Elizabeth Rizza1,2. Department of Physics and Astronomy, University of Missouri, Columbia, MO 65211; and Department of Astronomy,. New Mexico State ...
The Astronomical Journal, 126:119–142, 2003 July # 2003. The American Astronomical Society. All rights reserved. Printed in U.S.A.
SENSITIVE RADIO AND OPTICAL OBSERVATIONS OF z 0.2 RICH ABELL CLUSTERS Elizabeth Rizza1,2 Department of Physics and Astronomy, University of Missouri, Columbia, MO 65211; and Department of Astronomy, New Mexico State University, Las Cruces, NM 88003
Glenn E. Morrison2 Infrared Processing and Analysis Center, California Institute of Technology, Pasadena, CA 91125
Frazer N. Owen2,3 National Radio Astronomy Observatory, P.O. Box O, Socorro, NM 87801
Michael J. Ledlow2 Gemini Observatory, AURA, Casilla 603, La Serena, Chile
Jack O. Burns Center for Astrophysics and Space Astronomy, University of Colorado, Boulder, CO 80309-0839
and John Hill Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721 Received 2002 October 22; accepted 2003 April 1
ABSTRACT We present optical observations and radio continuum imaging data for a sample of rich, X-ray–bright Abell clusters at intermediate (z 0.2) redshifts. We find that the radio galaxy population varies substantially from cluster to cluster within this homogeneous sample. The spatial distribution of the high-luminosity radio galaxies (HLRGs; L1.4 > 1023 W Hz1) is very different from the low-luminosity radio galaxies (LLRGs; L1.4 1022.75 W Hz1), with the LLRGs displaying a flat spatial distribution in contrast to the centrally peaked HLRGs. A color-morphology classification shows that the HLRGs are composed primarily of galaxies with old stellar populations, whereas the LLRGs have a much more diverse composition. We do not see a correlation between the cluster radio fraction and cluster blue fraction. However, there is a moderate anticorrelation with richness, suggesting that a rich cluster is less likely to have radio-bright galaxies, whether the radio emission is due to active galactic nuclei or star formation. Key words: galaxies: active — galaxies: clusters: general — galaxies: starburst — radio continuum
A recent study using radio analysis to identify star-forming galaxies in two intermediate-redshift clusters (Owen et al. 1999; Owen & Dwarakanath 1999) found two strikingly different radio populations. The z = 0.25 cluster A2125 has a high blue fraction, according to Butcher & Oemler (1984), and complex X-ray morphology believed to be the result of a recent or ongoing major merger (Wang, Connolly, & Brunner 1997). The radio population also proved to be quite remarkable, consisting of 26 radio identifications yielding a radio fraction (the number of radio galaxies with MR < 21 divided by the number of cluster galaxies with MR < 21) of fRG = 0.19 (Owen et al. 1999). In contrast, a comparable mass cluster at z = 0.25, Abell 2645, has a very low blue fraction ( fB = 0.03) and low radio fraction ( fRG = 0.01). These two test cases suggested a possible connection between the BO mechanism (Butcher & Oemler 1984) that instigated star formation and blue galaxy colors and the mechanisms that initiated radio activity in the cluster galaxies. The follow-up studies by Morrison (1999, hereafter M99), Morrison et al. (2003b, hereafter R1), and Morrison & Owen (2003, hereafter R2) did not show a close tie between the blue fractions and radio fractions for their sample of z = 0.15–0.4 clusters, yet it did show an intriguing evolution in the radio galaxy population as a function of redshift. M99 found a significant increase in the number of
1. INTRODUCTION
Radio observations have become an increasingly important tool in the detection and study of star formation in clusters of galaxies. At 20 cm the radio luminosity function above 1023 W Hz1 is dominated by active galactic nuclei (AGNs), while below this power the majority of the sources are star-forming systems (e.g., Condon, Cotton, & Broderick 2002). The radio lifetimes are relatively short (108 yr) compared with typical crossing times of a cluster galaxy (109 yr) and may provide a unique signpost to recent and/or ongoing star formation. The radio sky is also significantly less crowded than the optical sky and, as a result, less susceptible to contamination. This makes it easier to identify active galaxies both in the core of the cluster as well as in the outer regions, which have been relatively poorly investigated in previous optical studies of active cluster galaxies. 1 Current address: 5801 South Dorchester Avenue, Apartment 6B, Chicago, IL 60637. 2 Visiting Astronomer, Kitt Peak National Observatory, National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation. 3 The National Radio Astronomy Observatory is a facility of the NSF operated under cooperative agreement by Associated Universities, Inc.
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low-power radio sources for z 0.4 clusters relative to the nearby galaxy population. These sources are presumably linked to an increase in the star formation in clusters at higher redshifts. There does not appear to be much evolution in the radio population at z d 0.3, yet there are exceptional clusters that exhibit very high radio fractions. In this paper, we extend the above analyses using a sample of 22 X-ray–bright clusters to further understand the radio galaxy population of intermediate-redshift clusters and the possible connection between the active cluster members and the overall dynamical state of the cluster. We present the optical and radio data for the cluster sample. The clusters were originally selected from the Abell catalog and include all clusters with richness class 2 and m10 17.6 (i.e., z e 0.2). The X-ray–brightest clusters in the sample were then selected using a ROSAT All-Sky Survey (RASS; Voges 1992) flux limit of 0.08 counts s1 in a 9000 detection cell. These count rates correspond to a luminosity cutoff of LX 4 1043 to 2 1045 ergs s1 for 0.2 < z < 0.34. The X-ray analysis for the first half of this sample was presented in Rizza et al. (1998). The complete sample is listed in Table 1. The paper is organized as follows: In x 2, we discuss the optical observations and analysis including our richness and blue fraction calculations. Section 3 describes our radio observations, source identification, and luminosity class determination. We present our morphological classification
4
Throughout the paper, we adopt H0 = 75 km s1 Mpc1 and q0 = 0.1.
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in x 4 and analyze the combined radio and optical properties of the radio galaxies. Finally, we discuss our results in x 5.
2. OPTICAL IMAGING
The optical imaging was performed using the Kitt Peak National Observatory (KPNO) 0.9 m telescope with the T2K CCD. The observations were acquired over 10 nights divided into two observing runs during 1998 April 30 to May 3 and 1998 November 19–23. The clusters were imaged in the Johnson B- and Cousins R-band filters for a total of 60 minutes each. The observations were divided into three 20 minute exposures and were dithered to facilitate the removal of cosmic rays and bad pixels. During the course of the night, about 10–15 Landolt standard stars were observed in both bands over a range in air mass in order to monitor the atmospheric conditions and serve as absolute flux calibrators. Several of the original nights were not photometric, so bootstrapped calibration frames were obtained using the New Mexico State University 1 m telescope at the Apache Point Observatory on 1998 July 5 and 12 and 1999 July 1, and the University of New Mexico Capilla Peak 2400 telescope on 2000 January 9. The data were reduced in the standard fashion using the software package IRAF. Once the images were reduced, their coordinate solutions were obtained using the task XTRAN within the AIPS software package. Positions of Guide Star Catalog stars in the field were used to calculate the solution, yielding errors in the x- and y-coordinate system that were, on average, less than an arcsecond for all
TABLE 1 Cluster Sample Cluster (1)
R.A. (J2000.0) (2)
Decl. (J2000.0) (3)
z (4)
Counts s1 (5)
N0.5 (6)
fB (7)
NRG (8)
fRG (9)
fLLRG (10)
A0141a .............. A0209 ............... A0360 ............... A0383 ............... A0773b .............. A0963a .............. A1111 ............... A1246 ............... A1550 ............... A1576 ............... A1703 ............... A1704b .............. A1758 ............... A2111b .............. A2219 ............... A2254 ............... A2294 ............... A2390 ............... A2552 ............... A2631 ............... A2645b .............. A2667a ..............
01 05 35.4 01 31 53.2 02 31 24.6 02 48 01.9 09 17 51.8 10 17 01.7 10 50 34.5 11 24 00.1 12 28 57.7 12 37 00.5 13 15 05.8 13 14 26.1 13 32 44.7 15 39 44.1 16 40 22.1 17 17 49.5 17 24 39.6 21 53 35.3 23 11 32.9 23 37 41.5 23 41 16.6 23 51 40.7
24 39 06 13 36 49 +06 59 07 03 32 10 +51 43 29 +39 02 52 02 36 20 +21 29 11 +47 37 58 +63 11 08 +47 37 58 +64 34 40 +50 33 10 +34 24 58 +46 42 34 +19 40 30 +85 53 21 +17 41 10 +03 38 33 +00 16 05 09 01 25 26 05 06
0.230 0.206 0.220 0.187 0.220 0.206 0.165 0.190 0.254 0.302 0.259 0.216 0.279 0.229 0.225 0.178 0.178 0.228 0.229 0.275 0.250 0.230
0.13 0.17 0.10 0.14 0.21 0.18 0.13 0.12 0.08 0.09 0.09 0.13 0.11 0.09 0.33 0.14 0.13 0.32 0.14 0.11 0.10 0.35
... 40 24 28 109 ... 21 43 15 28 56 22 42 61 43 32 34 46 27 24 65 ...
... 0.03 0.003 0.07 0.14 ... 0.05 0.10 0.05 0.06 0.08 0.08 0.12 0.18 0.24 0.17 0.05 0.07 0.25 0.29 0.07 ...
... 14(4) 14(3) 10(3) 8(5) ... 7(2) 6(4) 8(2) 10(2) 5(2) 5 11(2) 4(3) 14(4) 6(2) 6(0) 12(3) 16(0) 17(1) 7(4) ...
... 0.11 0.15 0.11 0.04 ... 0.13 0.08 0.18 0.07 0.05 0.08 0.07 0.02 0.08 0.05 0.11 0.08 0.14 0.20 0.04 ...
... 0.06 0.08 0.07 0.01 ... 0.07 0.03 0.02 0.01 0.01 0.05 0.01 0.01 0.04 0.05 0.06 0.06 0.12 0.13 0.00
Notes.—Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, and arcseconds. Col. (1): cluster ID; col. (2): right ascension; col. (3): declination; col. (4): redshift; col. (5): RASS count rate measured in a 9000 detect cell; col. (6): richness within 0.5 Mpc (see x 2.1); col. (7): blue fraction (see x 2.2); col. (8): number of radio galaxy identifications with the number of confirmed redshifts shown in parenthesis; col. (9): radio galaxy fraction; col. (10): low-luminosity radio galaxy fraction (see x 3.5). a X-ray data only. b Optical and radio analysis from M99, R1 and R2.
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Fig. 1a Fig. 1.—R-band images are shown for several of the clusters: (a) A209, gray-scale flux range = 2525.0–2800.0; (b) A360, gray-scale flux range = 1900.0– 2400.0; (c) A1246, gray-scale flux range = 1900.0–2100.0; (d ) A1550, gray-scale flux range = 1700.0–1900.0. The HLRGs are marked with squares, and the LLRGs with circles. Notice the large numbers of LLRGs for A209 and A360 in comparison with A1246 and A1550.
cases. After the coordinate solution was determined for both filters, the B-band images were shifted to match the Rband coordinate system. Figure 1 shows the R-band images for four clusters in the sample. Aperture photometry of the cluster galaxies was performed using the IRAF task PHOT and the zero-point magnitudes determined from the standard stars. We used a Gunn-Oke (GO) aperture, which is equal to a linear size of 26.2 h1 75 kpc in diameter. At redshifts relevant to this study, the GO aperture encompasses the majority of the galaxy light yet is small enough that there is negligible contamination in the aperture from neighboring galaxies. The magnitudes were measured in both B and R for the entire cluster frame and the BR colors were determined. A calibration error of 0.05 mag was adopted as the total photometric calibration error, including the transformation and systematic errors associated with the individual observation or background errors. The R-band errors are dominated by this external error, and 0.05 is formally adopted for Rmagnitude errors. The color errors, however, are sometimes dominated by the counting statistics in the B images,
especially for the fainter galaxies. We therefore include the errors for the colors for the radio galaxies, containing both the counting and calibration errors (see Table 3). The magnitudes were corrected for absorption using the values obtained from the NASA Extragalactic Database (NED).5 K-corrections for the measured magnitudes were then calculated using the K-correction versus galaxy type results from M99. The values were determined assuming the spectral synthesis models of Bruzual & Charlot (1993) over a range in redshift from z = 0.0–0.5. To maintain consistency with the Butcher-Oemler analysis (Butcher & Oemler 1978a, 1984), a color-magnitude (CM) correction was applied to the galaxy colors. This effect is redshift dependent and is larger at higher z. To apply the correction, a significant number of galaxies is needed. We
5 This research has made use of the NASA/IPAC Extragalactic Database, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.
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Fig. 1b
therefore decided to use the results from Gladders et al. (1998), who studied a large sample of clusters over a range in redshift with very high signal-to-noise ratio. A fit to their results produces the relationship m = 0.1715z 0.0424 (Morrison, Owen, & Ledlow 2003a, hereafter R3). This is consistent with our analysis for several clusters with high signal-to-noise ratios and was therefore adopted for our entire sample. The 0.9 m telescope images cover a linear extent of about 2–2.5 Mpc (at the average redshift of our clusters), which contains the vast majority of the cluster members yet does not provide for regions external to the cluster for use in background counts. Instead, the galaxy surface density of a field sky region centered on the Hubble Deep Field (HDF) was used to determine a background correction for all of our clusters. This field was observed in B and R using the 0.9 m telescope in the same manner as the object frames and was subjected to an identical reduction process (R3). The galaxy positions were again selected using AIPS as described above and aperture photometry was performed on all identified objects. This produced 3000 objects over the full field of the 0.9 m telescope frames (230 230 ). The photometry on these galaxies was then performed separately for each cluster in the following manner. First, the
HDF was ‘‘ placed ’’ at the redshift of the cluster to determine the angular size of a GO aperture. The magnitudes of the galaxies in the field were then measured with this aperture size. The aperture was adjusted for each cluster redshift, yielding unique magnitudes for all cases, and the final number of galaxies above a limiting magnitude was determined for each cluster individually. R3 determined that this method produced background counts with an error of about 23%. Although unavoidable under the circumstances, this error is rather high and must be kept in mind in interpreting some of the optical data. 2.1. Richness Counts The subsequent analyses followed the prescription outlined in R3. First, all galaxies within 0.5 Mpc of the cluster center (as defined by the peak in the X-ray emission) that had an absolute magnitude MR 20.5 (1.5 mag fainter than MR , the break in the R-band luminosity function; Ledlow & Owen 1995) were counted. All galaxies above this same detection limit in the background field (assuming it was placed at the redshift of the cluster) were then counted. The number of background counts was normalized by the ratio of the two areas, and this number was subtracted from the total cluster counts, yielding a value for N0.5. This richness is similar to the
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Fig. 1c
Bahcall richness count and has been shown to be a better representation of the cluster gravitational potential than the Abell richnesses (Bahcall 1981). The richness values for each cluster are listed in Table 1. 2.2. Blue Fraction Measurements Our method is a variant of Butcher & Oemler’s (1978a, 1978b, 1984) definition of a blue galaxy and calculation of the cluster blue fractions. Under the BO description, a galaxy is considered ‘‘ blue ’’ if its BV color in the rest frame of the cluster is greater than 0.2 mag bluer than the peak color of the old galaxy population. Only galaxies with MV 21 and located within the inner 30% of the cluster radial distribution were included in the count. In the method adopted here, SBANDS in IRAF was used to create a synthetic galaxy spectra Save, which has a color 0.2 mag bluer than a typical S0/E galaxy. This was done for a number of different redshifts, reflecting the evolving color dependence of z. Thus, for a given z, the BR value equal to (BV 0.2) is simply Save(z). The next step was to define the color of the ‘‘ old ’’ population at the cluster redshift. Initially, an estimate of the peak was determined by fitting the overall color distribution of the galaxies located within the inner 0.5 Mpc with MR 20.5 after correcting for absorp-
tion and the color-magnitude (CM) effect. This effect is due to the fact that more massive galaxies tend to be able to hold on to their supernovae-driven (and therefore enriched) gas more effectively than less massive galaxies. This results in massive (brighter) galaxies being more metal-rich and thereby redder in color. If this effect is not taken into account, the color distributions of the cluster galaxies appear broadened and somewhat skewed. The CM effect is a function of redshift, increasing in magnitude with larger values of z (Gladders et al. 1998). Using the peak as a guide, a range of colors was averaged to find the mode or peak of the cluster color distribution (ModeCl) ½ModeCl 4ðBRÞ ðBRÞ ½ModeCl þ 4ðBRÞ : ð1Þ After background correction the ‘‘ average ’’ color of the E/S0 population could be described as follows: hBRiE=S0 ¼
½ModeCl 4ðBRÞ ½ModeFLD 4ðBRÞ ; Ncluster Nfield
ð2Þ
where ModeFLD is the mode of the field distribution and
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Fig. 1d
Ncluster and Nfield are the number of galaxies in the cluster and field, respectively, within this color range (R1). From this average value, a galaxy was deemed ‘‘ blue ’’ by simply subtracting its color from the average and comparing it with the 4(BR) required at its redshift, as follows: hBRiE=S0 ðBRÞgalaxy 4ðBRÞz :
ð3Þ
A field correction was made by determining the number of ‘‘ blue ’’ galaxies in the background field using the same method and subtracting them from the cluster population. In so doing, it is possible for a cluster to have a negative blue fraction if the cluster population is very red. The blue fraction results are listed in Table 1, along with the richness results described above. This method was designed to mimic the BO analyses (Butcher & Oemler 1978b, 1984) in principle as well as in practice, as noted in R3. Only one of the clusters in this sample has a BO determined value, but for that case there is agreement within the error bars.
done using the MX multifiber spectrograph on the Steward Observatory 9000 Bok telescope located at the Kitt Peak National Observatory. The combined system has a 450 field ˚ at a of view with spectral coverage from 3500 to 7200 A ˚ . A detailed description of the MX resolution of 2.75 A Spectrometer is given in Hill & Lesser (1986, 1988). The observations were performed during several runs in 1997 October, 1998 February, and 1999 July. The data were reduced using software developed for the IRAF environment by J. Hill, with some modifications by other users, and the velocity analysis was performed using the IRAF task FXCOR in the RV package. This is an implementation of the Tonry & Davis (1979) technique, which is based on the correlation in the Fourier transform of the continuum absorption features. Although weather prevented us from obtaining the complete coverage we had intended, the new redshifts of the radio galaxies determined with MX are listed in Table 3. 3. RADIO ANALYSIS
3.1. Observations and Reduction Techniques 2.3. Optical Spectroscopy We conducted a program to obtain optical spectroscopy of each radio galaxy to confirm cluster membership. This was
The radio observations were performed with the Very Large Array (VLA) radio telescope at 20 cm using the A configuration. This setup was chosen because of its high
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TABLE 2 Radio Flux Limits Cluster (1)
R.A. (J2000.0) (2)
Decl. (J2000.0) (3)
z (4)
(lJ) (5)
log (L1.4) (6)
A0209 ............... A0360 ............... A0383 ............... A1111 ............... A1246 ............... A1550 ............... A1576 ............... A1703 ............... A1758 ............... A2219 ............... A2254 ............... A2294 ............... A2390 ............... A2507 ............... A2552 ............... A2631 ...............
01 31 53.2 02 31 24.6 02 48 01.9 10 50 34.5 11 24 00.1 12 28 57.7 12 37 00.5 13 15 05.8 13 32 44.7 16 40 22.1 17 17 49.5 17 24 39.6 21 53 35.3 22 56 35.9 23 11 32.9 23 37 41.5
13 36 49 +06 59 07 03 32 10 02 36 20 +21 29 11 +47 37 58 +63 11 08 +47 37 58 +50 33 10 +46 42 34 +19 40 30 +85 53 21 +17 41 10 +05 25 23 +03 38 33 +00 16 05
0.206 0.220 0.187 0.165 0.190 0.254 0.302 0.259 0.279 0.225 0.178 0.178 0.228 0.196 0.229 0.275
26.5 34.2 30.1 49.9 49.2 45.3 47.9 44.5 46.2 42.6 54.9 54.6 37.1 24.9 28.7 27.9
22.00 22.18 21.96 22.08 22.20 22.45 22.63 22.45 22.54 22.30 22.18 22.18 22.25 21.94 22.14 22.30
Note.—Col. (1): cluster; col. (2): right ascension; col. (3): declination; col. (4): redshift; col. (5): rms noise of radio map; col. (6): 4 radio luminosity limit.
spatial resolution (1>5) and broad field of view (2–3 Mpc in radius at the redshifts of the sample). The observations, obtained in two separate runs, were performed in spectral line mode in order to reduce the bandwidth smearing. The primary run for this project was on 1996 November 16 and 18. Each cluster was observed for 2–4 hr with 20–30 minute exposures. Secondary flux calibrators were observed every hour in order to account for phase errors and changes in the atmospheric and system conditions, and the absolute fluxes of the radio sources were calibrated by observing the primary flux calibrator 3C 48 at the beginning and end of the runs. The data were processed using standard procedures within the AIPS environment, where the maps were created using the task IMAGR after cleaning and selfcalibrating the data. The field itself was divided into four overlapping subfields (except for the case of A2390 for which a nine-subfield configuration was used), with their common corner centered on the core of the cluster. The UV data were divided into 10 minute time segments and processed individually in order to most accurately account for the three-dimensional effects. The final maps from each time segment were corrected for their threedimensional effects using the task OHGEO and then combined, weighting by 1/map, to produce the four composite maps for each cluster. The average primary beam of the maps was fitted, and then all of the fields were remapped using the same bmaj and bmin in IMAGR. The newest version of IMAGR is able to correct for threedimensional effects and thus makes the division of the UV data into time segments unnecessary. Unfortunately, this was only available during the reduction of the final cluster (A2390) for which it was used. A number of the clusters (A1111, A1246, A1550, A1576, A1703, A1758, A2219, A2254, A2294) were observed earlier as a trial run with shorter exposure times and are therefore not quite as deep as the new observations. Thus, the formal flux limit of the sample is 2 1022 W Hz1 with a subset that is complete to 1 1022 W Hz1. The limiting flux densities for the maps and their corresponding 4 limiting
luminosities are listed in Table 2 for all of the clusters, including the observations done in the trial run. 3.2. Source Identification The individual source identifications were determined using the task SAD within AIPS. All detections with peak flux densities above 4 for the given subfield were marked as potential sources. A Gaussian model was fitted to each source using the task IMFIT and the peak flux density, integrated flux density, and major and minor axes were recorded. For some very extended or complex structures, the fluxes were measured using TVSTAT instead. The fluxes and their errors were corrected for primary-beam diminution using the task PBCOR, and radio luminosities for the final identifications were calculated (see Table 3). The radio positions were correlated with the optical galaxy positions in order to create a catalog of optical/radio identifications. All unresolved sources (h < 300 ) required the optical position to be within 100 , sources with angular size h < 600 within 300 , and sources with h > 600 within 500 . The varying scale is based on the degrading positional accuracy of the radio data as a function of increasing angular size. The radio identifications were also required to have a projected linear distance d 2.5 Mpc from the cluster center, as defined by the X-ray centroid. A final cut was placed on the radio identifications to try to ensure that projected field galaxies were excluded. If a redshift for the galaxy was known, then the identification was included if it had |zgal zcluster| = 4z 0.01. If the redshift was not known, the identification was subjected to an optical magnitude cut. The results of R3 showed that few radio galaxies (aside from centrally dominant galaxies) had optical magnitudes MR 23.0. Therefore, excepting central galaxies, all identifications brighter than this limit were assumed to be foreground sources and were rejected. A lower limit of MR 21.0 was also set to avoid background contamination. This value is 1 mag fainter than MR (the R-band magnitude where the break in the luminosity function occurs) and places a conservative constraint on the
TABLE 3 Radio Identifications
ID (1)
R.A. (J2000.0) (2)
Decl. (J2000.0) (3)
z (4)
d (Mpc) (5)
MR (6)
BR (7)
log (L1.4) (W Hz1) (8)
h (arcsec) (9)
0.59a 0.05b 1.14 0.32 1.43 0.92 1.96 0.93 0.88 1.66 0.38 2.04 0.79 2.23
22.52 22.43 21.91 21.41 22.51 22.36 21.25 21.22 22.60 22.27 21.57 21.03 22.61 21.10
2.25(0.05) 2.20(0.05) 2.16(0.05) 2.14(0.06) 2.26(0.05) 2.26(0.05) 1.59(0.05) 2.12(0.06) 1.50(0.05) 2.08(0.05) 1.48(0.05) 2.37(0.06) 2.22(0.05) 2.35
23.88 24.25 23.51 23.25 23.25 22.53 22.82 22.28 22.54 23.45 22.56 22.34 22.68 22.65
Ext Ext Ext 7.88 2.93 4.90 4.00 0.00 2.13 0.98 1.05 0.00 5.00 2.12
0.14 0.21 0.46 1.97 1.64 0.16 0.92 0.52 0.80 0.29 0.74 1.99 2.17 2.39
23.30 22.41 21.42 21.85 22.97 21.37 22.36 22.41 22.73 21.63 21.67 21.00 22.56 21.90
2.42(0.05) 2.22(0.05) 1.71(0.05) 1.92(0.05) 1.56(0.05) 1.23(0.05) 2.28(0.05) 2.20(0.05) 2.34(0.05) 2.42(0.06) 1.48(0.05) 1.14(0.05) 2.30 2.30
24.02 24.07 22.43 22.76 22.41 22.71 22.47 23.24 22.51 23.37 22.91 24.72 22.71 25.48
3.51 Ext 1.88 0.00 0.00 2.09 2.03 0.84 1.42 9.10 3.84 1.22 2.40 3.50
0.09 0.63 0.95 0.96 1.63 0.38 0.92 1.72 2.31 2.21
23.15 22.24 21.94 21.70 21.37 21.88 22.26 21.17 22.47 22.20
2.15(0.05) 1.64(0.05) 2.20(0.06) 2.46(0.06) 1.75(0.06) 1.31(0.05) 2.06(0.05) 1.90(0.06) 2.18 2.18
24.49 22.77 22.22 22.60 22.30 22.55 23.87 22.42 23.36 22.82
0.99 2.96 0.00 2.00 0.00 3.72 0.96 2.33 1.64 1.67
1.18 1.58 1.37 0.08 1.28 2.07 2.07
22.03 22.92 21.25 22.66 21.35 21.21 22.43
2.33(0.06) 1.36(0.05) 1.30(0.05) 2.19(0.05) 1.90(0.06) 2.17 2.17
22.51 23.04 22.42 23.94 22.66 23.14 22.59
0.00 2.18 0.00 1.35 3.22 1.88 0.00
0.60 1.17 0.39 0.15 0.74 0.55
22.24 22.22 22.04 21.66 21.90 21.47
2.14(0.05) 2.20(0.05) 2.11(0.05) 2.24(0.06) 2.17(0.06) 2.07(0.06)
22.78 23.37 23.34 24.08 22.37 23.46
0.00 1.38 Ext Ext 0.00 Ext
A209 1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12............... 13............... 14...............
01 31 50.80 01 31 52.45 01 31 34.87 01 31 55.60 01 31 57.38 01 32 13.48 01 31 11.34 01 31 33.07 01 31 33.56 01 31 36.76 01 31 47.12 01 31 38.96 01 31 44.58 01 31 37.59
13 33 35.89 13 37 00.21 13 32 19.94 13 38 30.04 13 44 42.90 13 38 09.32 13 40 49.83 13 38 28.79 13 38 01.82 13 45 07.63 13 38 21.69 13 47 38.53 13 40 42.31 13 24 59.44
0.206 0.213 ... 0.214a ... ... ... ... ... 0.199a ... ... ... ... A360
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12............... 13............... 14...............
02 31 27.60 02 31 23.56 02 31 15.58 02 30 42.47 02 31 38.38 02 31 22.85 02 31 32.48 02 31 35.24 02 31 41.65 02 31 30.69 02 31 08.92 02 30 57.07 02 31 34.43 02 30 38.89
06 58 55.93 07 00 12.62 06 58 10.09 06 59 19.29 06 51 07.06 06 58 24.88 06 54 40.34 06 58 18.17 06 59 21.71 06 58 44.72 06 58 45.83 06 51 04.34 06 47 51.56 06 53 29.15
0.220c ... ... ... ... ... ... ... 0.217a 0.215a ... ... ... ... A383
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10...............
02 48 03.40 02 48 12.23 02 48 24.69 02 47 57.94 02 48 04.33 02 48 04.97 02 48 21.90 02 48 25.20 02 48 22.40 02 48 52.81
03 31 45.71 03 29 26.80 03 31 46.26 03 26 32.12 03 22 26.50 03 34 17.14 03 34 26.39 03 40 39.82 03 19 17.57 03 35 52.56
0.189a 0.196a ... ... ... ... 0.186a ... ... ... A1111
1................. 2................. 3................. 4................. 5................. 6................. 7.................
10 50 55.66 10 51 03.71 10 51 06.32 10 50 36.53 10 50 05.78 10 49 39.84 10 50 10.74
02 30 35.47 02 28 53.78 02 32 05.66 02 36 17.10 02 40 43.53 02 35 54.59 02 48 36.32
0.166a ... ... 0.165a ... ... ... A1246
1................. 2................. 3................. 4................. 5................. 6.................
11 23 48.02 11 24 22.59 11 23 46.64 11 23 56.19 11 24 12.71 11 24 04.12
21 32 57.43 21 31 32.11 21 28 10.35 21 28 57.61 21 29 35.36 21 27 28.06
... 0.189a 0.189a ... 0.190a 0.187a
126
TABLE 3—Continued
ID (1)
R.A. (J2000.0) (2)
Decl. (J2000.0) (3)
MR (6)
BR (7)
log (L1.4) (W Hz1) (8)
h (arcsec) (9)
1.76 0.97 0.11 1.36 0.95 0.76 0.24 0.26
21.34 21.20 22.35 22.16 21.42 21.01 22.00 22.74
2.86(0.08) 2.96(0.10) 2.67(0.06) 3.40(0.07) 2.28(0.06) 2.48(0.08) 2.56(0.06) 2.48(0.05)
23.25 22.58 23.52 23.10 23.08 23.82 23.58 24.03
2.09 0.00 Ext 0.75 1.31 0.47 1.15 1.54
0.83 0.06 2.28 0.09 1.24 0.86 0.52 0.74 1.67 0.75
22.11 23.15 22.94 22.46 21.92 22.13 22.80 21.71 22.61 23.21
2.60(0.06) 2.65(0.05) 2.67(0.05) 2.76(0.06) 1.71(0.05) 2.02(0.06) 2.08(0.05) 1.73(0.06) 2.31(0.05) 2.67(0.05)
24.19 23.62 23.52 24.58 23.07 23.29 22.90 22.77 23.73 24.25
3.39 8.15 7.32 2.52 0.00 3.71 1.85 1.22 1.34 2.57
1.15 0.46 0.29 0.39 1.21 0.09
22.54 21.96 21.85 21.81 22.39 22.29
2.41(0.05) 2.51(0.06) 2.40(0.06) 2.45(0.06) 1.97(0.05) 2.55(0.06)
23.41 23.51 24.11 22.84 23.04 24.80
4.33 0.00 Ext 0.00 0.00 Ext
0.31 0.46 0.57 1.35 1.70 0.93 1.43 0.34 0.38 1.28 1.86
23.04 21.93 21.77 22.13 21.77 22.82 21.35 22.40 22.04 22.46 22.44
2.64(0.05) 2.49(0.06) 2.22(0.06) 2.35(0.06) 2.20(0.06) 2.53(0.05) 1.79(0.06) 2.42(0.06) 2.67(0.07) 2.28(0.05) 2.45(0.06)
24.15 24.18 23.01 22.71 23.01 23.52 23.01 25.31 23.83 22.83 24.08
3.09 5.76 0.00 0.00 0.00 0.00 1.65 Ext Ext 0.00 3.28
1.80 0.77 1.16 0.27 0.15 0.16 1.42 1.65 1.80 1.85 0.42 0.19 1.54 0.11
22.18 21.76 21.46 21.67 21.42 21.95 22.61 22.58 21.65 22.08 22.18 22.57 22.08 21.64
2.41(0.05) 2.69(0.06) 2.33(0.06) 2.34(0.05) 2.16(0.06) 1.70(0.05) 2.42(0.05) 2.42(0.05) 2.43(0.06) 2.35(0.05) 2.06(0.05) 2.47(0.05) 2.12(0.05) 2.41(0.06)
22.61 22.62 23.47 23.87 23.08 22.58 22.71 22.54 22.97 24.15 24.37 24.05 22.54 23.71
0.00 0.00 0.00 8.93 6.04 0.00 2.23 1.43 0.00 Ext Ext 2.30 1.32 Ext
d (Mpc) (5)
z (4)
A1550 1................. 2................. 3................. 4................. 5................. 6................. 7................. 8.................
12 28 35.99 12 28 49.60 12 28 58.36 12 28 19.84 12 28 31.38 12 28 46.43 12 29 00.83 12 29 02.46
47 45 28.56 47 42 22.31 47 38 29.15 47 37 48.46 47 37 57.25 47 34 58.35 47 36 53.56 47 36 56.14
... ... ... 0.256a ... ... ... 0.263a A1576
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10...............
12 36 53.46 12 36 58.49 12 38 26.12 12 36 57.42 12 36 40.65 12 36 45.94 12 36 53.11 12 37 22.89 12 38 02.22 12 37 28.35
63 14 35.38 63 11 13.31 63 12 17.26 63 11 16.51 63 15 54.35 63 14 24.41 63 09 04.99 63 09 12.67 63 09 44.06 63 10 27.63
... ... ... 0.302d ... ... ... ... ... 0.293a A1703
1................. 2................. 3................. 4................. 5................. 6.................
13 15 06.39 13 15 16.30 13 15 02.35 13 15 15.32 13 14 36.36 13 15 07.84
51 54 28.07 51 50 23.83 51 50 20.51 51 48 04.47 51 45 26.15 51 48 58.18
0.274a ... ... 0.275a ... ... A1758
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11...............
13 32 38.38 13 32 39.40 13 32 33.90 13 32 41.03 13 32 41.07 13 33 02.06 13 33 23.09 13 32 52.09 13 32 53.01 13 33 12.96 13 32 01.39
50 33 36.08 50 34 32.50 50 30 50.02 50 26 37.10 50 25 03.23 50 29 28.24 50 30 28.53 50 31 48.00 50 31 37.07 50 29 01.43 50 28 01.53
0.279e ... ... ... ... ... ... ... ... ... 0.278a A2219
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12............... 13............... 14...............
16 41 13.65 16 40 37.90 16 40 06.25 16 40 11.58 16 40 21.11 16 40 22.99 16 40 56.80 16 41 07.68 16 41 11.54 16 40 44.56 16 40 31.67 16 40 23.77 16 39 32.62 16 40 15.93
46 42 04.19 46 45 08.30 46 48 23.16 46 43 17.97 46 43 24.47 46 42 17.95 46 39 10.22 46 40 37.45 46 40 03.46 46 34 07.17 46 42 22.25 46 42 10.41 46 41 44.98 46 42 31.68
... ... ... ... 0.219a ... 0.227a ... ... ... ... 0.229a ... 0.228f
127
TABLE 3—Continued
ID (1)
R.A. (J2000.0) (2)
Decl. (J2000.0) (3)
z (4)
MR (6)
BR (7)
log (L1.4) (W Hz1) (8)
h (arcsec) (9)
2.17 0.06 0.92 2.04 1.50 0.42
22.34 22.24 22.29 21.84 22.38 21.48
1.96(0.05) 1.59(0.05) 2.07(0.05) 1.50(0.05) 1.99(0.05) 1.49(0.05)
22.37 22.46 22.64 22.90 23.00 22.52
0.00 0.00 2.61 2.05 2.79 0.00
0.83 0.06 0.90 1.75 1.18 0.56
21.01 21.36 21.33 22.10 22.10 22.13
1.91(0.06) 2.41(0.07) 1.86(0.06) 1.95(0.05) 1.97(0.05) 2.42(0.05)
22.48 23.28 22.25 23.48 22.77 24.30
0.00 0.00 0.00 1.10 1.89 Ext
1.44 0.98 0.13 0.49 0.66 0.13 0.19 1.73 2.10 1.99 2.00 1.06
21.36 20.77 22.84 21.99 21.65 21.61 21.57 21.89 22.54 22.79 21.67 21.45
2.42(0.09) 1.93(0.09) 1.70(0.05) 1.71(0.06) 2.15(0.07) 2.43(0.08) 1.78(0.06) 1.80(0.06) 2.26(0.06) 2.37(0.06) 1.56(0.06) 2.10(0.08)
22.73 22.65 25.39 23.03 22.95 23.61 22.91 22.75 22.97 22.97 23.58 22.97
1.15 2.89 0.34 7.01 6.85 3.68 0.00 0.00 2.94 0.00 1.16 3.94
0.08 0.14 1.47 0.61 0.69 0.78 1.26 1.92 2.20 2.21 1.71 1.30 2.21 2.14 2.01 2.45
21.08 21.79 21.69 21.11 21.14 21.91 21.80 21.57 21.96 21.87 21.40 21.58 21.01 21.92 21.90 21.01
2.21(0.07) 2.26(0.06) 2.48(0.06) 2.17(0.07) 2.53(0.08) 2.53(0.06) 2.32(0.06) 2.17(0.06) 0.90(0.05) 2.04(0.05) 1.29(0.05) 1.55(0.05) 1.98(0.07) 2.10(0.05) 2.28(0.06) 2.36(0.10)
22.52 22.32 22.90 22.74 24.72 24.33 22.57 22.74 22.70 22.60 22.50 22.38 22.91 22.73 23.10 22.84
0.00 0.00 0.00 3.01 Ext Ext 0.00 0.00 2.27 1.97 0.00 0.00 1.24 1.21 1.76 1.98
d (Mpc) (5)
A2254 1................. 2................. 3................. 4................. 5................. 6.................
17 18 20.47 17 17 45.53 17 17 30.61 17 18 36.33 17 17 24.27 17 17 35.91
19 50 57.16 19 40 39.22 19 35 48.52 19 37 41.89 19 32 28.34 19 38 47.17
... ... 0.179a ... 0.173a ... A2294
1................. 2................. 3................. 4................. 5................. 6.................
17 23 59.63 17 24 13.40 17 28 01.62 17 32 22.56 17 17 56.84 17 22 48.31
85 58 22.77 85 53 33.36 85 57 02.93 85 59 48.20 85 56 15.04 85 56 21.51
... ... ... ... ... ... A2390
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12...............
21 54 03.00 21 53 18.62 21 53 36.82 21 53 45.59 21 53 49.68 21 53 35.19 21 53 31.42 21 52 59.83 21 54 07.63 21 53 29.21 21 53 36.91 21 53 46.99
17 37 41.66 17 44 18.19 17 41 43.69 17 41 47.82 17 41 20.50 17 41 50.48 17 41 33.56 17 38 14.71 17 33 29.55 17 30 59.51 17 30 49.00 17 36 28.15
... 0.234g 0.230a 0.215g ... ... ... ... ... ... ... ... A2552
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12............... 13............... 14............... 15............... 16...............
23 11 33.13 23 11 35.60 23 11 41.09 23 11 46.28 23 11 48.66 23 11 48.34 23 11 53.11 23 10 54.81 23 10 55.93 23 11 02.16 23 11 07.67 23 11 42.20 23 11 51.08 23 10 52.80 23 10 54.12 23 11 37.71
03 38 53.07 03 39 18.80 03 46 04.44 03 39 42.76 03 38 55.17 03 40 45.12 03 43 11.39 03 39 01.03 03 44 40.47 03 46 41.22 03 44 21.66 03 32 15.99 03 28 03.02 03 35 02.50 03 36 19.41 03 26 02.16
... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
z 0.2 RICH ABELL CLUSTERS
129
TABLE 3—Continued
ID (1)
R.A. (J2000.0) (2)
Decl. (J2000.0) (3)
z (4)
d (Mpc) (5)
MR (6)
BR (7)
log (L1.4) (W Hz1) (8)
h (arcsec) (9)
22.15 21.31 21.62 22.99 22.43 21.54 22.25 22.62 21.71 21.57 22.29 22.45 22.07 22.16 21.53 22.37 22.23
2.28(0.06) 2.58(0.10) 1.65(0.06) 2.37(0.05) 2.09(0.06) 1.74(0.06) 2.26(0.06) 2.42(0.06) 1.42(0.05) 2.22(0.07) 1.82(0.05) 2.28(0.06) 2.29(0.06) 1.70(0.05) 2.65(0.09) 0.82(0.05) 2.26(0.06)
23.42 22.70 22.51 22.90 25.24 22.96 22.86 22.94 22.93 22.96 23.58 24.14 22.93 22.72 23.00 22.94 23.19
0.95 0.00 0.00 0.00 Ext 2.23 4.92 0.00 2.40 0.00 3.1 Ext 0.00 0.00 1.16 3.5 0.89
A2631 1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12............... 13............... 14............... 15............... 16............... 17...............
23 37 41.34 23 37 54.39 23 37 56.58 23 37 39.73 23 37 40.12 23 37 07.14 23 37 15.08 23 37 23.62 23 37 32.84 23 37 33.39 23 37 35.66 23 37 48.00 23 37 51.97 23 38 04.26 23 38 10.17 23 37 12.70 23 37 19.05
00 18 33.80 00 22 20.89 00 17 18.36 00 16 15.98 00 16 40.00 00 18 28.62 00 16 19.84 00 16 20.81 00 16 33.79 00 22 37.14 00 25 22.27 00 14 25.50 00 14 21.88 00 08 26.81 00 13 40.51 00 16 38.27 00 13 14.26
... ... ... 0.278h ... ... ... ... ... ... ... ... ... ... ... ... ...
0.55 1.57 0.90 0.08 0.14 1.94 1.43 0.96 0.46 1.50 2.07 0.53 0.71 2.12 1.69 1.56 1.36
Note.—Col. (1): galaxy ID; col. (2): right ascension; col. (3): declination; col. (4): redshift; col. (5): linear projected distance from the cluster center; col. (6): absolute R-band magnitude; col. (7): galaxy color; col. (8): radio luminosity; col. (9): angular size of radio source. a Redshifts obtained with MX. b Kristian, Sandage, & Westphal 1978. c Quintana & Ramirez 1995. d Huchra et al. 1990. e Allen et al. 1992. f Owen, White, & Burns 1992. g Yee et al. 1996. h Crawford et al. 1995.
identifications, thus supporting inclusion of only true cluster members. The linear field of view for several of the 0.9 m images is less than 2.5 Mpc. This required outlying identifications and magnitudes to be determined from the Digital Sky Survey6 for the radio sources that were unobserved in our R-band images. The optical magnitudes for these cases are listed in Table 3, along with their colors (without errors). The sky survey magnitudes were determined using the zero point calculated by Bliton (2000) based on comparisons from published and 0.9 m magnitudes for a large sample of galaxies over the entire sky. The radio parameters of the final cluster identifications are listed in Table 3 along with the optical data. The radio sources are characterized by their absolute luminosity and angular resolution to maintain consistency with the selection criteria for cluster membership. To summarize, a radio identification was considered a cluster member and included in the ensuing analysis if 1. dID 2.5 Mpc from cluster center; 2. L1.4GHz 4 ; 3. 4z < 0.01, if known; otherwise,
6 The Digitized Sky Surveys were produced at the Space Telescope Science Institute under grant NAG W-2166. The images of these surveys are based on photographic data obtained using the Oschin Schmidt Telescope on Palomar Mountain and the UK Schmidt Telescope. The plates were processed into the present compressed digital form with the permission of these institutions.
4. MR 21.0; 5. MR 23.0 or is a centrally dominant galaxy. 3.3. Effects of Contamination Ideally, one would like to have redshift confirmation of cluster membership for every proposed radio identification. Unfortunately, this was not possible for the present study, so attempts were made to estimate the probability that field galaxies were erroneously included as cluster members despite the magnitude limits. The contamination estimate was based on the assumption that the background radio galaxy distribution scales the same way as the galaxy population. For each cluster the fraction of background or noncluster objects in each field was calculated by counting all of the cluster galaxies within a 2.5 Mpc radius divided by the number of background galaxies determined from our HDF image. It was then assumed that the noncluster radio galaxies are represented in this same ratio. From the fraction of contamination, the number of radio galaxies that were noncluster members was determined for each cluster. This value was summed for the entire sample and divided by the total area on the sky surveyed to generate an average background radio galaxy surface density. Using this as a measure of the confusing sources, the radio galaxy number and radio fraction were corrected for each cluster according to its area on the sky. The final errors in the counts are listed in column (6) of Table 4. The possible contamination by confusing sources as a function of radial distance from the cluster core was
130
RIZZA ET AL.
Vol. 126
TABLE 4 Radio Contamination Estimates Cluster (1)
R.A. (J2000.0) (2)
Decl. (J2000.0) (3)
z (4)
RG (5)
RGerr (6)
A0209 ............... A0360 ............... A0383 ............... A1111 ............... A1246 ............... A1550 ............... A1576 ............... A1703 ............... A1758 ............... A2219 ............... A2254 ............... A2294 ............... A2390 ............... A2552 ............... A2631 ...............
01 31 53.2 02 31 24.6 02 48 01.9 10 50 34.5 11 24 00.1 12 28 57.7 12 37 00.5 13 15 05.8 13 32 44.7 16 40 22.1 17 17 49.5 17 24 39.6 21 53 35.3 23 11 32.9 23 37 41.5
13 36 49 +06 59 07 03 32 10 02 36 20 +21 29 11 +47 37 58 +63 11 08 +47 37 58 +50 33 10 +46 42 34 +19 40 30 +85 53 21 +17 41 10 +03 38 33 +00 16 05
0.206 0.220 0.187 0.165 0.190 0.254 0.302 0.259 0.279 0.225 0.178 0.178 0.228 0.229 0.275
14(4) 14(3) 10(3) 7(2) 6(4) 8(2) 10(2) 5(2) 11(1) 14(4) 6(2) 6(0) 12(3) 16(0) 17(1)
3 2 3 4 3 2 1 2 2 2 3 3 2 2 2
Note.—Col. (1): cluster; col. (2): right ascension; col. (3): declination; col. (4): redshift; col. (5): number of radio galaxies (number with confirmed redshifts); col. (6): estimated error in radio galaxy count.
examined by scaling the background radio galaxy surface density to the area in each radial bin. The galaxy surface density in rich clusters is highest near the core and falls off in surface density approximately as r2. Because of this, the sources identified near the core are most likely cluster members, whereas the identifications in the outer regions of the search radius may suffer slightly more from contamination. Figure 2 illustrates the results of this analysis.
3.4. Luminosity Classes The radio luminosities were computed for each galaxy identification using the expression L ¼ 4D2L S ;
ð4Þ
where DL is the luminosity distance of the source and S is the integrated flux density. The log luminosities are listed in Table 3, along with the angular size of the source and its projected linear distance d in megaparsecs from the cluster core. When redshift information was not available, the radio identification was assumed to be at the adopted redshift of the cluster. High-powered radio galaxies are believed to be driven by a nuclear process that are associated with AGNs. The contribution of these types of sources to the radio luminosity function drops off with decreasing radio power and is dominated by the contribution of star forming galaxies below 1023 W Hz1. To distinguish between the two populations, a transition radio power was defined, where sources with L1.4 > 1023 W Hz1 are classified as high-luminosity radio galaxies (HLRGs) and sources with L1.4 1023 W Hz1 as low-luminosity radio galaxies (LLRGs). The association follows that the LLRGs are assumed to be primarily due to star formation and the HLRGs are attributed to AGN activity. Although the starforming sources increase in dominance below this division, to fully ensure exclusion of AGN a more conservative cutoff was adopted, whereby all sources with L1.4 1022.75 W Hz1 were classified as starburst radio galaxies (SBRGs). This region of the radio luminosity function is almost
exclusively dominated by star-forming systems and should have minimal contamination by AGNs. The spatial distributions of the radio galaxies were compared as a function of their radio luminosity classes. The surface density of the high- and low-luminosity sources is plotted in Figure 3, which shows the trend that the HLRGs are centrally concentrated, whereas the LLRGs have a much flatter distribution. Using the K-S test as a diagnostic, the two samples are different at greater than a 99.9% level. This trend was first noted in M99 and is supported by the extended sample presented here. The HLRG distribution is consistent with the Ledlow & Owen (1995) results for a large sample of radio galaxies residing in nearby rich clusters and Stocke et al. (1999) for more distant clusters. Those sources surveyed displayed a King (1962) distribution that fell off quickly with increasing radius. Fitting our profiles with King models results in distinctly different core radii for the two distributions. The core radius of the HLRGs is rc = 0.16 Mpc, while that of the LLRG distribution is rc = 0.62 Mpc. As a further comparison of the two distributions, their cumulative distribution functions were examined. By doing this, the range of radii at which the samples differed most significantly was determined. This region occurs at 0.65 Mpc, which is beyond the second bin in Figure 3. This confirms that the difference between the populations resides outside the core of the cluster and is not simply due to the strong central peak of the high-luminosity sources. 3.5. Radio Galaxy Fractions The radio galaxy fractions for all of the clusters were calculated in order to remove the richness of the cluster as a contributing factor in the detection of a radio-bright galaxy. This fraction is a measure of the number of radio galaxies as a function of the number of galaxies observed above a given optical magnitude: fRG ¼
NðLmin L1:4GHz Þ ; N2:5
ð5Þ
z 0.2 RICH ABELL CLUSTERS
No. 1, 2003
A360
A209
131
A383
4
5
4
3
4
3
3
A1111 3
2
2
2
1
2 1
1 1
-1 0.0
0 0
0
0
0.5
1.0
1.5
2.0
2.5
-1 0.0
-1
-1 0.5
1.0
1.5
2.0
2.5
0.0
0.5
2
1.5
2.0
2.5
0.0
0.5
A1576
A1550
A1246 3
1.0
4
5
3
4
1.0
1.5
2.0
2.5
2.0
2.5
2.0
2.5
A1703 3
2 3
2 1
1
2 1 1
0
0
0
0
-1 0.0
0.5
1.0
1.5
2.0
2.5
0.0
0.5
A1758
1.0
1.5
2.0
2.5
0.0
1.0
1.5
2.0
2.5
6
2.0
5
5
1.5
4
4
1.0
3
3
2
2
1
1
0
0 1.0
1.5
2.0
2.5
-1 0.0
1.5
3
0.5
1
0.0 0 -0.5
0.5
1.0
1.5
2.0
2.5
0.0
-1 0.5
1.0
1.5
2.0
2.5
2.0
2.5
0.0
0.5
1.0
1.5
A2631
A2552
3
1.0
2
A2390 4
0.5
A2294
-1.0 0.5
0.0
A2254
6
0.0
0.5
A2219
5
5
4
4
3
3
2
2
1
1
0
0
2 1 0 -1 0.0
0.5
1.0
1.5
2.0
2.5
-1 0.0
0.5
1.0
1.5
2.0
2.5
0.0
0.5
1.0
1.5
Fig. 2.—Contamination estimates for radio galaxy distribution. Histogram of radio galaxy numbers as a function of radius (in megaparsecs). The solid lines show the number of identifications at a given radius, and the shaded regions show the number after correcting for the estimated percentage of contamination in that annulus.
where N is the number of radio galaxies with luminosity L1.4GHz detected within the luminosity limits of the cluster at a radial distance d (where d 2.5 Mpc) and N2.5 is the number of galaxies with MR 21 and d 2.5 Mpc. This galaxy count, or richness, was measured using the same method as described above for N0.5. The radio fractions for the clusters are listed in Table 1 and shown graphically in Figures 4–6. There is a broad range in the radio fraction distributions for this sample, as first noted in the test study of A2645 and A2125 (Dwarakanath & Owen 1999; Owen et al. 1999). Table 1 also lists the breakdown of the radio fraction into luminosity class fLLRG and fHLRG. These fractions were computed in the same manner with the added restriction on the radio luminosity of log (L1.4GHz < 23.0) and log (L1.4GHz 23.0), respectively. Again, these fractions exhibit a relatively broad range in values, particularly the LLRGs. Of course, it must be kept in mind that the numbers presented here are upper limits because of the incompleteness in the redshift measurements. However, the magnitude limits placed on the sample greatly increase the probability that the radio galaxies are indeed cluster members.
4. RESULTS
4.1. Concentration Indices Although the imaging capabilities of the KPNO 0.9 m telescope do not allow for a detailed morphological analysis, an attempt was made to classify the systems as disk or elliptical galaxies by using concentration indices. The concentration index is a measure of the compactness of a galaxy as a function of its mean surface brightness. Specifically, the radii that enclose 30% and 70% of galaxy light were determined, and the ratio of these two radii yields the concentration index for the radio galaxy. This method has been used on images with 100 seeing to separate out the early- and late-type systems (Abraham et al. 1994; Owen et al. 1999). The concentration indices were determined using the IRAF task ELLIPSE in the STSDAS package. Elliptical models were fitted to the individual radio galaxies and the mean surface brightness (hli) was measured down to 24.5 mag arcmin2 (Owen & Laing 1989; Owen & White 1991). Contaminating sources were masked for the cases where a nearby star or galaxy was within the aperture. Some radio galaxies, however, were in very close projected pairs or
132
RIZZA ET AL.
Vol. 126
y
7
Surface Density (Mpc -2)
6 5
6
4 3 2
5
1 0 0.0
0.5
1.0 1.5 Distance from cluster center (Mpc)
2.0
2.5
Surface Density (Mpc -2)
2.5
4
3
2.0
2
1.5 1.0
1
0.5 0.0 0.0
0.5
1.0 1.5 Distance from cluster center (Mpc)
2.0
2.5
0 0.00
multiple systems. For these galaxies, it was not possible to accurately obtain their brightness distributions, and thus they were removed from the analysis. As shown in Figure 7, comparison of the concentration indices with the curves for an exponential and de Vaucouleurs profile facilitates classification of the galaxies as disks or E/S0’s. The concentration index analysis was accompanied by visual inspection, and for all classifiable cases the galaxy morphologies agreed with their concentration index typing. In a number of the cases, where the galaxies fell in the uncertain region between the two distributions, we were still able to categorize the
0.05
0.10
0.15
0.20
0.25
fRG
Fig. 3.—Surface density distribution for LLRGs and HLRGs. The HLRGs (top) are more centrally concentrated than the LLRGs (bottom). Both distributions are overlaid with the best-fitting King model.
Fig. 5.—Low-luminosity radio galaxy fractions ( frg for LLRGs)
galaxy upon visual inspection. The index value and mean surface brightness are listed in Table 5 for all the ‘‘ isolated ’’ radio galaxies. 4.2. Population-Age Classification of Radio Galaxies Expanding on the morphological classification, further categorization of the radio galaxy populations was attempted by pairing the optical morphologies with the 7
y 6
6 5
5 4
4 3
3
2
2
1
1
0 0.00
0.05
0.10
0.15
0.20
0.25
0 0.00
0.05
0.10
0.15
0.20
0.25
fRG
fRG
Fig. 4.—Total radio galaxy fractions ( frg for all radio galaxies)
Fig. 6.—High-luminosity radio galaxy fractions ( frg for HLRGs)
No. 1, 2003
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133
Fig. 7.—Concentration indices for the radio galaxies. The LLRGs and HLRGs are represented by crosses and filled circles, respectively, and the solid lines trace the distribution for an exponential (solid ) and de Vaucouleurs (dot-dash) profile.
optical colors of the galaxies. Actively star-forming systems are known to have blue colors, reflective of the young population of stars. In such systems, the emission lines may also dominate the blue spectrum and influence the integrated colors. Young stars are generally found in late-type or disk systems that contain substantial amounts of gas, providing fuel for ongoing star formation. On the other hand, earlytype or elliptical systems are predominantly composed of evolved, red stars. The lack of star formation in these galaxies implies that the radio emission is AGN in nature. Thus, in the absence of spectral information the combination of the galaxy color and morphology can yield a fair estimate of the stellar composition and therefore the nature of the radio emission. Using the color and morphology as
diagnostics, the radio sources were divided into three categories, as follows: 1. old: red galaxies with de Vaucouleurs profiles; 2. star-bursting (SB): blue galaxies with exponential profiles; 3. intermediate (int): blue galaxies with de Vaucouleurs profiles or galaxies with optical colors that are blueward of the mean cluster color (though not meeting the BO definition of blue; see x 2.2) paired with exponential profiles. The population classifications are listed in Table 5 for each of the radio identifications. Several radio sources were identified from the sky survey images and therefore do not
TABLE 5 Concentration Index Results ID (1)
R.A. (J2000.0) (2)
Decl. (J2000.0) (3)
z (4)
d (5)
C (6)
l (7)
BR (8)
Pop. (9)
0.59 0.05 1.14 0.32 1.43 0.92 1.96 0.93 0.88 1.66 0.38 2.04 0.79
0.514 0.612 0.572 0.514 0.606 0.528 0.350 0.551 0.541 0.565 0.450 0.434 0.573
21.82 22.50 22.76 22.51 22.16 22.49 22.48 22.64 21.43 22.78 22.25 22.76 22.32
2.25(0.05) 2.20(0.05) 2.16(0.05) 2.14(0.06) 2.26(0.05) 2.26(0.05) 1.59(0.05) 2.12(0.06) 1.50(0.05) 2.08(0.05) 1.48(0.05) 2.37(0.06) 2.22(0.05)
Old Old Old Old Old Old SB Old SB Old SB Old Old
0.14 0.21 0.46 1.97 1.64 0.16 0.92 0.52 0.80 0.29 0.74 1.99
0.534 0.639 0.521 0.522 0.457 0.525 0.517 0.611 0.597 0.612 0.413 0.665
22.41 22.05 21.81 22.55 21.94 21.92 22.25 22.46 22.33 22.20 21.81 22.54
2.42(0.05) 2.22(0.05) 1.71(0.05) 1.92(0.05) 1.56(0.05) 1.23(0.05) 2.28(0.05) 2.20(0.05) 2.34(0.05) 2.42(0.06) 1.48(0.05) 1.14(0.05)
Old Old SB Int SB Int Old Old Old Old SB Old
0.09 0.63 0.95 0.96 1.63 0.38 0.92 1.72
0.580 0.508 0.576 0.469 0.476 0.569 0.605 0.489
22.70 21.74 22.04 22.76 22.43 22.33 22.21 22.55
2.15(0.05) 1.64(0.05) 2.20(0.06) 2.46(0.06) 1.75(0.06) 1.31(0.05) 2.06(0.05) 1.90(0.06)
Old SB Old Old Int Int Old Int
1.18 1.58 1.37 0.08 1.28
0.639 0.681 0.635 0.605 0.430
22.20 21.29 22.56 22.36 21.99
2.33(0.06) 1.36(0.05) 1.30(0.05) 2.19(0.05) 1.90(0.06)
Old Int Int Old SB
0.60 1.17 0.39 0.15 0.74 0.55
0.556 0.561 0.612 0.535 0.517 0.434
22.31 22.43 22.55 22.34 22.17 22.98
2.14(0.05) 2.20(0.05) 2.11(0.05) 2.24(0.06) 2.17(0.06) 2.07(0.06)
Old Old Old Old Old Old
A209 1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12............... 13...............
01 31 50.80 01 31 52.45 01 31 34.87 01 31 55.60 01 31 57.38 01 32 13.48 01 31 11.34 01 31 33.07 01 31 33.56 01 31 36.76 01 31 47.12 01 31 38.96 01 31 44.58
13 33 35.89 13 37 00.21 13 32 19.94 13 38 30.04 13 44 42.90 13 38 09.32 13 40 49.83 13 38 28.79 13 38 01.82 13 45 07.63 13 38 21.69 13 47 38.53 13 40 42.31
0.206 0.213 ... 0.214 ... ... ... ... ... 0.199 ... ... ... A360
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12...............
02 31 27.60 02 31 23.56 02 31 15.58 02 30 42.47 02 31 38.38 02 31 22.85 02 31 32.48 02 31 35.24 02 31 41.65 02 31 30.69 02 31 08.92 02 30 57.07
06 58 55.93 07 00 12.62 06 58 10.09 06 59 19.29 06 51 07.06 06 58 24.88 06 54 40.34 06 58 18.17 06 59 21.71 06 58 44.72 06 58 45.83 06 51 04.34
0.220 ... ... ... ... ... ... ... 0.217 0.215 ... ... A383
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8.................
02 48 03.40 02 48 12.23 02 48 24.69 02 47 57.94 02 48 04.33 02 48 04.97 02 48 21.90 02 48 25.20
03 31 45.71 03 29 26.80 03 31 46.26 03 26 32.12 03 22 26.50 03 34 17.14 03 34 26.39 03 40 39.82
0.189 0.196 ... ... ... ... 0.186 ... A1111
1................. 2................. 3................. 4................. 5.................
10 50 55.66 10 51 03.71 10 51 06.32 10 50 36.53 10 50 05.78
02 30 35.47 02 28 53.78 02 32 05.66 02 36 17.10 02 40 43.53
0.166 ... ... 0.165 ... A1246
1................. 2................. 3................. 4................. 5................. 6.................
11 23 48.02 11 24 22.59 11 23 46.64 11 23 56.19 11 24 12.71 11 24 04.12
21 32 57.43 21 31 32.11 21 28 10.35 21 28 57.61 21 29 35.36 21 27 28.06
... 0.189 0.189 ... 0.190 0.187
134
TABLE 5—Continued ID (1)
R.A. (J2000.0) (2)
Decl. (J2000.0) (3)
z (4)
d (5)
C (6)
l (7)
BR (8)
Pop. (9)
1.76 0.97 0.11 1.36 0.95 0.76 0.24 0.26
0.469 0.468 0.650 0.584 0.397 0.529 0.213 0.677
22.62 22.95 22.61 22.29 22.27 22.55 22.72 22.92
2.86(0.08) 2.96(0.10) 2.67(0.06) 3.40(0.07) 2.28(0.06) 2.48(0.08) 2.56(0.06) 2.48(0.05)
Old Old Old Old Int Old Old Old
0.83 0.06 2.28 0.09 1.24 0.86 0.52 0.74 1.67 0.75
0.635 0.483 0.635 0.000 0.449 0.535 0.000 0.602 0.546 0.652
22.48 22.33 22.35 0.00 21.84 22.28 0.00 21.84 22.08 22.45
2.60(0.06) 2.65(0.05) 2.67(0.05) 2.76(0.06) 1.71(0.05) 2.02(0.06) 2.08(0.05) 1.76(0.06) 2.31(0.05) 2.67(0.05)
Old Old Old Old SB Int Old Int Old Old
1.15 0.46 0.29 0.39 1.21 0.09
0.613 0.614 0.626 0.633 0.591 0.591
22.57 22.56 22.95 22.30 22.12 21.99
2.41(0.05) 2.51(0.06) 2.40(0.06) 2.45(0.06) 1.97(0.05) 2.55(0.06)
Old Old Old Old Int Old
0.31 0.46 0.57 1.35 1.70 0.93 1.43 0.34 0.38 1.28 1.86
0.552 0.610 0.438 0.616 0.456 0.666 0.498 0.630 0.000 0.585 0.626
22.35 22.35 22.15 22.63 21.82 22.23 22.02 22.47 0.00 22.14 21.97
2.64(0.05) 2.49(0.06) 2.22(0.06) 2.35(0.06) 2.20(0.06) 2.53(0.05) 1.79(0.06) 2.42(0.06) 2.67(0.07) 2.28(0.05) 2.45(0.06)
Old Old Int Old Int Old SB Old Old Old Old
1.80 0.77 1.16 0.27 0.15 0.16 1.42 1.65 1.80 1.85 0.42 0.19 1.54 0.11
0.574 0.511 0.517 0.606 0.499 0.585 0.628 0.636 0.537 0.559 0.514 0.608 0.535 0.628
22.42 22.84 22.44 22.32 22.32 21.58 22.62 22.35 22.41 22.70 22.80 22.46 22.09 22.46
2.41(0.05) 2.69(0.06) 2.33(0.06) 2.34(0.05) 2.16(0.06) 1.70(0.05) 2.42(0.05) 2.42(0.05) 2.43(0.06) 2.35(0.05) 2.06(0.05) 2.47(0.05) 2.12(0.05) 2.41(0.06)
Old Old Old Old Old SB Old Old Old Old Old Old Int Old
A1550 1................. 2................. 3................. 4................. 5................. 6................. 7................. 8.................
12 28 35.99 12 28 49.60 12 28 58.36 12 28 19.84 12 28 31.38 12 28 46.43 12 29 00.83 12 29 02.46
47 45 28.56 47 42 22.31 47 38 29.15 47 37 48.46 47 37 57.25 47 34 58.35 47 36 53.56 47 36 56.14
... ... ... 0.256 ... ... ... 0.263 A1576
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10...............
12 36 53.46 12 36 58.49 12 38 26.12 12 36 57.42 12 36 40.65 12 36 45.94 12 36 53.11 12 37 22.89 12 38 02.22 12 37 28.35
63 14 35.38 63 11 13.31 63 12 17.26 63 11 16.51 63 15 54.35 63 14 24.41 63 09 04.99 63 09 12.67 63 09 44.06 63 10 27.63
... ... ... 0.302 ... ... ... ... ... 0.293 A1703
1................. 2................. 3................. 4................. 5................. 6.................
13 15 06.39 13 15 16.30 13 15 02.35 13 15 15.32 13 14 36.36 13 15 07.84
51 54 28.07 51 50 23.83 51 50 20.51 51 48 04.47 51 45 26.15 51 48 58.18
0.274 ... ... 0.275 ... ... A1758
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11...............
13 32 38.38 13 32 39.40 13 32 33.90 13 32 41.03 13 32 41.07 13 33 02.06 13 33 23.09 13 32 52.09 13 32 53.01 13 33 12.96 13 32 01.39
50 33 36.08 50 34 32.50 50 30 50.02 50 26 37.10 50 25 03.23 50 29 28.24 50 30 28.53 50 31 48.00 50 31 37.07 50 29 01.43 50 28 01.53
0.279 ... ... ... ... ... ... ... ... ... 0.278 A2219
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12............... 13............... 14...............
16 41 13.65 16 40 37.90 16 40 06.25 16 40 11.58 16 40 21.11 16 40 22.99 16 40 56.80 16 41 07.68 16 41 11.54 16 40 44.56 16 40 31.67 16 40 23.77 16 39 32.62 16 40 15.93
46 42 04.19 46 45 08.30 46 48 23.16 46 43 17.97 46 43 24.47 46 42 17.95 46 39 10.22 46 40 37.45 46 40 03.46 46 34 07.17 46 42 22.25 46 42 10.41 46 41 44.98 46 42 31.68
... ... ... ... 0.219 ... 0.227 ... ... ... ... 0.229 ... 0.228
135
TABLE 5—Continued ID (1)
R.A. (J2000.0) (2)
Decl. (J2000.0) (3)
z (4)
d (5)
C (6)
l (7)
BR (8)
Pop. (9)
2.17 0.06 0.92 2.04 1.50 0.42
0.608 0.397 0.627 0.488 0.676 0.405
22.36 22.58 22.28 21.98 22.23 21.91
1.96(0.05) 1.59(0.05) 2.07(0.05) 1.50(0.05) 1.99(0.05) 1.49(0.05)
Old SB Old SB Old SB
0.83 0.06 0.90 1.75 1.18 0.56
0.293 0.492 0.351 0.468 0.589 0.573
22.45 23.16 22.76 21.80 21.99 22.50
1.91(0.06) 2.41(0.07) 1.86(0.06) 1.95(0.05) 1.97(0.05) 2.42(0.05)
SB Old SB Int Int Old
1.44 0.98 0.13 0.49 0.66 0.13 0.19 1.73 2.10 1.99 2.00 1.06
0.520 0.474 0.580 0.570 0.579 0.575 0.537 0.656 0.605 0.616 0.683 0.000
22.67 22.711 21.912 21.851 22.527 22.203 22.577 22.491 22.597 22.310 22.547 0.000
2.42(0.09) 1.93(0.09) 1.70(0.05) 1.71(0.06) 2.15(0.07) 2.43(0.08) 1.78(0.06) 1.80(0.06) 2.26(0.06) 2.37(0.06) 1.56(0.06) 2.10(0.08)
Old Int Int SB Old Old Int Int Old Old Int Old
0.08 0.14 1.47 0.61 0.69 0.78 1.26 1.92 2.20 2.21 1.71 1.30 2.21 2.14 2.01 2.45
0.392 0.566 0.532 0.494 0.411 0.601 0.662 0.504 0.567 0.620 0.393 0.644 0.388 0.593 0.575 0.000
22.65 22.35 22.46 22.61 22.66 22.87 22.51 22.45 21.47 22.55 21.87 22.23 22.95 22.39 22.46 0.00
2.21(0.07) 2.26(0.06) 2.48(0.06) 2.17(0.07) 2.53(0.08) 2.53(0.06) 2.32(0.06) 2.17(0.06) 0.90(0.05) 2.04(0.05) 1.29(0.05) 1.55(0.05) 1.98(0.07) 2.10(0.05) 2.28(0.06) 2.36(0.10)
Old Old Old Old Old Old Old Old SB Int SB SB Int Old Old Old
A2254 1................. 2................. 3................. 4................. 5................. 6.................
17 18 20.47 17 17 45.53 17 17 30.61 17 18 36.33 17 17 24.27 17 17 35.91
19 50 57.16 19 40 39.22 19 35 48.52 19 37 41.89 19 32 28.34 19 38 47.17
... ... 0.179 ... 0.173 ... A2294
1................. 2................. 3................. 4................. 5................. 6.................
17 23 59.63 17 24 13.40 17 28 01.62 17 32 22.56 17 17 56.84 17 22 48.31
85 58 22.77 85 53 33.36 85 57 02.93 85 59 48.20 85 56 15.04 85 56 21.51
... ... ... ... ... ... A2390
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12...............
21 54 03.00 21 53 18.62 21 53 36.82 21 53 45.59 21 53 49.68 21 53 35.19 21 53 31.42 21 52 59.83 21 54 07.63 21 53 29.21 21 53 36.91 21 53 46.99
17 37 41.66 17 44 18.19 17 41 43.69 17 41 47.82 17 41 20.50 17 41 50.48 17 41 33.56 17 38 14.71 17 33 29.55 17 30 59.51 17 30 49.00 17 36 28.15
... 0.234 0.230 0.215 ... ... ... ... ... ... ... ... A2552
1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12............... 13............... 14............... 15............... 16...............
23 11 33.13 23 11 35.60 23 11 41.09 23 11 46.28 23 11 48.66 23 11 48.34 23 11 53.11 23 10 54.81 23 10 55.93 23 11 02.16 23 11 07.67 23 11 42.20 23 11 51.08 23 10 52.80 23 10 54.12 23 11 37.71
03 38 53.07 03 39 18.80 03 46 04.44 03 39 42.76 03 38 55.17 03 40 45.12 03 43 11.39 03 39 01.03 03 44 40.47 03 46 41.22 03 44 21.66 03 32 15.99 03 28 03.02 03 35 02.50 03 36 19.41 03 26 02.16
... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
z 0.2 RICH ABELL CLUSTERS
137
TABLE 5—Continued ID (1)
R.A. (J2000.0) (2)
Decl. (J2000.0) (3)
z (4)
d (5)
C (6)
l (7)
BR (8)
Pop. (9)
0.55 1.57 0.90 0.08 0.14 1.94 1.43 0.96 0.46 1.50 2.07 0.53 0.71 2.12 1.69 1.56 1.36
0.627 0.416 0.602 0.584 0.592 0.420 0.580 0.565 0.442 0.484 0.564 0.592 0.586 0.564 0.526 0.463 0.631
22.68 22.10 22.51 22.73 23.12 22.42 22.08 22.52 21.91 22.82 22.12 22.49 22.14 22.21 22.75 21.59 22.08
2.28(0.06) 2.58(0.10) 1.65(0.06) 2.37(0.05) 2.09(0.06) 1.74(0.06) 2.26(0.06) 2.42(0.06) 1.42(0.05) 2.22(0.07) 1.82(0.05) 2.28(0.06) 2.29(0.06) 1.70(0.05) 2.65(0.09) 0.82(0.05) 2.26(0.06)
Old Int Int Old Old SB Int Old SB Old Int Old Old Int Old SB Old
A2631 1................. 2................. 3................. 4................. 5................. 6................. 7................. 8................. 9................. 10............... 11............... 12............... 13............... 14............... 15............... 16............... 17...............
23 37 41.34 23 37 54.39 23 37 56.58 23 37 39.73 23 37 40.12 23 37 07.14 23 37 15.08 23 37 23.62 23 37 32.84 23 37 33.39 23 37 35.66 23 37 48.00 23 37 51.97 23 38 04.26 23 38 10.17 23 37 12.70 23 37 19.05
00 18 33.80 00 22 20.89 00 17 18.36 00 16 15.98 00 16 40.00 00 18 28.62 00 16 19.84 00 16 20.81 00 16 33.79 00 22 37.14 00 25 22.27 00 14 25.50 00 14 21.88 00 08 26.81 00 13 40.51 00 16 38.27 00 13 14.26
... ... ... 0.278 ... ... ... ... ... ... ... ... ... ... ... ... ...
Note.—Col. (1): galaxy ID; col. (2): right ascension; col. (3): declination; col. (4): redshift; col. (5): projected linear distance from cluster center; col. (6): concentration index; col. (7): mean surface brightness; col. (8): galaxy color; col. (9): population classification = old, star bursting (SB), intermediate (Int).
have either color or morphological information; no population classification was attempted for these cases. The results of the population analysis are presented in several different ways. Figure 8 displays the radio power distribution as a function of the stellar population. As expected, the HLRGs have predominantly old stellar populations, indicative of AGN radio emission. The LLRGs show considerably more diversity, with the majority of the sources characterized as old populations. We depict these results slightly differently in Figure 9. Here the radio galaxies are binned according to population and normalized by the number of sources in the luminosity class. For example, in Figure 9 the ‘‘ old ’’ population bin tells us that 80% of
Old
SB
the HLRGs fall into this category, as do 47% of the LLRGs. The decomposition of the overall sample is consistent with the population statistics for the more thoroughly studied case of A2125 (Owen et al. 1999). About 10% more old LLRGs were found in this sample relative to A2125 (47% compared with 38%), but the star-forming fractions are comparable, at 25%. This consistency is reassuring, given the different classification criteria in the two studies, and it lends support to the use of broadband colors plus optical morphology for population typing when faced with a lack of spectroscopic data. Further study of the composition of the LLRGs yields some suggestive trends. Figure 10 shows the background corrected fraction of the LLRGs that are ‘‘ old ’’ ( fOLD), ‘‘ SB ’’ ( fSB), and ‘‘ int ’’ ( fINT) plotted versus the richness of the clusters. The figure illustrates the positive correlation
Int
1.0 30 HLRG LLRG
0.8
20
0.6
10
0.4
0.2 0 22
23
24
25
Log(L 1.4 GHz) 0.0
Fig. 8.—Population distribution for entire sample. The high-luminosity end of the radio power distribution is almost exclusively composed of old populations. The low-luminosity end is much more diverse, consisting of old, intermediate, and star-forming populations.
Old
SB Population
Int
Fig. 9.—Fractional population breakdown for entire sample
138
RIZZA ET AL. 60
Vol. 126
60 1703
1703 50
50 2390 2219 209 2294
30
20
2390 1246 2219 1758
40 N 0.5
N 0.5
40
1246 1758
2254 383 1576 2552 2631 360
2254 30
1576
20
10
209 2294
383 2552 360 2631
10 0
20
40 60 80 Fraction "OLD"
100
120
0
60
20
40 60 80 Fraction "SB"
100
120
60 1703
1703
50
50
N 0.5
40
1246 1758
2219 40 N 0.5
1246 1758 209
2390
2294 2254 30
2552 360
2631
383 1576
20
2254 1576 383 2552 360 2631
30
20
10
2390 2219 209 2294
10 0
20 40 Fraction "INT"
60
0
20
40 60 80 100 Fraction "SB+INT"
120
Fig. 10.—Richness vs. fractional composition of the LLRGs. Note the correlation between N0.5 and fOLD and the anticorrection with fSB and fSB+INT.
between N0.5 and fOLD (see Table 6). As the richness of the cluster increases, so does the dominance of old populations in the composition of the LLRGs. The second panel in Figure 10 shows the other side of the same trend: the fraction of SB galaxies anticorrelates with the richness of the cluster. The intermediate population also suggests a similar trend. Finally, one can see that the anticorrelation also exists for the combination of star-forming and intermediate populations. It is likely that many of the intermediate objects are
actually star-forming systems, either without obvious or dominant disks (as in the case of S0’s) or with blue colors but uncertain morphological classifications from the concentration index analysis. Thus, it is of value to examine the ‘‘ SB+int ’’ as a ‘‘ younger ’’ population (relative to that defined as ‘‘ old ’’). The Spearman’s rank correlation coefficients and level of significance of the trends are listed in Table 6. The results support a correlation of fOLD and richness (and its complement, the anticorrelation of fSB+INT with N0.5).
TABLE 6 LLRG Composition/Richness Correlations
4.3. Radio Fraction Correlations
Relation (1)
(2)
Significance (%) (3)
N0.5 vs. fOLD ......................... N0.5 vs. fSB ............................ N0.5 vs. fINT .......................... N0.5 vs. fSB+INT .....................
0.7566 0.5556 0.5216 0.7566
99.72 95.13 93.25 99.72
Note.—Col. (1): relation; col. (2): correlation coefficient; col. (3): significance.
The radio fractions are plotted as a function of the cluster richness in Figure 11. Both the X-ray–bright sample (including the clusters that overlap with R1: A773, A1704, A2111, A2645) and the z = 0.1–0.25 clusters from R1 are included in the analysis to increase the signal-to-noise ratio. The former are shown as filled circles, and the latter as asterisks. The R1 sample is not necessarily X-ray faint, but the clusters simply were not selected to be X-ray bright. In addition, R1 includes some clusters with redshifts that are lower than initially targeted for the X-ray–bright sample. The Spearman’s
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correlation coefficients for the various comparisons are listed along with their significances in Table 7 for the X-ray– bright and combined samples. The radio fractions used in the subsequent analysis are the ‘‘ corrected ’’ fractions from x 3.3. These corrections have also been applied to the R1 clusters by neglecting their redshifts and starting from their original identification list. This was done in order to treat all of the clusters similarly and to correct for the incomplete redshift coverage of some of the non–X-ray–bright clusters. One suggestive relationship between the cluster richness and the overall radio fraction emerges in both the X-ray– bright and combined sample. If the radio detection rate
were simply a function of the number of galaxies surveyed, one would expect fRG to correlate with the richness measurements N0.5. As Figure 11 illustrates, this does not seem to be the case. Instead, there is a rather strong anticorrelation between the two. This is confirmation of the trend noted in R2 and means that the number of radio galaxies detected is not simply a function of the number of galaxies counted in the cluster. The anticorrelations with the radio luminosity–derived fractions ( fHLRG, fLLRG, and fSBRG) are also suggestive. This implies that the galaxies are more likely to be radio bright, whether they are AGN or star-forming, in lower richness clusters. Interestingly, this trend seems to
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TABLE 7 Radio Fraction Correlations X-Ray–Bright Sample
Combined Sample
Relation (1)
(2)
Significance (%) (3)
(4)
Significance (%) (5)
N0.5 vs. fRG .................. N0.5 vs. fHLRG .............. N0.5 vs. fLLRG............... N0.5 vs. fSBRG ............... N2.5 vs. fRG .................. N2.5 vs. fHLRG .............. N2.5 vs. fLLRG............... N2.5 vs. fSBRG ............... fB vs. fRG ..................... fB vs. fHLRG ................. fB vs. fLLRG .................. fB vs. fSBRG ..................
0.5462 0.4791 0.6288 0.5504 0.2922 0.3201 0.4244 0.3528 0.0176 0.0041 0.0186 0.0456
98.10 95.57 99.48 98.21 76.07 80.46 92.08 84.90 5.53 1.30 5.85 14.27
0.4787 0.4548 0.4037 0.3754 0.1837 0.3148 0.1159 0.1175 0.1143 0.1987 0.0637 0.0171
99.44 99.11 97.81 96.58 68.58 92.07 47.24 47.80 46.67 72.43 27.09 7.40
Note.—Col. (1): relation; col. (2): correlation coefficient for all clusters selected to be X-ray bright; col. (3): significance; col. (4): correlation coefficient for the combined sample; col. (5): significance.
extend down to richness class 0 systems and poor groups of galaxies. Preliminary analysis of a sample of 300 poor clusters also finds a strong anticorrelation between N0.5 and fRG at greater than 99% significance (Bliton 2000). The anticorrelation is not as strong when the radio fractions are compared with the richness measured at 2.5 Mpc (N2.5). This may in part be due to the increase in possible contamination at larger radii from the cluster core. A larger degree of contamination will increase the scatter in the relationship, diluting the trend. However, the lack of a strong correlation may also be an indication of a larger problem created by the lack of confirmed redshifts for the radio galaxy identifications, and this why we are only able to state that the trend is suggestive. The last series of plots (Fig. 12) show the relationship between the radio fractions and the cluster blue fraction. The test study of A2125 and A2645 (Dwarakanath & Owen 1999; Owen et al. 1999) suggested that fB and fRG were linked in these two instances and perhaps generally for BO clusters (Butcher & Oemler 1978b, 1984). A more extended sample, however, did not support this overall trend (R2). Instead there seems to be a broad range of fRG for intermediate-redshift to distant clusters, as is seen in Figure 12. Neither the X-ray–bright nor the combined sample shows significant correlations (see Table 7). 5. DISCUSSION
We have confirmed the result of R2 that the fraction of radio galaxies anticorrelates with the richness of the cluster for both the large combined sample and the X-ray–bright sample. This trend implies that a galaxy in a richer cluster is less likely to be radio-bright whether the emission is due to star formation or AGN activity. In keeping with this, the very high radio galaxy fraction measured for the test case A2125 is not caused by sampling more galaxies; instead, the high fraction seems to be due to the cluster environment. These results go hand in hand with the correlation between N0.5 and fold for LLRGs discussed above. Dynamical processes, such as galaxy interactions, rampressure stripping and harassment or tidal influence from
the cluster potential, all efficiently remove the gas from the outer regions of the galaxy and truncate star formation. This reduces the fraction of star-forming systems after an initial burst at earlier epochs (as seen in BO clusters and in the radio results for z 0.4 clusters from Morrison et al. 2003c, hereafter R4). Studies also show that interactions/ mergers appear to be extremely efficient at funneling large amounts of gas into the merger nuclei (e.g., Barnes & Hernquist 1991, 1992; Mihos & Hernquist 1994), which could eventually lead to powering an AGN. Galaxy harassment has the same consequence of efficient funneling of gas into the central regions of galaxies (Lake, Katz, & Moore 1998). This model even predicts an increase in the AGN population in BO clusters where galaxy harassment appears to be in effect. Abadi, Moore, & Bower (1999) find that rampressure stripping will also deplete the gas in the outer regions of the galaxy on relatively short timescale, while the inner core or bulge will retain its gas leading to nuclear starbursts or AGNs. In all of these scenarios, the galaxies first undergo a period of activity before evolving into a more dormant state, as is well known from studies of the color (Butcher & Oemler 1984; Rakos & Shombert 1995) and morphological evolution (Dressler et al. 1994a, 1994b, 1997; Couch et al. 1994, 1998) of clusters. Regardless of the nature of the radio emission, the lifetimes of the sources are expected to be short: 108 yr for star formation (Condon 1992) and approximately less than a few times 108 yr for AGN emission (e.g., Martini & Weinberg 2001). In popular CDM models of structure formation the central regions of rich galaxy clusters form at z e 1 (Kravtsov & Klypin 1999). This is supported by direct observations of rich clusters at z 1 (e.g., Postman et al. 1996; Luppino & Kaiser 1997). Since dynamical processes either induce activity that consumes the gas in the galaxies or strip the gas away, it is natural to expect suppression of both AGN and starbursting radio galaxies in the central 0.5–1 Mpc of rich clusters after several billion years (i.e., by z 0.2). The study of radio galaxies presented here shows a link between the richness of the cluster and the degree of activity in the constituent galaxies. It therefore appears that the fraction of radio galaxies in clusters is a not
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only a function of the dynamic state of the system (i.e., merging or infall). It may also depend on the evolutionary state of the cluster galaxies that can be related back to the richness of the system. Our results indicate that the radio galaxy fractions are not correlated with the cluster blue fraction. Initial expectations of a relationship between fRG and fB were motivated by the trial study of A2125 and A2645; yet analysis by M99 and R3 and R1 also found a lack of an overall correlation, despite a few suggestive individual cases. As noted in their study, A2125 appears to be an unusual cluster and may be experiencing ongoing merger activity which is enhancing low-level
radio emission. A further explanation may be that the blue fractions are measured in the inner region of the cluster (see x 2.2), whereas the radio galaxies are distributed over a much broader spatial extent. Several of our clusters have particularly low fB in their cores, which accentuates this discrepancy. In addition, through the radio observations we are probing very recent star formation, since the lifetimes of the electrons are on the order of 108 yr (Condon 1992). In comparison, the blue galaxy colors have a longer duration. We are also including AGNs in our radio fractions, whereas the blue fractions only trace the star-forming systems. And, finally, both fB and fRG are relatively noisy measurements, which, when
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compared with one another, would decrease the signal in a possible relationship. One important point that merits reiteration is the fact that we do not have complete redshift confirmation of cluster membership for all of the radio identifications. Corrections for this were attempted by subtracting an average background radio galaxy surface density scaled to the cluster angular size (see x 3.3); yet this is just one model representing the possible contamination. In adopting this model, it was assumed that the background objects have the same probability of being radio sources as the cluster galaxies. This assumption may be an overestimate, leading to a conservative treatment of the radio galaxy errors. However, if the degree of contamination has been underestimated, it may in fact conspire to produce the anticorrelation we observe between the radio fraction and core richness. The lack of redshifts dominates our uncertainty as well as contributes to the intrinsic scatter in the correlations. Another point to consider is that the anticorrelation is strongest for N0.5, which measures the core richness of the cluster. Although this richness count was initially adopted since it is least likely to suffer from background contamination, the radio fractions were calculated by determining the radio galaxy distribution within 2.5 Mpc in radius from the cluster center. Since the LLRGs tend to reside in the outer regions of the system, by using N0.5 we may be comparing the radio galaxies to a slightly different galaxy population than is representative of the richness count. This could therefore be complicating the interpretation of an anticorrelation. These issues can be addressed by obtaining redshift confirmation of cluster membership for the radio identifications and until this is done, the correlations found in our analysis are to be considered preliminary.
6. CONCLUSIONS
We have carried out an analysis to characterize the nature of the radio galaxy emission in a sample of intermediate-
redshift clusters using optical colors and morphological classifications of the host galaxies. We find that the LLRGs have a much broader spatial distribution than the HLRGs, in agreement with R2. These two distributions are statistically different at 99.99% level of significance. A first-order comparison of the radio distributions with the optical analysis implies that the HLRGs are consistent with the red galaxy population and the LLRGs with the spatial distribution of the blue galaxy population. The LLRGs are composed of a more diverse mix of galaxy populations/ages. The surface density of these sources shows that they have a broad spatial distribution and are not concentrated near the core, as are the HLRGs. About 25% of these sources reside in blue, disky galaxies, which are classically associated with star formation and BO-type objects. Yet almost half of these sources reside in red, elliptical galaxies, implying a possible AGN connection. The number of old LLRGs correlates with the richness, whereby higher richness clusters have higher fractions of old LLRGs. This implies that the cluster environment influences the age or degree of processing of the galaxies and may be a factor in determining the nature of the low-luminosity radio galaxy emission. We find that richer clusters tend to have lower radio fractions. This trend, first noted in M99, is detected for both the combined and X-ray–bright samples and implies that at z d 0.25 galaxies residing in less rich clusters are more likely to be radio bright. However, there does not appear to be a correlation between the radio fractions and the cluster blue fractions. We would like to thank Jon Holtzman and Dave Summers for their help in obtaining calibration frames for several of our cluster fields. We would also thank National Optical Astronomy Observatory and National Radio Astronomy Observatory for their support of these observations and for the development of the software used to analyze the data. This work has been supported by a grants from the National Science Foundation and NASA to J. O. B.
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