submitted to ApJS

Subaru/XMM-Newton Deep Survey (SXDS) : Optical Imaging
and Photometric Catalogs 1
Hisanori Furusawa1 , George Kosugi2 , Masayuki Akiyama1 , Tadafumi Takata2 , Kazuhiro
Sekiguchi2 , Ichi Tanaka1 , Ikuru Iwata3 , Masaru Kajisawa2 , Naoki Yasuda4 , Mamoru Doi5 ,Masami Ouchi6,7 , Chris Simpson8 , Kazuhiro Shimasaku9,10 , Toru Yamada11 , Junko
Furusawa1 , Tomoki Morokuma2 , Catherine M. Ishida1 , Kentaro Aoki1 , Tetsuharu Fuse1 ,Masatoshi Imanishi2 , Masanori Iye2 , Hiroshi Karoji2 , Naoto Kobayashi5 , Tadayuki
Kodama2 , Yutaka Komiyama1,2 , Yoshitomo Maeda12 , Satoshi Miyazaki1,2 , YoshihikoMizumoto2 , Fumiaki Nakata2 , Jun’ichi Noumaru1 , Ryusuke Ogasawara2 , Sadanori
Okamura9,10 , Tomoki Saito13 , Toshiyuki Sasaki1 , Yoshihiro Ueda14 , and Michitoshi Yoshida3

–2–
ABSTRACT
We present multi-waveband optical imaging data obtained from observations of the Subaru/XMM-Newton Deep Survey (SXDS). The survey ﬁeld, centered at R.A. = 02h 18m 00s , decl. = −05◦ 00 00 , has been the focus of a wide
range of multi-wavelength observing programs spanning from X-ray to radio
wavelengths. A large part of the optical imaging observations are carried out
with Suprime-Cam on Subaru Telescope at Mauna Kea in the course of Subaru Telescope “Observatory Projects”. This paper describes our optical observations, data reduction and analysis procedures employed, and the characteristics of the data products. A total area of 1.22 deg2 is covered in ﬁve contiguous sub-ﬁelds, each of which corresponds to a single Suprime-Cam ﬁeld of
1

Subaru Telescope, National Astronomical Observatory of Japan, 650 North A’ohoku Place, Hilo, HI
96720 USA; furusawa@subaru.naoj.org
2

National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588 Japan

3

Okayama Astrophysical Observatory, National Astronomical Observatory of Japan, 3037-5 Honjo, Kamogata, Asakuchi, Okayama 719-0232 Japan
4

Institute for Cosmic Ray Research, University of Tokyo, 5-1-5 Kashiwa-no-Ha, Kashiwa, Chiba 277-8582
Japan
5

Institute of Astronomy, School of Science, University of Tokyo, 2-21-1 Osawa, Mitaka, Tokyo 181-8588
Japan
6

Observatories of the Carnegie Institution of Washington, 813 Santa Barbara Street, Pasadena, CA 91101

7

Carnegie Fellow

8

Astrophysics Research Institute, Liverpool John Moores University, Twelve Quays House, Egerton Wharf, Birkenhead CH41 1LD UK
9

Department of Astronomy, School of Science, University of Tokyo, 7-3-1 Hongo, Bunkyo, Tokyo 113-0033
Japan
10

Reseach Center for the Early Universe, School of Science, University of Tokyo, 7-3-1 Hongo, Bunkyo,
Tokyo 113-0033 Japan
11

Astronomical Institute, Graduate School of Science, Tohoku University, Aramaki, Aoba, Sendai, Miyagi
980-8578 Japan
12

Institute of Space and Astronautical Science, Japan Aerospace Exploration Agency, Sagamihara, Kanagawa 229-8510 Japan
13

Physics Department, Graduate School of Science, Ehime University, 2-5 Bunkyou, Matsuyama,Ehime
790-8577 Japan
14

Department of Astronomy, Kyoto University, Sakyo-ku, Kyoto 606-8502 Japan

–3–view (∼ 34 × 27 ), in ﬁve broad-band ﬁlters B, V, Rc , i , z  to the depths of
B = 28.4, V = 27.8, Rc = 27.7, i = 27.7 and z  = 26.6 (AB, 3 σ, φ = 2 ). The
data are reduced and compiled into ﬁve multi-waveband photometric catalogs,
separately for each Suprime-Cam pointing. The i -band catalogs contain about
900,000 objects, making the SXDS catalogs one of the largest multi-waveband
catalogs in corresponding depth and area coverage. The SXDS catalogs can beused for an extensive range of astronomical applications such as the number density of the Galactic halo stars to the large scale structures at the distant universe.
The number counts of galaxies are derived and compared with those of existing
deep extragalactic surveys. The optical data, the source catalogs, and conﬁguration ﬁles used to create the catalogs are publicly available via the SXDS web
page (http://www.naoj.org/Science/SubaruProject/SXDS/index.html).
Subject headings: cosmology: observations — large-scale structure of universe —galaxies: evolution — galaxies: formation — galaxies: photometry

1.

INTRODUCTION

Understanding the formation and evolution processes of the individual galaxy and thegrowth of the large scale structures (LSSs) in the universe is one of the major goals in
extragalactic astronomy today. In the scheme of a typical Λ-CDM model, the growth ofstructures is governed by the gravitational growth of initial ﬂuctuations of dark matter.
The baryonic material cools in these dark matter structures and grows through hierarchicalclustering to galaxies and clusters of galaxies. The subsequent evolution of galaxies is closely
connected to their environments, which, of course, relate to the LSSs, where those galaxies arelocated. The ultimate goal of extragalactic surveys is to trace these evolutionary processes
by well-deﬁned statistical galaxy samples.
Optical imaging is arguably the cornerstone of any extragalactic survey, since it provides
identiﬁcations and positions of celestial objects for follow-up spectroscopy. In recent years,
there have been a number of deep imaging survey projects which devote signiﬁcant amountsof telescope time, such as Great Observatories Origins Deep Survey (GOODS: Giavalisco etal. 2004) and Cosmic Evolution Survey (COSMOS: HST treasury project: Feldmann et al.
2006; Scoville et al. 2007). These surveys provide multi-waveband galaxy samples at faint1

Based on data collected at Subaru Telescope, which is operated by the National Astronomical Observatory of Japan.

–4–
magnitude. Then, the photometric redshift techniques (Furusawa et al. 2000; Bolzonella etal. 2000; Feldmann et al. 2006) are frequently used to pre-select candidates for spectroscopy.
Although a great deal of information can be obtained from the optical data alone, thevalue of the data set grows signiﬁcantly as data at other wavelengths from other facilities
are added. We therefore elected to use a multi-wavelength approach for the Subaru/XMMNewton Deep Survey (SXDS; Sekiguchi et al. 2007, in preparation) from the very start,
ensuring that our chosen ﬁeld would be accessible and suitable for observations at all wavelengths. An equatorial ﬁeld, centered at R.A. = 02h 18m 00s , decl. = −05◦ 00 00 is chosen totie up with the deep X-ray observations with XMM-Newton observatory (Ueda et al. 2007),and also due to the accessibility by all major observatories, both existing and planned. The
major multi-wavelength observation programs on the SXDS ﬁeld (hereafter SXDF) include
deep radio imaging with the VLA (Simpson et al. 2006), sub-mm mapping with SCUBA(Coppin et al. 2006; Mortier et al., 2005), mid-infrared observations with Spitzer Space
Telescope (Lonsdale et al. 2003), deep near-infrared imaging with the UKIRT/WFCAM(Foucaud et al. 2006; Dye et al. 2006; Lawrence et al. 2007), and the X-ray observations with XMM-Newton observatory. Importantly, the survey ﬁeld of an infrared ultra deep
survey (UDS; Foucaud et al. 2006) covering 0.77 square degrees as part of the UKIDSS
project (Lawrence et al. 2007) is centered on the SXDF. It is expected that our extensive
multi-wavelength data set will provide photometric redshifts accurate to ∆z  0.1 over awide range of redshift, as well as detailed spectral energy distributions for the vast majorityof the objects in the ﬁeld.
The SXDS has been undertaken as a part of the Subaru Telescope “Observatory Projects”,
in which a large amount of observing times are devoted to carry out intensive survey programs by combining observing times rewarded to builders and observatory staﬀ of the Subaru
Telescope. Note that Subaru Deep Field (SDF; Kashikawa et al. 2004) is another observatory
project, which targets on a ﬁeld diﬀerent from the SXDF. The data set of the SDF is slightly
deeper than the SXDF, but concentrates on one-ﬁfth narrower ﬁeld than the SXDF. Thus,
these surveys complement each other and they can be used for a wide variety of studies.
In this paper, we describe the observation details and data reduction and analysis of
our optical imaging survey with the prime-focus camera on Subaru Telescope, as well as theresultant data products. A detailed description of the survey strategies, scientiﬁc objectives,and multi-wavelength survey plans is given in a companion paper (Sekiguchi et al., 2007, in
preparation). In Section 2, we describe optical imaging observations by Subaru Telescope,and in Section 3 we explain the data reduction procedure employed. We present creation ofmulti-waveband photometric catalogs including object detection and aperture photometryin Section 4, and calibrations of photometry and astrometry in Section 5. In Section 6, we–5–
discuss characteristics of the catalogs such as limiting magnitude, detection completeness,
and number counts of galaxies. The summary is given in Section 7. Throughout this paper,
all the magnitudes are expressed in the AB system unless otherwise mentioned.

2.

OBSERVATIONS AND DATA

The optical imaging observations of the SXDF are carried out using the prime-focus camera (Suprime-Cam; Miyazaki et al. 2002) on Subaru Telescope in the period from September
2002 to September 2005. Suprime-Cam is a 5 × 2 mosaic of ten MIT/LL buttable 2k×4k
CCD camera. It has a 34 × 27 ﬁeld of view with a projected pixel scale of 0.202 arcsec. The
layout of the pointings is arranged as a cross shape so that each of the North-South and theEast-West directions has an extent of ∼ 1.3 deg (Figure 1). This corresponds to a ﬁeld spanof a transverse dimension of ∼ 75 Mpc at z ∼ 1 and ∼ 145 Mpc at z ∼ 3 in the comoving
scale for (h, ΩM , ΩΛ ) = (0.7, 0.3, 0.7). The coverage of the comoving volume is ∼ 3.2 × 107
Mpc3 for z = 0 − 3. Table 1 lists ﬁeld center coordinates of the ﬁve pointings. Hereafter, theﬁve pointings are referred to as C (Center), N (North), S (South), E (East), and W (West),
with respect to their relative positions on the sky.
Observations of the SXDF are performed using ﬁve broad-band ﬁlters, B, V , Rc , i , andz  to cover the entire wavelength range observable with Suprime-Cam. The transmission
curves of the ﬁve ﬁlters employed in the SXDS, convolved with the CCD sensitivities, reﬂectivity of the primary mirror, transmissions of corrector lenses, and atmospheric transmission,
are given in Figure 2. Also, given in Figure 2 are the redshifted spectral energy distributions(SEDs) of an early-type galaxy at z = 0.5 and 1.2 and a late-type galaxy at z = 0.5 and
5.0 for references. With photometric data sampled in the 5 bands, we can investigate color
properties of red passively evolving galaxies to z ∼ 1 (see Kodama et al. 2004; Yamada et al.
2005) and young star-forming galaxies at high redshift such as lyman break galaxies (LBG)(Ouchi et al. 2004, 2005). The observations used in this paper span observing runs over a
period of 3 years. A total of 160 hours are allocated to the SXDF imaging observations as apart of Subaru Telescope “Observatory Projects” (see also Kashikawa et al. 2004) combinedwith observatory staﬀ times. In addition to these times, we use 14 hours of the observing
times oﬀered by the Supernova Cosmology Project (Doi et al. 2003; Lidman et al. 2005),
which add the total exposure time of the center and the west pointings in the i band (i -C
and i -W). Moreover, in the course of another survey program for Lyα emitting galaxies(Ouchi et al. 2005, 2007), 3 hours are devoted to the V -band imaging. In total, 133 hours
are used for on-source exposures for the ﬁnal data set described in this paper. The complete
log of the observations is given in Table 2.

–6–
The coordinates of the ﬁve pointings are carefully chosen so that the resultant images
uniformly cover the entire SXDF, though we analyze each of the pointings separately inthe present study. Also, to eliminate bad pixels, gaps between individual CCD chips andcosmic ray events, circular dithering pattern is used for each pointing. This dithering pattern
employs a circular motion with radii of 60 to 120 arcsec around the center position of eachpointing, combined with slight oﬀsets to the center positions of the circular motion. As an
example, the dithering pattern for the Rc band is plotted for each pointing in Figure 3.

3.

DATA REDUCTION

The SXDS images obtained are processed using an in-house pipeline software package(See Ouchi 2003). We use it in almost the same way that successfully used for the SDF(Kashikawa et al. 2004). This pipeline is based on a software package developed for SuprimeCam data reduction (Yagi et al. 2002) and several IRAF tasks (geomap, geotran, etc.) areused in the process of a geometrical transformation and alignment of ﬁnal stacked images. We
use a standard mosaic-CCD data reduction procedure, which is brieﬂy summarized below.
First, we apply a bias correction to the raw data. The median count of the overscan
region is computed for each line, which represents a typical bias level at that line, and thismedian count is subtracted from the counts in all the pixels in that line. The overscan regions
are trimmed after the bias subtraction. We assign a ﬂag number (-32768) to the saturated or
detected bad pixels, and these pixels are ignored in the following reduction processes. Since
pixels in the outside edges of each frame are likely to be aﬀected by noises which can not becorrected by the following ﬂatﬁelding process, they are also masked with the ﬂag number.
Flat frames are constructed from the normalized object frames taken during the sameobserving run. We use more than 150 object frames to create ﬂat frames. If the numberof object frames is less than 150 frames, object frames taken in diﬀerent observing runs
within 3 months or twilight ﬂat frames taken in the same run are combined to increase thesignal-to-noise ratios (S/Ns) of the ﬂat frames for that period. The median value of theeach pixel is used to create the ﬂat frames. Pixels which could be aﬀected by bright objects
and vignetting by the auto-guider probe are ﬂagged out before processing the ﬂat frames.
Then, object frames are divided by the ﬂat frames. The peak-to-peak ratios of backgroundlevels over the entire ﬂat ﬁelded frames are less than 2%. We do not expect any signiﬁcantsystematic errors by the ﬂatﬁelding that aﬀect accuracies of the following analyses.
After the ﬂatﬁelding process is done, the distortion correction is applied to each frame
using the 5-th order polynomial transformation formula derived from Miyazaki et al. (2002).

–7–
The changes in shapes and sizes of objects in the frames due to the distortion are negligible
compared with the size of PSFs. Each pixel is transformed so that the surface brightness ofthe pixel is conserved, since the relative ﬂux in each pixel has been already corrected by the
ﬂatﬁelding process. Positional displacement of objects along the perpendicular orientation
due to the atmospheric eﬀect is also corrected at the same time.
The sky background is subtracted from each distortion-free frame, We divide an image
into temporary meshes with 64 × 64 pixels, corresponding to 12.9 arcsec in one direction.
The mode of counts in each mesh is adopted as the sky background levels after smoothed bya median ﬁlter with a 3 × 3 kernel. This global sky subtraction should not aﬀect photometricaccuracies of compact objects such as galactic stars and faint galaxies. However, caution
must be taken when we apply these values to extended objects such as nearby galaxies withlow surface brightness.
The stacking of the images is performed for each of the pointings separately in thefollowing way. First, we choose 100 to 200 unsaturated stellar objects with a peak ﬂux of> 500ADU in each frame distributing evenly over the entire frame. Second, we determinea reference frame of all the reduced frames. For the other frames, relative positions androtation (∆x, ∆y, ∆θ) and ﬂux ratios to the reference frame are determined based on a leastsquare method using these stellar objects commonly detected in the frames. Third, for all
the frames, the positional shifts and rotations are corrected. Here, we do not use frames
which have the ﬂux ratios of less than 30% of that of the reference frame. Finally, all theframes are co-added by calculating ﬂux-weighted average values in each pixel with a 3-σclipping process. The clipping process eﬀectively removes cosmic rays and bad or saturated
pixels, and satellite trails etc., and should not aﬀect the resulting total ﬂuxes of objects
because the diﬀerence in counts among frames with various seeing sizes is mostly within the3-σ threshold.
After the co-adding of each band images, all the ﬁve-band (B, V, Rc , i , z  ) images are registered into identical positional coordinates as we intend to execute multi-waveband aperture
photometry. The positional registration is performed by the geometrical transformation by
3-rd order polynomials ﬁtting using the positions of stellar objects common to the co-added
stacked images. The r.m.s. deviation of residuals of transformed positions with respect tothe reference positions in the Rc -band images is approximately 0.1-0.15 arcsec.
For each pointing, the geometrically matched images are smoothed with Gaussian ﬁlters
iteratively so that the PSF sizes of all the images become the same as that in the worst image
used in this study. The PSF sizes of the resultant images are 0.80, 0.84, 0.82, 0.82, and 0.82
arcseconds in SXDS-C, SXDS-N, SXDS-S, SXDS-E, and SXDS-W, respectively.

–8–
Finally, regions with low S/Ns or those aﬀected by saturation or overﬂow of electrons
are carefully checked and assigned as ﬂagged regions which are indicated in photometriccatalogs. The resultant PSF sizes and total usable areas of each pointing is summarized inTable 3.

4.

MULTI-WAVEBAND PHOTOMETRIC CATALOGS

We construct multi-waveband photometric catalogs of the objects identiﬁed in the SXDF.

4.1.

Object Detection

Object detection and photometry are performed using SExtractor version 2.3.2 (Bertin& Arnouts 1996). Before performing the object detection and photometry, the outskirtsof the stacked images are trimmed in order to eliminate any failure such as inappropriate
estimation of sky levels and memory overﬂow during extraction of objects by SExtractor due
to a high noise level. After trimming the outskirts, each individual image has similar size of
approximately 10, 800 × 8, 100 pixels. We use the detection criteria that at least ﬁve pixels
which have values of > 2σ of sky background r.m.s. noise must be connected in order to be
included in our object list. Positions of detected objects are translated to sky coordinates
based on the astrometric calibration which will be discussed in Section 5.2.

4.2.

Aperture Photometry

After the object detection is completed in one of the ﬁve bands, which we call ‘detection
band’, multi-waveband aperture photometry is carried out with SExtractor for the other fourbands using the same apertures at the same positions as those in the detection band. We
perform aperture photometry with ﬁxed circular apertures (MAG APER) of 2 and 3 arcseconds
in diameter, and variable elliptical apertures (MAG AUTO and MAG BEST) utilizing the Kron’sﬁrst moment scheme with apertures’ radii of 2.5 rKron (Kron 1980). Measured ﬂuxes of objects
are converted into magnitude using the photometric zeropoints, which will be described in
Section 5.1. The extracted magnitude is considered to be the asymptotic total magnitudeof galaxies, which should sample at least 94% of ﬂux of a galaxy (Bertin & Arnouts 1996).
Other useful parameters for objects including shape, size, and stellarity and so on, are also
extracted by SExtractor at the same time. All the object parameters which are listed in the

–9–catalogs are summarized in Table 4.
In this procedure, B-, V -, Rc -, i -, and z  -selected multi-waveband catalogs are createdfor each pointing separately, which means that each object has measurements in the ﬁvebands (B, V, Rc , i , z  ) in each catalog. Thus 25 catalogs (= 5 pointings times 5 detection
bands) in total are created.

4.3.

Flagged Areas

Although we have ﬂagged out bad pixels from the image frame in the reduction procedure, there still remain (1) areas strongly aﬀected mainly by overﬂow of electrons or envelopesof extended objects, and (2) areas weakly aﬀected by large halos from very bright stars or
noisy regions. In our catalogs, the objects which are located in these areas are ﬂagged 1 and
2, respectively. Otherwise, the objects located in the clean area have a ﬂag value of 0.

The basic parameters used in the object detection and aperture photometry are summarized in Table 5. The complete conﬁguration ﬁles for SExtractor can be found at http://www.naoj.org/Sci
Note that we adopt default parameters of SExtractor for the source extraction unless we do
not ﬁnd any problem with the output values. The parameters adopted are suitable for extraction of moderate to faint compact galaxies, and not heavily optimized to extraction ofextended objects or very faint objects. Thus the reliability of the current catalogs for thoseobjects should be regarded with some caution.

5.
5.1.

CALIBRATIONSPhotometric Calibration

Photometric zeropoints in all the ﬁve bands are determined in the following way. First,
we compare our observation with the earlier calibrated observations for the same overlapping
ﬁeld, which were made by the Suprime-Cam instrument development team (Ouchi et al.
2001, 2004). The southern half of our center pointing (SXDS-C) is overlapped with thenorthern half of the image obtained by Ouchi et al. (2001). They determined the photometriczeropoints in the Suprime-Cam band system based on observations of photometric standard
stars SA92 and SA95 for B, V and Rc bands and a spectrophotometric standard star SA95-42
for i and z  bands.
We estimate the photometric zeropoints for our center pointing (SXDS-C) by comparing
magnitudes (MAG AUTO) of ∼ 200 objects which are commonly detected in the two images and– 10 –
have FWHMs of smaller than 1.2 arcsec with a magnitude range of 20.0 to 24.0. Then, we
extend these zeropoints to the surrounding four pointings by using objects in the overlapping
areas in the same manner as for the determination of the zeropoint for the SXDS-C. In ourcatalogs, we list the instrumental magnitude obtained in the Suprime-Cam band system.
Hence, no transformation of magnitude considering color terms of band responses are conducted. The color terms between the Suprime-Cam band system and the standard system
are basically small – for the reddest stars, of order 0.1 in the B, V, i , z  bands, and slightly
large (∼ 0.3) in the Rc band.
To test internal consistency of our zeropoint calibrations, we examine colors of stellar
objects in our photometric catalogs and those computed by convolving typical stellar SEDs
covering the complete ranges of spectral types (Gunn & Stryker 1983) with Suprime-Cam
response curves. We ﬁnd that it is necessary to apply corrections to the zeropoints of∆V = +0.03 and ∆z  = +0.03 for all the ﬁve pointings, and ∆i = +0.05 and +0.02 for theSouth and East pointings (i -S and i -E), respectively. After these corrections are applied,the colors of the stellar objects in our catalogs agree with those computed for the SEDs from
Gunn & Stryker (1983) within 0.03 mag in all the ﬁve bands (Figure 4). The systematic
errors adjusted here may be due to relatively small number of stellar objects used in thedetermination of the photometric zeropoints.
To ensure the accuracy of our zeropoints, we conduct CCD photometric observations ofa part of the SXDF in B, Rc , i and z  bands with the UH88-inch telescope at Mauna Kea
on October 9 and 10, 2002. We compare magnitudes of objects in B, Rc , i and z  bands
obtained by the UH88-inch observations with our catalog values. We ﬁnd that our catalog
values agree within 0.05 mag r.m.s. for B and Rc bands and within 0.1 mag r.m.s. for iand z  bands. The relatively large diﬀerence seen in i and z  bands are probably due tothe diﬀerence in the band responses between the UH88-inch system and the Suprime-Cam
system, and low S/Ns of shallow SXDS images used for comparison with the UH88-inch data.
To conﬁrm accuracy of the photometric zeropoints in the i and z  bands in our catalogs,
follow-up observations are made for the two bands with the 1.0-m telescope at the UnitedStates Naval Observatory (USNO). The observations at the USNO were performed on thephotometric night of December 3, 2004 in the course of photometric calibration for a high-z
supernovae search (Yasuda et al. 2007, in preparation). We ﬁnd that a systematic diﬀerence
between zeropoints in our catalogs and those determined by the USNO observation is as
small as 0.03 mag r.m.s.
Thus we conclude that the uncertainties of calibrated photometric zeropoints of the
SXDS images are 0.03 − 0.05 mag r.m.s. The adopted photometric zeropoints are listed inTable 2.

– 11 –
5.2.

Astrometric Calibration

Astrometric calibration is performed using ∼200 stars with magnitude of Jvega < 16.5
derived from 2-MASS point source catalogs (PSC; Cutri et al. 2003), in each of the ﬁvestacked Rc -band images. Since images in the other bands have been geometrically transformed, the positions of objects detected on the images should coincide with those found in
the Rc -band images (see Section 3). The world coordinates for the images are calculated
for the Rc -band images and written in the FITS headers of all the stacked images with the
IRAF tasks CCMAP and CCSETWCS. The r.m.s. uncertainties of the determined coordinates across the FOV with respect to the PSC positions are on the order of 0.2 arcsec inboth RA and Dec in each pointing (Table 6). The positional error of the PSC is reported as∼ 0.07 arcsec.
We examine possible systematic diﬀerences in the calculated world coordinates between
diﬀerent pointings by using objects in the overlapping areas of two pointings. Figure 5 andFigure 6 show the diﬀerences in RA and Dec of point-like objects with FWHMs < 1.2 arcsec
in the magnitude range R = 20.5 − 24.0 detected both in the SXDS-C image and othersurrounding images. From the ﬁgures, the world coordinates assigned to each of the ﬁve
pointing images are in good agreement with one another with accuracy of about 1 arcsec
at the edges of the images. The diﬀerences of the world coordinates between two pointingsstudied here are thought to be for the worst cases, since they are derived based on only
the objects in the overlapping areas which are located at the edge of each image. Slopes
and oﬀsets in the residuals are seen in each panel of the ﬁgures. These disagreements inthe positions of objects between two overlapping images come from possible tilts, oﬀsets
and/or diﬀerence in geometrical scale of the world coordinate systems between the two
images. Residuals of distortion in the images and the atmospheric eﬀect, which have notbeen completely removed in the reduction process, may cause the errors in determinationof the world coordinate systems. Nevertheless, the SXDS catalogs have a good enoughpositional accuracy to perform follow-up spectroscopy.

The astrometrically and photometrically calibrated images of the SXDS in ﬁve bands (B, V, Rc , i , z  ) described in this section are released via SXDS Web Page (http://www.naoj.org/Science/S
The raw image data are also available to the public via Subaru Telescope ARchives System(STARS: https://stars.naoj.org).

– 12 –
6.

PROPERTIES OF THE CATALOGS
6.1.

Galactic Extinction

All the magnitudes listed in the catalogs are not corrected for the Galactic extinction.
The magnitude attenuation in each band estimated for the central position (02h 18m 00s , −05◦ 00 00 ;J2000) based on Schlegel et al. (1998), assuming an extinction curve with RV = AV /E(B −V ) = 3.1 is as follows: AB = 0.091, AV = 0.070, ARc = 0.056, Ai = 0.044, and Az  = 0.031.
Table 7 summaries the magnitude attenuation in each band for all the pointings. The Galactic extinction above must be taken into account when studying extragalactic objects.

6.2.

Limiting Magnitudes

We examine 3-σ limiting magnitudes measured for the 2 -diameter apertures in all
ﬁve bands (mlim column in Table 2) in the following manner. First, counts in unit of
ADU fallen into the 2 -diameter apertures are measured at approximately 10,000 positions
randomly selected and spreading over the entire image. Next, a histogram of the counts
is produced and only the faint side of the histogram is ﬁt by the Gaussian function, as
the bright-side tail is composed of not only the sky background but also photons fromthe objects. The σ of the best-ﬁt Gaussian is regarded as the 1-σ sky ﬂuctuation of theimage for the 2 -diameter apertures. We ﬁnd that limiting magnitude in each band is
B = 28.4, V = 27.8, Rc = 27.7, i = 27.7 and z  = 26.6 (AB, 3 σ, φ = 2 ) in the deepest
images of the ﬁve pointings.

6.3.

S/N Distribution Map

To examine the homogeneity of S/Ns of the SXDS data across the entire ﬁeld, weinvestigate the S/N distribution map in the Rc -band images. First, we calculate the sky
ﬂuctuation per 2 -diameter aperture within a mesh with a size of 350 × 350 pixels at eachposition in the same manner as in estimation of the limiting magnitudes. The mesh is shiftedwith a step of 175 pixels to cover the entire ﬁeld of view. Then, the ﬂuctuations for the 2”diameter aperture are converted to S/Ns for objects with an Rc magnitude of 27.5 at eachposition.
Figure 7 shows the S/N distribution maps for 27.5-mag objects thus obtained for eachpointing. The contour lines superimposed on the ﬁgures represent S/N=3 positions. First,
we see that the brightest sources on the images signiﬁcantly aﬀect the S/Ns around the– 13 –
sources. These areas are ﬂagged in the catalogs. Second, each image shows almost an
axisymmetric pattern of the S/N distribution when we ignore the low S/N areas due to thebrightest sources. This pattern reﬂects the vignetting of the prime focus. Third, the pattern
is slightly distorted for each image, probably due to bright sources, vignetting by the autoguider probe, diﬀerence in the quantum eﬃciency among CCDs, and so on. However, thediﬀerence in S/Ns between in such distorted areas and in clean areas with higher S/Ns ison the order of 20% at the maximum. This diﬀerence is negligible in most studies even for
faint sources. Thus, we can securely utilize the data with quite homogeneous S/Ns acrossthe ﬁeld of view.

6.4.

Completeness and False Detection Rate

Completeness of the object detection as a function of magnitude for each band and foreach pointing is estimated. The detection completeness is determined based on a Monte
Carlo simulation by adding artiﬁcial objects which have Gaussian proﬁles with their Poisson
photon noise into a stacked image used to create the catalogs and then detect them again
in the same manner as in creating catalogs. The detection rate thus obtained is regarded
as the detection completeness. Here the artiﬁcial objects are added to random positionsin the eﬀective area of the image. The FWHMs of the Gaussian proﬁles are set to be thesame as the PSF sizes representative of the images which are derived from Table 3. The
artiﬁcial objects contaminated by any neighbor sources listed in the catalogs are removed in
calculation of the detection rate. If such blended objects are not removed from the sample,
the detection rate should slightly decrease. However, since the contamination by neighborobjects leads to misidentiﬁcation of the artiﬁcial objects in the object detection process, we
adopt the approach which excludes the blended artiﬁcial objects, which was also employedin Kashikawa et al. (2004).
The completeness determined by this procedure is shown in Figure 8 for all of the ﬁvebands. For the ﬁve pointings, the curves of the detection completeness drop similarly withincreasing magnitude. We can summarize that the detection completeness in each image is
at approximately 50% for objects with the 5-σ magnitude and 30% for objects with the 3-σmagnitude, with a slight diﬀerence depending on the band.
On the other hand, rates of the false object detection are also estimated as a function ofmagnitude. We generate negative images of each pointing and in each band by multiplying
all the counts of the stacked images by −1, and then perform object detection for the negative
images in the same manner as in creating the catalogs. In this process, detected objects are
considered to be spurious objects. Thus we deﬁne the false detection rate as the number of

– 14 –
the spurious objects divided by the total number of objects in the catalogs in each magnitudebin, i.e., Nspurious (m)/Ntotal (m) (Figure 9). It is seen that a contribution of the false detection
is negligible, which does not exceed 0.5% if any, in the magnitude range brighter than the3-σ limiting magnitude.

6.5.

Magnitude Diﬀerences among the Catalogs

In the areas where two images overlap with each other, a large fraction of objects aredetected and their magnitudes are measured in both of the pointings. We can check systematic diﬀerence in magnitude of those objects between the two catalogs. Since photometriczeropoints of the catalogs are determined so that MAG AUTO of objects in the overlapping
areas should coincide between any two catalogs (see Section 5.1), the magnitude diﬀerencein MAG AUTO is negligible. MAG BEST should be identical to MAG AUTO for isolated objects.
Therefore magnitude diﬀerences only for the 2 - and 3 -diameter aperture magnitudes areinvestigated in the following manner.
First, we choose only compact objects with FWHMs < 1.2 arcsec which are listed inboth the z  -selected catalog for SXDS-C and that for another pointing which overlap witheach other, then measure the diﬀerence in magnitude of the objects by subtracting magnitudein the SXDS-C catalog by that in the other. Then, we determine the oﬀset in the magnitude
diﬀerence, i.e., systematic magnitude diﬀerence, by ﬁtting the magnitude diﬀerence for arange of 21.0 to 23.5 mag by a least square method. Table 8 lists the systematic magnitude
diﬀerence thus obtained for the 2 - and 3 - diameter aperture magnitudes in each band foreach pair of catalogs. We note that 2 - and 3 -diameter aperture magnitudes indicate small
systematic diﬀerences of ≤ 0.05 mag between the catalogs. We think that this diﬀerence is
probably due to diﬀerence in the PSF shape of the stacked images among the pointings.
In order to see the eﬀects of the systematic diﬀerence in the aperture magnitudes, we plot
diﬀerences in colors of objects between the catalogs measured with the 2 -diameter aperturemagnitude (Figure 10). In this ﬁgure, for each of the four pairs of catalogs, the diﬀerencesin colors (B − V, V − Rc , Rc − i , i − z  ) are plotted. We see the systematic diﬀerence in the2 -diameter aperture colors of objects, which is no larger than 0.1 mag, due to the diﬀerencein the aperture magnitude between the catalogs. Since the magnitudes in the catalogs are
not adjusted for this diﬀerence, the magnitude diﬀerence between the catalogs should betaken into account in the case that the 2 - or 3 -diameter aperture magnitude is used. For
MAG AUTO and MAG BEST, which are considered as asymptotic total magnitude, the magnitude
diﬀerence is negligible.

– 15 –In the same manner as mentioned above, the r.m.s. of the magnitude diﬀerences foreach band are calculated as a function of MAG AUTO, which shows no systematic diﬀerence.
Here the detection band for catalogs are chosen to be the same as the band for which themagnitude diﬀerence is investigated. For instance, to obtain the magnitude diﬀerence in the
B band, the B-selected catalogs are used. Again, only the objects with FWHMs < 1.2 arcsec
are used for measurement of the magnitude diﬀerence. The r.m.s. of magnitude diﬀerences
measured in the above way are shown in Figure 11 for each band. It is seen that the r.m.s.
steadily increases with increasing magnitude. The r.m.s. of the magnitude diﬀerence can be√converted into random magnitude error by dividing it by 2. We ﬁnd that the magnitude
errors thus estimated are entirely consistent with those expected from our calculation of
the limiting magnitude discussed in Section 6.2. Note that the magnitude errors listed inour catalogs are computed by SExtractor assuming the simple Poisson statistics for the sky
background ﬂuctuations. The sky background noises measured in photometry apertures forthe stacked images are likely to be larger than those estimated by SExtractor. Therefore
the magnitude errors in the catalogs are underestimates and are not suitable for immediate
analysis such as SED ﬁtting and photo-z estimation.

6.6.

Number Counts of Galaxies

To examine the characteristics of the catalogs we compare the number counts of galaxies
with data from previous surveys.
We calculate the number counts of galaxies by correcting our raw counts of detected
objects against the detection completeness estimated in Section 6.4. The detailed studies
of density ﬂuctuations of galaxies based on the number counts or luminosity functions ofgalaxies can be found in elsewhere (e.g., Yamada et al. 2005). Star/galaxy separationis conducted based on FWHM and 2 -diameter ﬁxed aperture magnitude of the objects,
for which the eﬀect of the Galactic extinction (Table 7) is corrected. The criteria for the
separation is determined so that both objects which have sizes as same as PSF of the images
and saturated objects are eﬀectively removed. For fainter magnitudes (> 22.5 − 23 mag), no
separation is performed since a large fraction of galaxies have small sizes similar to those ofstars, and contribution to the number count by stars is small at the faint magnitudes. The
criteria for the separation adopted is summarized in Table 9.
The results of the number counts of galaxies are as follows: First, internal comparisonof the number counts among the ﬁve pointings of the SXDS are shown for each band in
Figure 12 to Figure 16. In each ﬁgure, raw number counts of the galaxy samples extracted
by the star/galaxy separation process above are plotted with open symbols. Eﬀective ar-– 16 –
eas used to calculate the number counts per unit area are given in Table 3. The number
counts of galaxies which are corrected for the detection completeness are also superimposedwith ﬁlled symbols. Here, the correction is performed by using the detection completeness as a function of magnitude which has been estimated for the Gaussian proﬁles, i.e.,
N(m)corrected = N(m)observed /completeness(m). Error bars are calculated on the assumptionof simple Poisson noise. We can say that no large systematic diﬀerence is seen in the superposition of the number counts of galaxies among the ﬁve pointings in B, V , and Rc bands down
to the 3-σ magnitudes. Similarly, in the other two bands, the number counts are consistent
with one another among the ﬁve pointings within the Poisson errors at magnitudes brighter
than ∼ 25 mag. However, diﬀerences in the number counts at fainter magnitudes are foundby a factor of 1.4 in the i band and 1.7 in the z  band. This might be partly due to theﬁeld-to-ﬁeld variation in the number counts of galaxies. However, since we have used only
the Gaussian proﬁles with FWHM∼ 0.8 arcsec in the simulation of the detection completeness, the completeness is likely to be overestimated at a faint magnitude range where mostobjects have extended proﬁles. So, we cannot conclude that the diﬀerences in the number
counts at the faint end seen in the i and z  bands indicates the ﬁeld-to-ﬁeld variation.
Next, the number counts in the SXDF are compared with results by previous surveys(Figure 17 to Figure 21). Here, we calculate mean number counts of galaxies corrected for thedetection completeness for the entire SXDF using the above sample for the ﬁve pointings, andplot them with ﬁlled circles with error bars. The error bars are again based on the Poissonnoise. The results in other surveys are derived from Kashikawa et al. (2004), Capak et al.
(2007), and Metcalfe (2007)2 and the references cited therein. A complete list of referencescompared in the ﬁgures is as follows: SDF – Kashikawa et al. 2004, HDFs/WHDFs –Metcalfe et al. 2001, Arnouts99 – Arnouts et al. 1999, SDSS – Yasuda et al. 2001, CADIS –Huang et al. 2001, Kummel&Wagner01 – Kümmel & Wagner 2001, Tyson88 – Tyson 1988,
Arnouts01 – Arnouts et al. 2001, Capak04 – Capak et al. 2004, Smail95 – Smail et al. 1995,
SH93 – Steidel & Hamilton 1993, Hogg97 – Hogg et al. 1997, COSMOS – Capak et al. 2007,
Postman98 – Postman et al. 1998, Lilly91 – Lilly, Cowie, & Gardner 1991.
From the comparison of the number counts among various surveys, we understand that
our sample is consistent with results in other surveys for blank ﬁelds for the following reasons.
First, the mean number counts of the SXDS show a good agreement with those from previous
surveys down to magnitude where the number counts start to turn oﬀ. Second, it reproduces
the faint-end data points derived from Hubble Deep Field (HDF). Again, the detection
completeness simulated with the Gaussian proﬁles might cause a slightly large uncertainty
2

http://star-www.dur.ac.uk/~nm/pubhtml/counts/counts.html

– 17 –of the corrected number counts at the faint end in each band. Thus, we conclude that the
SXDS catalogs in the present study have characteristics similar to those in previous surveys
and can be applied to general statistical studies on faint objects ranging from Galactic objects
to high-redshift galaxies.
We would like to mention how the Poisson error on the number counts is decreased byincrease of the survey area. Figure 22 shows a comparison of the Poisson error associatedwith the number counts of galaxies among three surveys which are based on diﬀerent survey
areas, namely the HDF-North (5.3 arcmin2 ), GOODS-North (170 arcmin2 ) and SuprimeCam 1 FOV (918 arcmin2 ). The uncertainty in the number counts due to the Poisson error
is shown with a pair of lines for each of the surveys. For reference, the number count of
SXDS is superimposed with a thick solid line. From the ﬁgure, we clearly see that the
Poisson error decreases with increasing survey area. In particular, at the bright magnitudesof 21 − 23 mag, a large uncertainty by a factor of > 3 in the number count is found for
the HDF-North area (dotted lines). For the same magnitude range, the uncertainty for aﬁeld of view of Suprime-Cam (dot-and-dash lines) is 8 − 15 times as small as that for the
HDF-North area and 1.5 − 2.5 times as small as that for the GOODS-North area. Even at
the faintest magnitudes of > 27 mag, a signiﬁcant diﬀerence in the uncertainty from one
to another survey area is seen. Thus it is emphasized that the SXDS data, which consistsof ﬁve pointings of Suprime-Cam ﬁeld of view, is a useful data set for studies on celestial
objects which dominate a large fraction of the number counts without suﬀering from the
Poisson error and the ﬁeld-to-ﬁeld variation as well, e.g., a study of evolution of the LSSs.

7.

SUMMARY

The optical imaging observations of the SXDS project are carried out using the SuprimeCam on Subaru Telescope in B, V, Rc , i and z  bands. The SXDF has a contiguous area
coverage of ∼1.2 deg2 , which consists of ﬁve Suprime-Cam pointings.
The photometric zeropoints are determined with an absolute accuracy of no larger than
0.05 mag r.m.s. in the photometry. The r.m.s. of the astrometric accuracies across the ﬁeld
is on the order of 0.2 arcsec in both RA and Dec. The systematic diﬀerences in the adopted
world coordinates between diﬀerent pointings are within about 1 arcsec at the outer edgeof the ﬁeld of view. Thus, our SXDS catalogs have a good enough positional accuracies toperform follow-up spectroscopy.
Multi-waveband photometric catalogs of detected objects are created for each band andfor each pointing. Each of the catalogs contains more than a hundred sixty thousand of

– 18 –
objects. The catalogs have quite homogeneous S/Ns across the ﬁeld. The achieved limiting
magnitudes in each band are B = 28.4, V = 27.8, Rc = 27.7, i = 27.7 and z  = 26.6
(AB, 3 σ, φ = 2 ).
The detection completeness as a function of magnitude is estimated by a Monte-Carlosimulation assuming a Gaussian proﬁle. The number counts of galaxies in the SXDF ineach band are computed by correcting the detection completeness, which are consistent witheach other among the ﬁve pointings with a slight diﬀerence at faint magnitudes. The mean
number counts of galaxies averaged over the ﬁve pointings show a good agreement withresults from previous surveys down to the faint-end magnitude.
With the aid of the wide coverage area of SXDS data, the uncertainty of the number
counts of galaxies due to the Poisson error is greatly decreased. It is emphasized that the
SXDS data is extremely useful for pursuing studies on celestial objects spreading in a wide
ﬁeld without suﬀering from the Poisson error and the ﬁeld-to-ﬁeld variation. The SXDS
catalogs can be applied to studies, scientiﬁc objectives of which range from the Galactic
objects to the large scale structures of the universe.
The optical data, the compiled photometric catalogs, and conﬁguration ﬁles used to create the catalogs have been released to the public and can be retrieved from the public dataarchives server of National Astronomical Observatory of Japan, http://www.naoj.org/Science/SubaruPr
Acknowledgments
We thank Dr. Masafumi Yagi for his kind suggestions in analyzing the data based on
the software package developed by himself. We thank Dr. Nobunari Kashikawa for fruitful discussions. The anonymous referee would also be appreciated for his/her constructive
suggestions. We would like to convey our gratitude to the Subaru Telescope builders andstaﬀ for their invaluable help to complete the observing run with Suprime-Cam. Data reduction/analysis in this work was in part carried out on ”sb” computer system operated by
the Astronomical Data Center (ADC) and Subaru Telescope of the National Astronomical
Observatory of Japan. Use of the UH 2.2-m telescope for the observations is supported by
NAOJ.

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– 21 –

N

-05:00:00
E(J2000)W

1.3 degree
02:18:00(J2000)Fig. 1.— Three color (B-, Rc - and i -band) composite optical image of the Subaru/XMMNewton Deep survey Field (SXDF). The area consists of ﬁve Suprime-Cam pointings, corresponding to a 1.22 square degree ﬁeld.

– 22 –

Fig. 2.— Band response curves of the SXDS Suprime-Cam imaging observations (upper
panel) and the typical spectra of an early-type galaxy (labeled as “elliptical”) and a latetype galaxy (labeled as “disk”) at the diﬀerent redshifts (lower panel).

– 23 –

Fig. 3.— Dithering patterns for the Rc -band observations for ﬁve pointings. The center(0, 0) position of each pointing corresponds the coordinates given in Table 1.

– 24 –

Fig. 4.— Color-color plots for the stellar objects in the SXDF (red dots) and Gunn-Stryker
stars (blue triangles). The SXDF colors agree with the colors of Gunn-Stryker stars to theaccuracy of 0.03 mag or better.

– 25 –

Fig. 5.— Residuals in the RA (ﬁlled circles) and in the Dec (crosses) of the World-Coordinate
values in the RA direction in 4 overlapping areas between the center pointing (SXDS-C) andthe north (SXDS-N: upper left panel), SXDS-C and the south (SXDS-S: upper right panel),
SXDS-C and the east (SXDS-E: lower left panel), and SXDS-C and the west (SXDS-W:
lower right panel). Only the objects with the FWHM< 1.2 arcsec are plotted.

– 26 –

Fig. 6.— Same as Figure 5, but the residuals in the Dec direction.

– 27 –

Fig. 7.— Contour maps of the sky S/N distribution in the Rc band. The deviation (in ADU)
of sky background ﬂuctuations per 2 -diameter aperture are measured and converted to the
S/Ns for objects with a magnitude of 27.5 at each position.

– 28 –

Fig. 8.— Object detection completeness of the ﬁve pointings as a function of magnitude ineach band.

– 29 –

Fig. 9.— False object detection rate for the ﬁve pointings as a function of magnitude in
the ﬁve bands. The perpendicular dashed lines indicate the 3-sigma limiting magnitude
measured in 2 -diameter aperture in the center pointing (SXDS-C).

– 30 –

Fig. 10.— Residuals in colors measured for the 2 -diameter aperture magnitudes of theobjects common in two separate pointing catalogs.

– 31 –

Fig. 11.— Random magnitude diﬀerence as a function of magnitude in each band. The
random magnitude diﬀerences are computed by comparing ﬂuxes of compact objects with
FWHMs < 1.2 arcsec between the central pointing and the four surrounding pointings.

– 32 –

Fig. 12.— B-band galaxy number counts for the ﬁve pointings. Open symbols are the raw
numbers extracted using the star/galaxy separation process. Filled symbols are the numbers
corrected for the detection completeness estimated with Gaussian artiﬁcial sources.

– 33 –

Fig. 13.— V -band galaxy number counts for the ﬁve pointings. Symbols are the same as in
Figure 12.

– 34 –

Fig. 14.— Rc -band galaxy number counts for the ﬁve pointings. Symbols are the same as
in Figure 12.

– 35 –

Fig. 15.— i -band galaxy number counts for the ﬁve pointings. Symbols are the same as in
Figure 12.

– 36 –

Fig. 16.— z  -band galaxy number counts for the ﬁve pointings. Symbols are the same as in
Figure 12.

– 37 –

Fig. 17.— Comparison of the galaxy number count of the SXDS (blue ﬁlled circles) withother major optical imaging surveys. The SXDS values are corrected for the detection
completeness and the mean value for the ﬁve pointings is plotted.

– 38 –

Fig. 18.— Same as Figure 17, but for V -band data.

– 39 –

Fig. 19.— Same as Figure 17, but for Rc -band data.

– 40 –

Fig. 20.— Same as Figure 17, but for i -band data.

– 41 –

Fig. 21.— Same as Figure 17, but for z  -band data.

– 42 –

Fig. 22.— B-band number count uncertainties of galaxies, due to the Poisson error, for the
Hubble Deep Field (HDF), GOODS North ﬁeld (GOODS-N), a single Suprime-Cam ﬁeld
(SPCAM) are plotted with the SXDS number count (solid line). The uncertainty for eachsurvey is represented by a pair of lines plotted in the ﬁgure.

– 43 –

Table 1. COORDINATES OF SXDS FIVE POINTINGS
Pointing

RA (J2000)

Dec (J2000)

Position Angle of FOV

Center
North
South
East
West

02h 18m 00s .000
02h 18m 00s .000
02h 18m 00s .000
02h 19m 47s .070
02h 16m 12s .930

−05◦ 00 00. 00−04◦ 35 00. 00−05◦ 25 00. 00−05◦ 00 00. 00−05◦ 00 00. 00

North is up
North is up
North is up
East is up
West is up

Note. — Field names for each pointing and their positional coordinates. PA, the orientation of the FOV on the sky are chosen so that the overhead time to rotate the position angle
are minimized and eﬀective area with high signal-to-ratios are is maximized.

Exp. Time
(min)
345
330
330
330
330

319

313

321

291

293

248
232
232
232

Band-Pointing

B-C

B-N

B-S

B-E

B-W

V -C

V -N

V -S

V -E

V -W

Rc -C

Rc -N

Rc -S

Rc -E

0.76

0.74

0.78

0.76

0.72

0.76

0.82

0.80

0.72

0.78

0.82

0.82

0.84

0.80

PSF Size(arcsec)

34.25934.21934.27634.31533.64933.63933.64333.64833.63934.71634.69834.70634.70134.723Zeropts for Images(magAB ADU−1 )

27.51

27.67

27.74

27.57

27.72

27.77

27.75

27.65

27.78

28.21

28.06

28.33

28.39

28.09

mlim
(magAB )173,586

178,416

174,365

188,079

205,915

192,951

186,545

203,299

213,851

197,770

179,478

189,916

176,372

197,317

Nobj

Date of Observations

2002, Sep. 30, Oct. 01/07
2003, Nov. 17
2002, Nov. 01/09/30, Dec. 06
2003, Nov. 17
2002, Nov. 01/09/30, Dec. 06
2003, Nov. 17
2002, Nov. 02/09/30, Dec. 06

2003, Oct. 02/24/26
2004, Jan. 16, Oct. 09, Dec. 10/14/15
2005, Jan. 05
2003, Oct. 02/26, Nov. 17
2004, Jan. 16, Dec. 10/14/15
2005, Jan. 05, Sep. 28
2003, Oct. 02/24/26
2004, Jan. 16, Dec. 10/14/15
2005, Jan. 05, Sep. 28
2003, Oct. 02/26
2004, Jan. 16, Dec. 10/14/15
2005, Jan. 05, Sep. 28
2003, Oct. 02/26, Nov. 17
2004, Jan. 16, Dec. 10/14/15
2005, Jan. 05, Sep. 28

2002, Sep. 29/30, Oct. 01
2003, Nov. 17
2002, Nov. 02/05/27, Dec. 01/07
2003, Nov. 17
2002, Nov. 02/04/05/27, Dec. 01/07
2003, Nov. 17
2002, Nov. 04/05, Dec. 01/07
2003, Nov. 17
2002, Nov. 04/05/27, Dec. 01/07
2003, Nov. 17

Table 2. SUMMARY OF OBSERVATIONS AND DATA OF THE SXDS

– 44 –

232

647
440
309
368
598

217
252
184
267
311

Rc -W

i -C

i -N

i -S

i -E

i -W

z  -Cz  -Nz  -Sz  -Ez  -W
0.74

0.74

0.76

0.70
0.74

0.82

0.74

0.68

0.76

0.78

0.82

PSF Size(arcsec)

32.77632.74332.25833.07632.27834.08733.98634.04634.04234.05534.247

Zeropts for Images(magAB ADU−1 )

26.60

26.49

26.39

26.57
26.64

27.58

27.49

27.47

27.66

27.62

27.53

mlim
(magAB )167,748

160,415

167,779

183,324
163,324

178,543

175,404

187,791

176,394

181,352

186,648Nobj

2002, Sep. 29/30, Oct. 01, Nov. 04/05
2002, Nov. 04/05/10
2003, Sep. 21/22, Oct. 01/02
2002, Nov. 04/05/10
2003, Sep. 21/22, Oct. 01
2002, Nov. 04/05/10
2003, Sep. 21/22, Oct. 01/02
2002, Nov. 04/05/10
2003, Sep. 21/22, Oct. 01/02

2002, Sep. 29/30, Nov. 01/02/05/09/27/29, Dec. 06/07
2003, Oct. 20/21
2002, Sep. 29/30, Nov. 01/02/09/29, Dec. 07
2003, Sep. 22, Oct. 02/21
2002, Sep. 29/30, Nov. 01/02/09/29
2003, Sep. 22, Oct. 02
2002, Sep. 29/30, Nov. 01/02/09/29, Dec. 07
2003, Sep. 22, Oct. 02/21
2002, Sep. 29/30, Nov. 01/02/05/09/27/29, Dec. 06/07
2003, Oct. 20/21

2003, Nov. 17
2002, Nov. 02/09/30, Dec. 06
2003, Nov. 17

Date of Observations

Note. — Explanations of columns: (column 1) Filter pass-band (B, V, Rc, i , or z  ) and the pointed ﬁeld (center = C, north= N, south =S, east = E, and west =W). (2) Total exposure time in minutes. (3) The PSF size of the stacked image. (4)Photometric zeropoint. (5) A 3-sigma limiting magnitude in AB magnitude, measured by random sampling with 2 -diameter
aperture. (6) Number of objects detected in the image. (7) Date of observations.

Exp. Time(min)

Band-Pointing

Table 2—Continued– 45 –

– 46 –

Table 3. IMAGE SIZE AND THE AREA COVERED BY THE SXDS STACKED
IMAGES

Pointing

PSF Size(arcsec)

Area Covered(arcmin2 )Eﬀective Area(arcmin2 )Center (C)North (N)South (S)East (E)West (W)
Total

0.80
0.84
0.82
0.82
0.82
...

1041.91011.71008.4986.4
1017.74404.4

968.4
951.8
967.3
917.3
913.9
4057.0Note. — Properties of stacked images of the SXDS. The table lists pointing ﬁeld names,
ﬁnal PSF sizes, total areas covered by the stacked images, and eﬀective usable areas after
excluding the masked regions in each pointing. Five images in diﬀerent bands in eachpointing have the same PSF size. The total and eﬀective areas covered by the 5 pointings
taking account of overlapping areas between each pointing are 4404.4 arcmin2 and 4057.0
arcmin2 , respectively.

– 47 –

Table 4. LIST OF PARAMETERS IN THE SXDS OPTICAL IMAGING OBJECTCATALOGS
Parameter Name

Unit

Description

ID
X
Y
RA
Dec
KRON RADIUS
A IMAGEB IMAGEELLIPTICITY
THETA IMAGE
ISOAREA IMAGE
IsophotalMag(B)IsophotalMag(V)IsophotalMag(R)IsophotalMag(i)IsophotalMag(z)IsophotCorMag(B)IsophotCorMag(V)IsophotCorMag(R)IsophotCorMag(i)IsophotCorMag(z)2.0ApertureMag(B)2.0ApertureMag(V)2.0ApertureMag(R)2.0ApertureMag(i)2.0ApertureMag(z)3.0ApertureMag(B)3.0ApertureMag(V)3.0ApertureMag(R)3.0ApertureMag(i)3.0ApertureMag(z)sequential unique ID of the object in the catalog
pixels X coordinate of detection
pixels Y coordinate of detection
h:m:s right ascension of detection
d:m:s declination of detection
pixels Kron radius rK deﬁned in Bertin & Arnouts 1996
pixels 2nd order moment along the major axis
pixels 2nd order moment along the minor axis
isophotal weighted ellipticity of the object (=A IMAGE/B IMAGE)
degree position angle of the major axis, counter-clockwised, 0.0=X axis
pixels area of the lowest isophotemag isophotal magnitude in B
mag isophotal magnitude in V
mag isophotal magnitude in Rc
mag isophotal magnitude in i’
mag isophotal magnitude in z’
mag corrected isophotal magnitude in B deﬁned by Bertin & Arnouts 1996
mag corrected isophotal magnitude in V deﬁned by Bertin & Arnouts 1996
mag corrected isophotal magnitude in Rc deﬁned by Bertin & Arnouts 1996
mag corrected isophotal magnitude in i’ deﬁned by Bertin & Arnouts 1996
mag corrected isophotal magnitude in z’ deﬁned by Bertin & Arnouts 1996
mag φ 2 ﬁxed-aperture magnitude in B
mag φ 2 ﬁxed-aperture magnitude in V
mag φ 2 ﬁxed-aperture magnitude in Rc
mag φ 2 ﬁxed-aperture magnitude in i’mag φ 2 ﬁxed-aperture magnitude in z’
mag φ 3 ﬁxed-aperture magnitude in B
mag φ 3 ﬁxed-aperture magnitude in V
mag φ 3 ﬁxed-aperture magnitude in Rc
mag φ 3 ﬁxed-aperture magnitude in i’mag φ 3 ﬁxed-aperture magnitude in z’– 48 –

Table 4—Continued
Parameter Name

Unit

Description

MAG AUTO(B)MAG AUTO(V)MAG AUTO(R)MAG AUTO(i)MAG AUTO(z)MAG BEST(B)MAG BEST(V)MAG BEST(R)MAG BEST(i)MAG BEST(z)Err(IsophotalMag)(B)
Err(IsophotalMag)(V)
Err(IsophotalMag)(R)
Err(IsophotalMag)(i)Err(IsophotalMag)(z)Err(IsophotCorMag)(B)
Err(IsophotCorMag)(V)
Err(IsophotCorMag)(R)
Err(IsophotCorMag)(i)Err(IsophotCorMag)(z)Err(2.0ApertureMag)(B)
Err(2.0ApertureMag)(V)
Err(2.0ApertureMag)(R)
Err(2.0ApertureMag)(i)Err(2.0ApertureMag)(z)Err(3.0ApertureMag)(B)
Err(3.0ApertureMag)(V)
Err(3.0ApertureMag)(R)
Err(3.0ApertureMag)(i)Err(3.0ApertureMag)(z)Err(MAG AUTO)(B)

mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag
mag

automatic-aperture magnitude in B
automatic-aperture magnitude in V
automatic-aperture magnitude in Rc
automatic-aperture magnitude in i’automatic-aperture magnitude in z’
MAG AUTO(B) if there are no neighbors, or IsophotCorMag(B) other
MAG AUTO(V) if there are no neighbors, or IsophotCorMag(V) othe
MAG AUTO(Rc) if there are no neighbors, or IsophotCorMag(R) othe
MAG AUTO(i’) if there are no neighbors, or IsophotCorMag(i) otherw
MAG AUTO(z’) if there are no neighbors, or IsophotCorMag(z) otherisophotal magnitude rms error in B
isophotal magnitude rms error in V
isophotal magnitude rms error in Rc
isophotal magnitude rms error in i’isophotal magnitude rms error in z’corrected isophotal magnitude rms error in B
corrected isophotal magnitude rms error in V
corrected isophotal magnitude rms error in Rc
corrected isophotal magnitude rms error in i’corrected isophotal magnitude rms error in z’
φ2 ﬁxed-aperture magnitude rms error in B
φ2 ﬁxed-aperture magnitude rms error in V
φ2 ﬁxed-aperture magnitude rms error in Rc
φ2 ﬁxed-aperture magnitude rms error in i’
φ2 ﬁxed-aperture magnitude rms error in z’
φ3 ﬁxed-aperture magnitude rms error in B
φ3 ﬁxed-aperture magnitude rms error in V
φ3 ﬁxed-aperture magnitude rms error in Rc
φ3 ﬁxed-aperture magnitude rms error in i’
φ3 ﬁxed-aperture magnitude rms error in z’
automatic-aperture magnitude rms error in B

– 49 –

Table 4—Continued
Parameter Name
Err(MAG AUTO)(V)
Err(MAG AUTO)(R)
Err(MAG AUTO)(i)Err(MAG AUTO)(z)Err(MAG BEST)(B)
Err(MAG BEST)(V)
Err(MAG BEST)(R)
Err(MAG BEST)(i)Err(MAG BEST)(z)FLUX MAX(B)FLUX MAX(V)FLUX MAX(R)FLUX MAX(i)FLUX MAX(z)FWHM IMAGE(B)
FWHM IMAGE(V)FWHM IMAGE(R)FWHM IMAGE(i)FWHM IMAGE(z)CLASS STAR(B)
CLASS STAR(V)
CLASS STAR(R)
CLASS STAR(i)CLASS STAR(z)FLAGS(B)FLAGS(V)FLAGS(R)FLAGS(i)FLAGS(z)FLAG AREA

Unit

Description

mag automatic-aperture magnitude rms error in V
mag automatic-aperture magnitude rms error in Rc
mag automatic-aperture magnitude rms error in i’mag automatic-aperture magnitude rms error in z’
mag MAG BEST magnitude rms error in B
mag MAG BEST magnitude rms error in V
mag MAG BEST magnitude rms error in Rc
mag MAG BEST magnitude rms error in i’mag MAG BEST magnitude rms error in z’
ADU peak surface brightness in B
ADU peak surface brightness in V
ADU peak surface brightness in Rc
ADU peak surface brightness in i’
ADU peak surface brightness in z’
pixels FWHM of proﬁle from a gaussian ﬁt in B
pixels FWHM of proﬁle from a gaussian ﬁt in V
pixels FWHM of proﬁle from a gaussian ﬁt in Rc
pixels FWHM of proﬁle from a gaussian ﬁt in i’pixels FWHM of proﬁle from a gaussian ﬁt in z’
stellarity index in B: 0.0=Galaxy and 1.0=Starstellarity index in V: 0.0=Galaxy and 1.0=Starstellarity index in R: 0.0=Galaxy and 1.0=Starstellarity index in i’: 0.0=Galaxy and 1.0=Starstellarity index in z’: 0.0=Galaxy and 1.0=Starextraction ﬂags in B deﬁned by SExtractor
extraction ﬂags in V deﬁned by SExtractor
extraction ﬂags in Rc deﬁned by SExtractor
extraction ﬂags in i’ deﬁned by SExtractor
extraction ﬂags in z’ deﬁned by SExtractor
index for aﬀected area: 0=Clean, 1=STRONG HALO, 2=WEAK HALO

– 50 –

Note. — The list of ﬁelds listed in the optical multi-waveband catalogs. All the ﬁelds arecommonly included in all the catalogs created for each detection and each pointing.

– 51 –

Table 5. THE SEXTRACTOR PARAMETERS USED TO CREATE THE SXDSCATALOGS
Parameter

Value

DETECT MINAREA
DETECT THRESHANALYSIS THRESHFILTER
DEBLEND NTHRESH
DEBLEND MINCONTBACK SIZEBACK FILTERSIZEBACKPHOTO TYPEBACKPHOTO THICKSATUR LEVEL
GAIN

5
2.0
2.0
N
32
0.005
32
3
GLOBAL
24
50000.0
2.6

Note. — Parameters for SExtractor used to create catalogs. The above values are common
to all the SExtractor runs regardless of detection bandpasses.

– 52 –

Table 6. RESULT OF ASTROMETRIC CALIBRATION

Pointing

rms (RA)(arcsec)rms (Dec)(arcsec)

Number of Stars

Center
North
South
East
West

0.160
0.220
0.251
0.261
0.217

0.151
0.218
0.207
0.225
0.201

233
208
194
199
218

Note. — The resultant r.m.s. of astrometric calibration. The number of stars used for
calculation of the world coordinates is shown in the far right column.

– 53 –

Table 7. THE GALACTIC EXTINCTION IN THE SXDF
Pointing

E(B − V )

AB

AV

ARc

Ai

Az 

SXDS-CSXDS-NSXDS-SSXDS-ESXDS-W

0.021
0.019
0.023
0.021
0.021

0.091
0.083
0.100
0.091
0.091

0.070
0.064
0.077
0.070
0.070

0.056
0.052
0.062
0.056
0.057

0.044
0.040
0.048
0.044
0.044

0.031
0.029
0.034
0.031
0.031

Note. — The magnitude attenuation due to the Galactic extinction in each band calculated
for the central coordinates of each pointing based on Schlegel et al. (1998). Since magnitudes
in the released catalogs are not corrected for the Galactic extinction, the above values must
be taken into account when studying extragalactic objects.

– 54 –

Table 8. THE SYSTEMATIC DIFFERENCES IN APERTURE MAGNITUDESBETWEEN THE DIFFERENT CATALOGS OF THE SXDS
Pointing

B(2 )

B(3 )

V (2 )V (3 )Rc (2 )Rc (3 )i (2 )i (3 )z  (2 )z  (3 )

C-N
C-S
C-E
C-W

-0.04
-0.05
-0.02
+0.01

-0.02
-0.02
-0.01
+0.01

-0.04
-0.02
-0.02
-0.03

-0.01
-0.00
-0.01
-0.01

-0.03
-0.03
-0.01
+0.01

-0.01
0.00
-0.00
+0.00

+0.03
+0.03
+0.05
-0.01

+0.02
+0.00
+0.01
-0.00

-0.02
+0.00
-0.01
-0.01

-0.01
+0.01
-0.00
-0.01

Note. — The systematic diﬀerences in aperture magnitudes of objects between the catalog of SXDS-C and that of another surrounding pointing. For each band, the magnitude
diﬀerences in 2 - and 3 -diameter aperture magnitude are listed. The column of pointing
denotes the pair of pointings compared. For instance, the line for C-N indicates magnitude
diﬀerences calculated by magnitudes in SXDS-C minus those in SXDS-N.

– 55 –

Table 9. CRITERIA USED FOR STAR/GALAXY SEPARATION
Band

Criteria

B
V
Rc
i
z

(m = 19.5 − 23.0, σ > 4.5, σ > −2.5 (m − 20.3) + 4.5) or (m > 23.0)(m = 19.5 − 23.0, σ > 4.5, σ > −2.5 (m − 20.0) + 4.5) or (m > 23.0)(m = 19.5 − 22.5, σ > 4.5, σ > −2.5 (m − 20.3) + 4.5) or (m > 22.5)(m = 19.5 − 22.5, σ > 4.5, σ > −2.5 (m − 20.1) + 4.5) or (m > 22.5)(m = 18.5 − 22.5, σ > 4.5, σ > −3.0 (m − 19.3) + 4.5) or (m > 22.5)

Note. — The criteria used in star/galaxy separation procedure. Objects satisfying thesecriteria are recognized as galaxies. In the criteria list, m denotes 2”-diameter ﬁxed aperturemagnitude and σ is FWHM of objects in the corresponding band.