90 65 56 29

18.0–18.5..... 121 103 100 43

18.5–19.0..... 203 145 152 107

19.0–19.5..... 282 263 208 184

19.5–20.0..... 425 359 337 262

Total ........... 1396 1085 966 662

Fig. 9.—Spatial distributions of cataloged galaxies in the magnitude ranges of RK < 16:3 (left) and 16:3 < RK < 18 (right). Open circles represent isolatedgalaxies, and closed circles represent deblended galaxies. Large crosses indicate positions of three cD galaxies, NGC 4874, NGC 4889, andNGC 4839.


axies. Note, however, that not all cluster member galaxiesare low surface brightness galaxies; there are cluster membergalaxies which have high surface brightness.

Figure 12 shows effective radius as a function ofmagnitude for galaxies in the Coma field and the SA 57 field.Galaxies which have large effective radii in the Comafield correspond to low surface brightness galaxies seen inFigure 11.

Figure 13 shows profile shape parameter as a function ofmagnitude for galaxies in the Coma field and the SA 57 field.The distribution becomes wider as magnitude goes fainter,especially the number of large-N galaxies increases. This isconsistent with the result that dwarf elliptical galaxies preferexponential radial profiles rather than de Vaucouleurs pro-files. Similar distribution is also seen in the Virgo cluster(e.g., Binggeli & Jerjen 1998).

Figure 14 is a color-magnitude diagram for galaxies in theComa field and the SA 57 field. The color-magnitude rela-tion of giant elliptical galaxies (13.RK . 16) is clearly seenin the Coma field, and it extends to fainter magnitude. Colorcriteria for spectroscopic sample selection are based on thisdiagram, as is described in Paper II.

6. SUMMARY

We carried out a deep photometric and spectroscopic sur-vey of wide areas in the Coma cluster region. The surveydata are useful to investigate the stellar population proper-ties of galaxies in a wide magnitude range and the photo-metric and spectroscopic properties of galaxy population in

Fig. 10.—Same as Fig. 9 but for 18 < RK < 19 (left) and 19 < RK < 20 (right)

Fig. 11.—Effective surface brightness as a function of Kron magnitudefor galaxies in the Coma fields (top) and the SA 57 field (bottom). Large dotsrepresent isolated galaxies, and small dots represent deblended galaxies.


different environments within the cluster. We present theresults in a series of papers.

This paper summarizes the imaging observations andphotometric data reduction procedures of our survey. Wecarried out imaging observation with Mosaic CCD CameraII (MCCDII), which was attached to the prime focus of the4.2 m William Herschel Telescope (WHT) at La Palma,Canary Islands. The pixel scale at the prime focus of WHTis 0>21 and 50� 32 arcmin2 field of view is attained withfour different exposures.

Our observations covered a large field of view (2.22 deg2)from the cluster center to the outskirts (see Fig. 3), and ourphotometry goes down to a significantly deep limiting mag-nitude (RK ’ 23 mag). In this paper, we analyzed part ofthe survey area (1.32 deg2) where both R and B band imagesare available. We identified 3147 galaxies in this subareadown to RK ¼ 20 mag, which is the limit of secure star-gal-axy discrimination. The data are accurate to less than 0.1mag in magnitude and 0>5 in position. We carried out sur-face photometry and compiled a photometric catalog forthese 3147 galaxies together with 662 galaxies identified inthe control field SA 57. Distributions of photometricparameters, such as effective surface brightness, effectiveradius, profile shape parameter, and color, are plotted as afunction of magnitude.

Observations on which the series of papers are based werecarried out with the 4.2 m William Herschel Telescope andthe Mosaic CCD camera in a collaborative program

Fig. 12.—Effective radius as a function of Kron magnitude for galaxiesin the Coma fields (top) and the SA 57 field (bottom). Large dots representisolated galaxies, and small dots represent deblended galaxies.

Fig. 13.—Profile shape parameter as a function of Kron magnitude forgalaxies in the Coma fields (top) and the SA 57 field (bottom). Large dotsrepresent isolated galaxies, and small dots represent deblended galaxies.

Fig. 14.—Color-magnitude diagrams for galaxies in the Coma fields(top) and the SA 57 field (bottom). Large dots represent isolated galaxies,and small dots represent deblended galaxies.


between Japan and UK under the Memorandum of Under-standing between Japan and the United Kingdom Concern-ing Cooperation in Ground-Based Astronomy. We thankthe observatory staff of the Isaac Newton Group for theirexcellent technical support which made MCCDII observa-

tion with the WHT possible. We are grateful to the referee,Prof. Greg Bothun, for his many valuable comments andsuggestions. Y. K. acknowledges the Hayakawa Fund ofthe Astronomical Society of Japan for the support of histravel.

REFERENCES

Bernstein, G. M., Nichol, R. C., Tyson, J. A., Ulmer, M. P., & Wittmann,D. 1996, AJ, 110, 1507

Bertin, E., &Arnouts, S. 1996, A&AS, 117, 393Binggeli, B., & Jerjen, H. 1998, A&A, 333, 17Biviano, A., Durret, F., Gerbal, D., Le Fevre, O., Lobo, C., Mazure, A., &Slezak, E. 1995, A&AS, 111, 265

Butcher, H. R., &Oemler, A. Jr. 1984, ApJ, 285, 426———. 1985, ApJS, 57, 665Carter, D., et al. 2001, ApJ, in press (Paper V)Castander, F. J., et al. 2001, AJ, 121, 2331Colless,M., &Dunn, A.M. 1996, ApJ, 458, 435Doi,M., Fukugita,M., & Okamura, S. 1993,MNRAS, 264, 832———. 1995, ApJS, 97, 59Dressler, A. 1980, ApJ, 236, 351Godwin, J. G., Metcalfe, N., & Peach, J. V. 1983, MNRAS, 202, 113(GMP)

Godwin, J. G., & Peach, J. V. 1977,MNRAS, 181, 323Gregory, S. A. 1975, ApJ, 199, 1Irwin,M., &McMahon, R. 1992, Gemini, 37, 1Jørgensen, I., & Franx,M. 1994, ApJ, 433, 553Karachentsev, I. D., Karachentsev, V. E., Richter, G. M., & Vennik, J. A.1995, A&A, 296, 643

Kashikawa, N., Sekiguchi, M., Doi, M., Komiyama, Y., Okamura, S.,Shimasaku, K., Yagi,M., &Yasuda, N. 1998, ApJ, 500, 750

Kashikawa, N., Sekiguchi, M., Yagi, M., Yasuda, N., Okamura, S.,Shimasaku, K., &Doi,M. 1995a, in IAU Symp. 167, NewDevelopmentsin Array Technology and Applications, ed. D. A. Philip (Dordrecht:Kluwer), 345

Kashikawa, N., Yagi, M., Yasuda, N., Okamura, S., Shimasaku, K., Doi,M., & Sekiguchi,M. 1995b, ApJ, 452, L99

Kent, S.M., &Gunn, J. E. 1982, AJ, 87, 945Komiyama, Y., Doi,M., &Okamura, S. 2001, PASP, submittedKron, R. G. 1980, ApJS, 43, 305Landolt, A. U. 1992, AJ, 104, 340

Lobo, C., Biviano, A., Durret, F., Gerbal, D., Le Fevre, O., Mazure, A., &Slezak, E. 1997a, A&A, 317, 385

———. 1997b, A&AS, 122, 409Maddox, S. J., Efstathiou, G., Sutherland, W. J., & Loveday, J. 1990,MNRAS, 243, 692

Mazure, A., Casoli, F., Durret, F., & Gerbal, D. 1998, A New Vision of anOld Cluster: Untangling ComaBerenices (Singapore: World Scientific)

Mazure, A., Proust, D.,Mathez, G., &Mellier, Y. 1988, A&AS, 76, 339Mellier, Y., Mathez, G., Mazure, A., Chauvineau, B., & Proust, D. 1988,A&A, 199, 67

Mobasher, B., et al. 2001, ApJS, 137, 279 (Paper II)Okamura, S., et al. 1996, J. KoreanAstron. Soc., 29, S375Poggianti, B.M., et al. 2001a, ApJ, 562, 689 (Paper III)Poggianti, B.M., et al. 2001b, ApJL, 563, 118 (Paper IV)Racine, R. 1996, PASP, 108, 699Rood, H. J. 1970, ApJ, 162, 333Rosner, S., & Bastian, U. 1991, PPM Star Catalog, Publ. Astron. Rechien-Institut Heidelberg

Sandage, A., Binggeli, B., & Tammann, G. A. 1985, AJ, 90, 1759Secker, J., &Harris, W. E. 1997, PASP, 109, 1364Sekiguchi,M., Nakaya, H., Kataza, H., &Miyazaki, S. 1998, Exp. Astron.,8, 51

Sersic, J. L. 1968, Atlas de Galaxias Australes (Cordoba: Obs. Astron. Cor-doba)

Strom, K.M., & Strom, S. E. 1978, AJ, 83, 73Thompson, L. A., &Gregory, S. A. 1993, AJ, 106, 2197Trentham, N. 1998,MNRAS, 293, 71Ulmer, M. P., Bernstein, G. M., Martin, D. R., Nichol, R. C., Pendleton,J. L., & Tyson, J. A. 1997, AJ, 112, 2517

Visvanathan, N., & Sandage, A. 1977, ApJ, 216, 214White, S. D.M., Briel, U. G., &Henry, J. P. 1993,MNRAS, 261, L8Yagi, M., Kashikawa, N., Sekiguchi, M., Doi,M., Yasuda, N., Shimasaku,K., &Okamura, S. 2002, AJ, in press

Zwicky, F. 1937, ApJ, 86, 217

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