Äîêóìåíò âçÿò èç êýøà ïîèñêîâîé ìàøèíû. Àäðåñ îðèãèíàëüíîãî äîêóìåíòà : http://hea-www.harvard.edu/~pgreen/Papers/Q2345.ps Äàòà èçìåíåíèÿ: Wed Sep 14 22:07:19 2005 Äàòà èíäåêñèðîâàíèÿ: Mon Oct 1 23:58:24 2012 Êîäèðîâêà: Ïîèñêîâûå ñëîâà: ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï ï

CHANDRA OBSERVATIONS OF THE QSO PAIR Q2345+007: BINARY OR MASSIVE DARK LENS?
Paul J. Green, 1 Chris Kochanek, Aneta Siemiginowska, Dong­Woo Kim, Maxim Markevitch,
John Silverman, 1 and Anil Dosaj 1
Harvard­Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138; pgreen@cfa.harvard.edu
Buell T. Jannuzi 1
National Optical Astronomy Observatory, P.O. Box 26732, Tucson, AZ 85726­6732
and
Chris Smith 1
Cerro Tololo Inter­American Observatory, National Optical Astronomy Observatory, Casilla 603, La Serena, Chile
Received 2001 October 29; accepted 2002 February 4
ABSTRACT
The components of the wide (7>3) separation quasar pair Q2345+007A,B (z ¼ 2:15) have the most strik­
ingly similar optical spectra seen to date (Steidel & Sargent) yet no detected lensing mass, making this system
the best candidate known for a massive (#10 14 M # ) dark matter lens system. Here we present results from a
65 ks Chandra observation designed to investigate whether it is a binary quasar or a gravitational lens. We
find no X­ray evidence for a lensing cluster to a (0.5--2 keV) flux limit of 2 # 10 #15 ergs cm #2 s #1 , which is con­
sistent with lensing only for a reduced baryon fraction. Using the Chandra X­ray observations of the quasars
themselves, together with new and published optical measurements, we use the observed emission properties
of the quasars for further tests between the lens and binary hypotheses. Assuming similar line­of­sight
absorption to the images, we find that their X­ray continuum slopes are inconsistent (# A ¼ 2:30 ×0:36
#0:30 and
# B ¼ 0:83 ×0:49
#0:44 ) as are their X­ray--to--optical flux ratios. The probability that B su#ers intrinsic absorption
su#cient to account for these spectral di#erences is negligible. We present new optical evidence that the flux
ratio of the pair is variable, so the time delay in a lens scenario could cause some of the discrepancies. How­
ever, adequately large variations in overall spectral energy distribution are rare in individual QSOs. All new
evidence here weighs strongly toward the binary interpretation. Q2345+007 thus may represent the highest
redshift example known of an interaction­triggered but as­yet unmerged luminous active galactic nucleus.
Subject headings: gravitational lensing --- quasars: absorption lines --- quasars: individual (Q2345+007) ---
X­rays: general --- X­rays: individual (Q2345+007)
1. INTRODUCTION
The density and evolution of massive halos (M > 10 14
M # ) is a powerful probe of the cosmological model. The
number of massive clusters depends exponentially on the
amplitude of the power spectrum when normalized by # 8 ,
the amplitude on a scale of 8 h #1 Mpc. These halos can be
detected through an overdensity of galaxies (e.g., Postman
et al. 1996 for a modern example), from X­ray emission
(e.g., Ebeling et al. 2000), or with the Sunyaev­Zeldovich
e#ect (e.g., Holder et al. 2000). All these methods depend on
emission or absorption from the baryons in the halo. Alter­
natively, we can detect such massive halos using gravita­
tional lensing. This can be done either with weak lensing
surveys (Wittman et al. 2001) or with surveys for multiply­
imaged background sources with wide image separations
(usually defined by D# > 3 00 ; Kochanek, Falco, & Munoz
1999, hereafter KFM99).
Unlike the other methods of finding clusters, gravita­
tional lensing can also find `` dark '' clusters where there is
mass but no detectable baryons. Some candidates have been
found in weak lensing (shear­selected) surveys (Umetsu &
Futamase 2000; Erben et al. 2000), but only one such cluster
has been spectroscopically confirmed to date (Wittman et
al. 2001). Most of the debate about dark halos has focused
on the population of wide­separation quasar pairs
(WSQPs). All the lens candidates with separations D# < 3 00
have identifiable primary lens galaxies in deep NICMOS
observations (see synopses in KFM99 or Mortlock & Web­
ster 2000). Above 3 00 there are 17 objects. Four are bona fide
gravitational lenses, with similar optical/radio flux ratios
and negligible spectroscopic di#erences, as well as a plausi­
ble lensing galaxy. Eight more pairs, with discrepant radio/
optical flux ratios for a lens, or greater than 3 # velocity dif­
ferences, are very probably binary quasars. The remaining
five objects are WSQPs with similar spectra, essentially
identical redshifts, and no visible lens galaxy or cluster.
Were four pairs found to be dark clusters, they would imply
that the dark clusters were just as numerous as normal
clusters.
The problem with simply interpreting the WSQPs as dark
clusters is that we expect to find WSQPs even in the absence
of dark lensing as examples of binary quasars in which the
central engines of two nearby galaxies are simultaneously
active. We know that many of the pairs are binary quasars
because they have discrepant flux ratios as a function of
wavelength or greater than 3 # velocity di#erences. We can
prove from statistical analyses of the relative numbers of
1 Visiting Astronomer, Kitt Peak National Observatory and/or Cerro
Tololo Inter­American Observatory, National Optical Astronomy Observ­
atory, operated by the Association of Universities for Research in Astron­
omy, Inc. (AURA), under cooperative agreement with the National
Science Foundation; pgreen@cfa.harvard.edu.
The Astrophysical Journal, 571:721--732, 2002 June 1
# 2002. The American Astronomical Society. All rights reserved. Printed in U.S.A.
721

radio­quiet/radio­quiet, radio­quiet/radio­loud, and radio­
loud/radio­loud pairs that most of the remaining WSQPs
are binary quasars (KFM99). Since the number of WSQPs
is 100 times higher than expected from simple extrapola­
tions of the quasar­quasar correlation function, galaxy
interactions are crucial to creating binary quasars (see
KFM99; Mortlock & Webster 2000) and can be used as a
tool to study the triggering of nuclear activity in galaxies
(e.g., Osterbrock 1993).
Nonetheless, we still have puzzling examples of WSQPs
whose members have startlingly similar optical spectro­
scopic properties but no evidence for a lens. One example is
the optical quasar pair Q2345+007A,B (Tyson et al. 1986).
With a separation of 7>3, Q2345+007 is the most prominent
`` dark matter '' gravitational lens candidate. Optical spectra
of the two image components show exceptional intrinsic
similarity in both line profile and velocity (e.g., z ¼ 2:15,
Dv B#A ¼ 15 # 20 km s #1 ; Steidel & Sargent 1991, hereafter
SS91). Slight di#erences may be explainable as the com­
bined e#ect of time variability in the source and time delays
produced by the lens (Small, Sargent, & Steidel 1997). After
nearly 20 yr of study including deep optical and infrared
imaging (Nieto et al. 1988; Weir & Djorgovski 1991; Bonnet
et al. 1993; Gopal­Krishna et al. 1993; McLeod, Rieke, &
Weedman 1994; Fischer et al. 1994; Pello et al. 1996) and
VLA radio imaging (Patnaik, Schneider, Narayan 1996),
this pair remains the most obstinate mystery. Either mas­
sive, concentrated dark matter near the line of sight is acting
as a gravitational lens or two neighboring quasars have vir­
tually identical spectral characteristics yet significantly dif­
ferent luminosities.
In this paper we present a deep X­ray image of
Q2345+007 obtained with the Chandra X­Ray Observatory
(Chandra). Using the superb sensitivity and resolution of
Chandra we address two questions. First, we search the field
for extended X­ray emission from hot baryons in the lens,
which allows us to detect an optically dark halo. We cannot
detect a baryon­free halo, since the X­ray emission depends
on the hot baryons in the halo, but we can detect halos with
reduced baryon fractions. Second, we test whether the X­
ray flux ratios and spectra of the images are consistent with
the lens hypothesis. We describe the Chandra observations
next in x 2 and supporting Kitt Peak optical imaging in x 3.
We discuss the search for extended emission from a lensing
cluster in x 4, and contrast the properties of the two quasars
in x 5, summarizing in x 6.
2. Chandra OBSERVATIONS AND ANALYSIS
Q2345+007 was first observed by Chandra on 2000 May
26 for 25.2 ks using the back­illuminated S3 chip of the
Advanced CCD Imaging Spectrometer (ACIS) in faint,
timed­event mode. We o#set the pointing by 1 in spacecraft
Y­coordinate from the Chandra aim point to allow for imag­
ing of any extended emission all within one ACIS node. It
was observed again on 2000 June 27 in the same configura­
tion for 52.6 ks. However, because of an incorrectly config­
ured bias­only run in the previous ACIS segment, event
rates were about 20 times as high as expected from the chips
(I1 and I3) processed by Front End Processors (FEP) 0 and
3. The resultant telemetary situation created up to 50% dead
time in these chips. Telemetry dropouts are fully accounted
for in the final exposure times, but the aim­point chip S3
was in any case not a#ected.
We have used data reprocessed (2001 in April) at CXC. 2
We then ran XPIPE (Kim et al. 2000), which was specifically
designed for the Chandra Multiwavelength Project
(ChaMP; Green, Forster, & Kuraszkewiecz 2000). Data
screening applied in the CXC level 2 processing excludes
events with bad grades (mostly cosmic­ray events) and
events with status bits set such as bad pixels and columns.
Additional bad pixels and columns were excluded by exam­
ining events in the chip coordinates in each chip. To remove
time intervals of high background rates, we make a light
curve and exclude those time intervals with 3 # or greater
fluctuation above the mean background count rate. Given
di#erent characteristics between BI and FI chips, this is
done separately for each ACIS chip. The remaining expo­
sure time for the back­illuminated S3 chip is 64,998 s.
For source detection, we have applied the wavdetect algo­
rithm (Freeman et al. 2002), available in the Chandra Inter­
active Analysis of Observations (CIAO) software package. 3
Wavdetect is more reliable than the traditional celldetect
algorithm for finding individual sources in crowded fields
and for identifying extended sources, but the algorithm
tends to detect spurious sources near the detector edge. To
avoid such false detections, we generate an exposure map
for each chip (assuming a monoenergetic distribution of
source photons at 1.5 keV) and apply an exposure threshold
to be 10%. After performing various tests to find the most
e#cient parameters (see Kim et al. 2000), we select a signifi­
cance threshold parameter to be 10 #6 , corresponding to one
spurious source per CCD chip. We use scale parameters of
1, 2, 4, 8, 16, 32, and 64 pixels to cover a wide range of
source scales, thus accommodating the point­spread func­
tion variation as a function of o#­axis angle. For other
parameters, we have used default values given in CIAO2.1.
As shown in Figure 1, Q2345+007A,B are well­resolved
and strongly detected in the image. The Chandra astrometry
for the pair corresponds closely (#0>3) to the optical coun­
terpart positions, as expected from Chandra's absolute
aspect quality (Aldcroft et al. 2000). We perform our own
aperture photometry at the wavdetect positions of the 46
detected sources on this chip, which yields sources between
7 and 608 net counts (0.5--8 keV), and signal­to­noise ratio
(S/N) from 1.6 to 24.
3. MOSAIC OPTICAL IMAGING OBSERVATIONS
Several lensing cluster candidates have been suggested
from analysis of deep optical images. As part of ChaMP for
follow­up of Chandra serendipitous sources (Green et al.
2000), on UT 2000 September 29 and 2001 August 22 we
obtained images in Sloan g 0 , r 0 , and i 0 filters at the CTIO 4 m
Blanco Telescope using the wide­field MOSAIC camera,
which has eight 2048 # 4096 chips in a 4 # 2 array compris­
ing a 36 0 # 36 0 field of view. We reduced the images using
the MSCRED (v4.1) package (Valdes & Tody 1998; Valdes
1998) in the IRAF environment (Tody et al. 1986). This pro­
ceeds via the usual initial calibration, subtracting median­
combined bias frames and dividing by similarly prepared
dome flats. Additionally, we correct for electronic cross talk
2 CXCDS versions R4CU5UPD14.1, along with ACIS calibration data
from the Chandra CALDB 2.0b.
3 CIAOmay be downloaded from http://asc.harvard.edu/ciao.
722 GREEN ET AL. Vol. 571

between pairs of CCDs sharing readout electronics. We then
make a single image in each filter by median­combining
multiple object frames, e#ectively rejecting all of the celestial
objects in the final frame. From this we create a fringe­cor­
rection frame in each filter with the large­scale variation
removed and interactively subtract scaled versions from
every exposure in each given filter. We then generate the sky
flat for large scale corrections, again via median filtering,
this time of fringe­corrected object frames. This sky flat is
divided into every object frame.
From our MOSAIC images, we determined the (previ­
ously unpublished) positions of optically identified lens can­
didates discussed below, listed for convenience in Table 1.
These optical positions are based on an astrometric solution
of rms #0>1 obtained by matching detected objects to the
GSC2.2. 4 On each observing run, we obtained for this field
three dithered images per filter. After we flag bad pixels and
charge bleeds from severely saturated stars in each image,
we project them onto the tangent plane. Removal of remain­
ing large­scale gradients and scale di#erences between dith­
ered images allows for combination into a final single
stacked image in each filter. These nights were probably
photometric, since the scaling values were close to 1.
We used the MOSAIC images to determine the optical
flux ratios of the quasar pair in our three filters. These are
tabulated as magnitude di#erences in Table 2, along with
previously published values, for convenience. We discuss
the evident variability in x 5.3.
4. THE SEARCH FOR THE LENS
We now search for X­rays from any extended halo cen­
tered on the WSQP, and also measure X­rays from pub­
lished optical cluster candidates. In x 4.1 we describe our
search for X­ray emission from any lensing halo, and in
x 4.2 we compare the extended X­ray emission to the opti­
cally identified cluster candidates inferred from the distribu­
tion of galaxies in the field (Bonnet et al. 1993; Pello et al.
1996). In x 4.3 we discuss the implications of our measure­
ments for the lens hypothesis.
4.1. Search for Extended X­Ray Emission
No significant extended emission sources are evident to
the eye on the ACIS­S3 image. When searching for faint
extended sources, however, it is important to minimize
background contamination. The ACIS particle background
increases significantly below 0.5 keV and again at high ener­
gies. To optimize detection of a weak cluster signal, we first
filtered the cleaned, combined image to include only pho­
tons between 0.5 and 2 keV. We then masked out pixels
within a radius encompassing 95% of the encircled energy
around all point sources detected by wavdetect. This image
is divided by the appropriate ACIS­S3 exposure map, which
takes account instrumental features (e#ective area, quan­
tum e#ciency, telescope vignetting) as well as relative expo­
sure due to the dither pattern. Further division by the
exposure time yields a normalized image with pixel values in
photons cm #2 s #1 . To facilitate the search for nearby
B
A
Fig. 1.---Left: This 9 # 9 0 ACIS­S3 (0.3--8 keV) image of Q2345+007 (north up, east to the left) was smoothed with a 3 pixel (1>5) Gaussian for clarity and
shows the QSO pair on chip S3, about 1 north and 30 00 west of the chip center. While numerous point sources are detected, no extended cluster emission is dis­
tinguishable around the quasar pair on the image. Right: A 1 0 # 1 0 close­up of Q2345+007A,B shows that the pair is well resolved, with no detectable emission
from nearby point sources.
4 The Guide Star Catalogue­II is a joint project of the Space Telescope
Science Institute and the Osservatorio Astronomico di Torino. Space Tele­
scope Science Institute is operated by the Association of Universities for
Research in Astronomy, for the National Aeronautics and Space Adminis­
tration under contract NAS 5­26555.
No. 2, 2002 QSO PAIR Q2345+007 723

extended sources, we smoothed the point­source--subtracted
image using a 10 00 FWHM Gaussian. Figure 2 shows several
marginal excesses at #20% above the background level of
2 # 10 #10 photons cm #2 s #1 arcsec #2 (corresponding to 0.16
counts s #1 across the entire chip). Several low­level (#10%)
features in the normalized image in Figure 2 partially coin­
cide with features in the exposure map. This may be because
the telescope e#ective area varies in the 0.5--2 keV bandpass
and the assumed spectrum used to compute the exposure
map (monoenergetic 1.5 keV) may not represent the inci­
dent spectra very well across the image. More detailed simu­
lations for detecting extended structures in ACIS images are
called for, but this is beyond the scope of this paper. None
of the features were detected by our wavelet detection analy­
sis, so we do not consider any of the apparent fluctuations
to be significant unless they appear clearly in radial profiles.
In Table 1 we tabulate the X­ray positions (based on the
peak flux in the smoothed image) of apparent extended
excesses near Q2345+007. We summed the X­ray counts in
10 00 annuli centered on a point midway between the two
TABLE 1
Measured Positions of QSOs and Possible Cluster Centers
Center
R.A.
(2000.0)
Decl. a
(2000.0) Separation b Reference
Optical
QSO A,O ........................... 23 48 19.6 00 57 21.6 . . . 1
QSO B,O............................ 23 48 19.2 00 57 17.7 . . . 1
Center ................................ 23 48 19.4 00 57 19.2 . . . 1
C1 ...................................... 23 48 17.2 00 57 22 33 2
G1...................................... 23 48 16.5 00 58 08.3 66 2
G2...................................... 23 48 17.2 00 57 31.1 36 2
Z075 .................................. 23 48 16.6 00 57 31 44 3
Z120 .................................. 23 48 19.3 00 57 57 38 3
X­Ray
QSOA ............................... 23 48 19.6 00 57 21.4 0.3 1
QSO B................................ 23 48 19.2 00 57 17.5 0.2 1
CXO J234817.6+005717 ... 23 48 17.6 00 57 17 27 1
CXO J234816.9+005811 ... 23 48 16.9 00 58 11 64 1
CXO J234812.7+005813 ... 23 48 12.7 00 58 13 114 1
Note.---Units of right ascension are hours, minutes, and seconds, and units of declina­
tion are degrees, arcminutes, and arcseconds.
a Coordinates with subarcsecond precision are listed for objects with well­defined
centroids from our optical imaging.
b For QSOs, separation of X­ray centroids from optical centroids. For putative optical
clusters, separations are in arcseconds from the center of the QSO pair center.
References.---(1) This paper; (2) Bonnet et al. 1993; and (3) Pello et al. 1996.
TABLE 2
Magnitude Differences for Q2345+007A,B
Epoch Band m B #mA Reference Epoch Band m B #mA Reference
1981.92 ....... B 1.43 # 0.06 5 1991.96 ....... R 1.25 # 0.09 4
V 1.43 # 0.05 5 1992.73 ....... V 1.13 # 0.12 4
1981.98 ....... r 1.54 # 0.13 5 1992.88 ....... B 1.30 # 0.06 4
1982.57 ....... r 1.44 # 0.15 5 1992.93 ....... K 1.3 # 0.1 7
1982.90 ....... r 1.41 # 0.13 5 1993.76 ....... J 1.55 3
1989.65 ....... B 1.30 # 0.14 6 1993.76 ....... K 1.51 3
g 1.31 # 0.07 6 1998.73 ....... u* 1.62 # 0.12 8
r 1.24 # 0.06 6 g* 1.63 # 0.03 8
i 1.08 # 0.06 6 r* 1.60 # 0.05 8
1989.67 ....... r 1.28 # 0.04 6 i* 1.62 # 0.07 8
1989.69 ....... g 1.32 # 0.03 6 z* 1.37 # 0.16 8
r 1.27 # 0.04 6 2000.75 ....... g* 1.88 # 0.01 1
i 1.19 # 0.04 6 r* 1.85 # 0.01 1
1989.96 ....... K 1.11 # 0.17 4 i* 1.81 # 0.01 1
1990.79 ....... B J 1.24 3 2001.63 ....... g* 1.64 # 0.01 1
1990.79 ....... R 1.19 3 r* 1.65 # 0.01 1
1990.79 ....... I 1.14 3 i* 1.60 # 0.01 1
Note.---We include only ratios published with accurate dates in B band or redder.
References.---(1) This paper; (2) Bonnet et al. 1993; (3) Pello et al. 1996; (4) Gopal­Krishna et al. 1993; (5) Sol et al.
1984; (6) Weir & Djorgovski 1991; (7) McLeod et al. 1994; (8) SDSS Early Data Release.
724 GREEN ET AL. Vol. 571

QSO images and extending to a 2 0 radius, and we detect no
excess over the background. To derive an upper limit on the
extended source flux, we denote source counts S and back­
ground counts B (where B has been estimated from an area
A B but normalized to the source area A S ). 5 The random
1 # deviation in counts is # ¼ S × kB
Ï ÷ 12
, where
k ¼ 1 × A S =A B
Ï ÷
½ #. If we define a source count upper limit
at N standard deviations, then our source count upper limit
becomes
S ¼ N# ¼
N
2 # N × #################### N 2
× 4kB
p # :
We estimate the background from an area A B 4A S and
thereby derive a 95% upper limit of fewer than 44 counts
from any cluster within a circle of 1 0 radius centered on the
WSQP, corresponding to a count rate of 7 # 10 #4 s #1 . Con­
sidering a Raymond­Smith model with a reasonable range 6
of parameters and Galactic absorption, the corresponding
upper limit on the 0.5--2 keV flux is around 2 # 10 #15 ergs
cm #2 s #1 . This illustrates Chandra's excellent sensitivity to
faint sources, even when extended. 7
4.2. Observed Cluster Constraints
Our 0.5--2 keV flux upper limit of 2 # 10 #15 ergs cm #2 s #1
within 1 0 of the WSQP clearly yields a stringent limit on the
existence of any lensing cluster. Assuming a T ¼ 2 3 keV
spectrum, this upper limit corresponds at z # 1 to a 0.5--2
keV (rest­frame) luminosity of 1:2 # 10 43 ergs s #1
(1:6 # 10 43 for # ¼ 0:3). An r ¼ 1 0 aperture at z ¼ 1 corre­
Fig. 2.---Chandra ACIS­S3 Image of the 9<7 field surrounding Q2345+007, with north up, east to the left. We removed counts within these apertures and
then divided by an exposure map before smoothing with a 10 00 Gaussian. The resulting gray scale shows a mean flux of #2 # 10 #10 photons cm #2 s #1 pixel #1 ,
with values across the image ranging from about 0.7 to 5.3 in those units. Circles with sizes representing the point­spread function (95% encircled energy) mark
the positions of point sources detected by CIAO wavdetect. The positions of the twin QSOs are evident just northeast of center. The large contours show linear
levels (from 500 to 900 cm 2 , in steps of 100) in the exposure map. Positions of putative optically identified galaxy clusters are marked in bold type.
5 Source counts are derived from total counts T in area A S via
S ¼ T # B.
6 The unabsorbed flux is derived, i.e., from outside the Galaxy. Changes
in the assumed spectral models from 2 to 8 keV result in changes of #15%
in the value of the derived flux.
7 If there is any residual flux from the WSQP scattered beyond the radius
containing 95% of the total energy, our estimated cluster flux upper limit is
conservative.
No. 2, 2002 QSO PAIR Q2345+007 725

sponds to 0.5 Mpc (0.6 Mpc for # ¼ 0:3) and encloses about
half of the total luminosity for a typical cluster with the sur­
face brightness described by a #­model with a core radius of
0.25 Mpc and # ¼ 0:6 (Jones & Forman 1984). With this
correction, our upper limit corresponds to a typical lumi­
nosity of a galaxy group with T # 2 3 keV (e.g., Hwang et
al. 1999 and references therein), so our estimate is self­con­
sistent. All luminosities, sizes, and distances are calculated
assuming H 0 ¼ 50 km s #1 Mpc #1 , # 0 ¼ 1:0, and # ¼ 0:0
unless otherwise noted (we also give # 0 ¼ 0:3 values). This
represents by far the strongest constraint on the X­ray lumi­
nosity of any dark lens candidate to date (see Chartas et al.
2001).
It is certainly possible that a massive cluster that is not
centered on the WSQP could also produce the observed
image splitting. Therefore, we also investigate apparent flux
excesses in the vicinity (<1>5) of the quasar pair. However,
the required mass of the lensing cluster would rise with
angular distance from the WSQP, so that our flux limits
constrain the existence of such a cluster even more strongly
at larger angles.
The closest significant excess of extended X­ray emission
to the WSQP is CXO J234817.6+005717. Centered about
27 00 from the center of the WSQP, it has a peak flux of about
0.0544 counts cm #2 s #1 arcsec #2 . The radial profile of this
excess, determined with possible nearby extended sources
removed, is shown in Figure 3 (see also Fig. 4). The excess
contributes just 11 # 5 counts above background. At just
2.2 #, we find this source to be of questionable certitude. A
more credible 2 # upper limit to the flux yields 9 # 10 #16 ergs
cm #2 s #1 from 0.5--2 keV.
Bonnet et al. (1993) identified a cluster center they label
C1 from a weak gravitational shear field pattern suggesting
a lens velocity dispersion # v # 1200 km s #1 . The peak X­ray
flux of the above extended source corresponds to a position
just 8 00 from C1. Pello et al. (1996) claimed an excess of gal­
axies with photometric redshifts z # 0:75, which is 13 00 from
C1 and 21 00 from CXO J234817.6+005717. However, cent­
roids may naturally have somewhat di#erent positions
because optical galaxies may not follow the overall mass dis­
tribution, and the cluster also may not be virialized. Assum­
ing that all three objects (C1, the Pello et al. z ¼ 0:75 optical
galaxy excess, and CXO J234817.6+005717) can be identi­
fied as the same object, the 0.5--2 keV luminosity upper limit
is 2:8 # 10 42 ergs s #1 (3:6 # 10 42 ergs s #1 for # ¼ 0:3), more
similar to an isolated elliptical galaxy or a small group than
to a rich # v # 1200 km s #1 cluster, suggesting that the large
velocity dispersion estimate is the result of a line­of­sight
projection.
In any case, at a transverse distance of #220 h 50 kpc, such
a mass is too small to produce the observed pair separation
by lensing.
Bonnet et al. also identified the galaxy G2 8 as having
a position consistent with C1 within the errors (#9 00 from
the best centroid of C1). At 15 00 from G2, CXO
J234817.6+005717 is clearly not consistent with emission
from that galaxy. Another apparent X­ray excess CXO
J234816.9+005811 is an arcminute from the WSQP cent­
roid, but 7 00 from the galaxy G1 identified by Bonnet et al.
(1993). However, our radial profile centered either on CXO
J234816.9+005811 or on the position of G1 reveals no sig­
nificant excess above background, even when the other
nearby extended sources are excised.
4.3. Comparison to Required Lens
We start by assuming that the WSQP is a lens produced
by a simple, singular isothermal sphere (SIS; see Schneider,
Ehlers, & Falco 1992). The image separation
D# ¼ 8#Ï# v =c÷ 2
D LS =D OS depends only on the velocity dis­
persion of the potential # v and the ratio of the comoving dis­
tances between the lens and the source, D LS , and the
observer and the source, D OS . Since we failed to detect a lens
cluster, we will get the most conservative limits if we assume
that the lens lies at the `` minimum flux redshift,'' the red­
shift that would minimize the observed X­ray flux. If we
neglect K­corrections, the flux from the lens is
F ¼ L
4#D 2
OL Ï1 × z l ÷ 2 / D OS r H
DOLD LS Ï1 × z l ÷
# # 2
; Ï1÷
where r H is the Hubble radius c=H 0 . This flux diverges at
low redshift because of the proximity of the cluster and at
high redshift because of the mass of the cluster. For # 0 ¼ 1
the flux is minimized at z l ¼ 0:92, which we will round to
z l ¼ 1 for simplicity.
Now using a SIS model for the lensing mass, the large
image separation of 7>3 implies a cluster velocity dispersion
of # v ¼ 860 km s #1 or a cluster mass of 1:3 # 10 14 M # . We
emphasize that this is the minimum enclosed mass or disper­
Fig. 3.---Radial profile is centered at the position of CXO
J234817.6+005717, the excess nearest the QSOs in Fig. 4. We used succes­
sive 10 00 annuli in a cleaned 0.5--2 keV image, excluding both detected point
sources and regions with apparent source excesses in Fig. 2. Vertical error
bars are 1 #, and horizontal bars represent the range of each annulus. The
background level (#0.01 counts per 0>5 pixel) has not been subtracted. The
profile displays a possible flux excess in the first bin, significant at 2.5--3 #
above background.
8 G1 is the brighter galaxy. Labels in Bonnet et al. (1993) for G1 and G2
are incorrectly swapped in all but their Fig. 2.
726 GREEN ET AL. Vol. 571

sion required at this redshift to induce the observed pair sep­
aration. Combining the LX ­# v relation from, e.g., Mulchaey
& Zabludo# (1998) and the LX ­T relation from, e.g., Mar­
kevitch (1998), and neglecting any possible cosmological
evolution of these relations for a qualitative estimate, we
obtain LX Ï0:5 2 keV÷ # 2 # 10 44 ergs s #1 and T # 5 keV
for such a cluster. At z ¼ 1, this corresponds to
f X Ï0:5 2 keV÷ # 4 # 10 #14 ergs cm #2 s #1 (3 # 10 #14 for
# ¼ 0:3). Again assuming a typical cluster brightness distri­
bution and dividing by 2 to convert to the r ¼ 1 aperture,
we can see that our 95% flux limit is an order of magnitude
below this minimum required flux estimate. If the lens were a
dark cluster lacking not only galaxies but also gas, the gas
fraction would have to be #3 times lower than that in
known clusters. At low redshift, all well­studied clusters (at
least in the relevant range of radii and masses) have similar
values of the gas fraction (e.g., Mohr, Mathiesen, & Evrard
1999; Vikhlinin, Forman, & Jones 1999) so such a deviation
appears to be extremely unlikely. We conclude that a single
cluster acting as a lens is not a plausible scenario for
Q2345+007.
In other lenses clearly due to a combination of a cluster
and a galaxy (particularly Q0957+561; Keeton et al. 2000),
a massive, luminous lens galaxy dominates the image split­
ting. Here we see no such candidate galaxy, even in the
infrared, to a limit of approximately L # =10 near redshift
unity. Such a galaxy, unless completely di#erent from all
other known lens galaxies (e.g., Kochanek et al. 2000, Xan­
thopoulos et al. 1998), must make a negligible contribution
to the overall image separation and modifies our estimate of
the critical radius of the putative dark halo only by a factor
of 1 # 2b gal =D# with b gal 5 1>0.
5. ARE THEY IMAGES OF THE SAME QUASAR?
We have failed to find a lens. While a dark matter lens is
not ruled out, the pair could also be shown to be a binary
quasar by evidence that the components' spectral energy
distributions are di#erent. We explore this by first determin­
ing the X­ray properties of the two quasars in x 5.1, followed
by a discussion of the possible e#ects of absorption and
extinction on the flux ratios in x 5.2. In x 5.3 we compare the
flux ratios from the near­infrared to the X­ray bands to con­
clude that they are probably di#erent quasars. We discuss
the remaining puzzle of the strikingly similar optical/UV
spectra of the components in x 5.4.
23:48:12
15
18
21
24
27
0:57:00
56:15
55:30
45
58:30
59:15
B
A
G1
z120
z075
C1
Fig. 4.---An r 0 image from the CTIO 4 m telescope from UT 2000 September 29 of a 4<25 field surrounding Q2345+007. North is up, east is to the left, and a
J2000.0 coordinate grid is displayed. The large contours show the same X­ray flux levels as in Fig. 2. Positions of the QSO images A and B and of putative opti­
cally identified galaxy clusters are marked in bold type.
No. 2, 2002 QSO PAIR Q2345+007 727

5.1. The X­Ray Properties of the Quasars
The total broadband (0.3--8 keV) counts from the QSOs
are 358:5 # 20 and 54:8 # 8 for components A and B,
respectively. The broadband X­ray flux ratio is 6:55 # 0:50.
As a first check of whether the QSOs have similar
X­ray spectra, we measure their hardness ratio,
HR ¼ ÏH # S÷=ÏH × S÷, using the hard counts (H ) from
2.5--8.0 keV and the soft counts (S) from 0.3--2.5 keV. While
di#erences in instrument and telescope calibration are com­
pletely negligible at these spacings, the benefit of comparing
hardness ratios directly are that no spectral models need be
fitted to the data. We find HRA ¼ #0:89 # 0:02 and
HR B ¼ #0:55 # 0:11, which are significantly di#erent (at
3 #).
However, although optical magnitudes suggest little dif­
ferential reddening (see x 5.2 below), absorption along dif­
ferent sight lines to the QSO might still produce a significant
di#erence in HR if the absorbers are not dusty. We also cal­
culate the X­ray flux ratio in a bandpass that is much less
subject to absorption. In the 1.6--8 keV band (5--25 keV in
the QSO rest frame) the X­ray flux ratio is smaller,
2:3 # 0:5. This large di#erence from the broadband flux
ratio is expected given their di#ering hardness ratios.
A test for di#erential absorption requires X­ray spectral
modeling. For spectral analyses, we used the latest
(CALDB2.7) ACIS­S3 FITS embedded function (FEF) files
and corresponding gain tables. The new FEFs are analytic
fits based on a physical model of the back­illuminated CCD,
developed by the MIT ACIS/IPI team, and adjusted to the
on­orbit gain of the S3 chip as determined by the flight cali­
bration sources. These FEFs are used to generate a response
matrix file (RMF), which maps the incident photon energy
to ACIS pulse height (deposited charge). An ancillary
response file (ARF) calibrates the e#ective collecting area of
a specified source region on ACIS as a function of incident
photon energy. We extract ACIS PI spectra from a 3>5
region around each QSO, using the psextract script
described in the standard thread for CIAO2.1. This script
creates an aspect histogram file and the RMF and ARF cali­
bration files appropriate to the source position on chip
(which is time­dependent due to dither) and CCD tempera­
ture (#120 C).
We first group the Chandra spectra into bins containing
at least 10 counts each. We test four spectral models to test
permutations that link or decouple the power law and
absorbing columns of Q2345+007A,B. Results for all four
models are compiled in Table 3. Model 1 fixes (z ¼ 0)
absorption at the Galactic value but allows the power­law
slope and flux normalization to vary for each QSO as
NÏE÷ ¼ A i E ## e #N Gal
H #ÏE÷ photons cm #2 s #1 keV #1 :
In this formula, A i is the normalization and # i is the power­
law photon index for the individual QSOs; N H is the equiva­
lent Galactic neutral hydrogen column density
3:81 # 10 20 cm #2 , which characterizes the e#ective absorp­
tion (by cold gas at solar abundance), with #ÏE÷ being the
corresponding absorption cross section (Morrison &
McCammon 1983). For determining the best­fit parameter
values, we use Powell optimization with Primini statistics
(Kearns, Primini, & Alexander 1995). This fit yields a typi­
cal continuum slope #A ¼ 2:19 # 0:15 for the brighter com­
ponent (# 2
# ¼ 1:1). The fainter QSO shows a best­fit of
# B ¼ 0:79 # 0:4 (# 2
# ¼ 0:6). These confirm our impression
from the hardness ratio that the observed energy distribu­
tions are not consistent.
There are not su#cient counts from B to independently
fit both its absorbing column and power­law slope. How­
ever, a somewhat stronger test than model 1 is available by
assuming that A and B have identical total absorbing col­
umns and fitting that column simultaneously (model 2). The
result is a column NH ¼ 5:3 # 3:1 # 10 20 cm #2 , which yields
slopes #A ¼ 2:30 ×0:36
#0:30 and # B ¼ 0:83 ×0:49
#0:44 . Therefore, if the
X­ray emission from A and B su#er identical absorption
along the line of sight, we can confidently say that their
intrinsic spectral energy distributions are inconsistent. Con­
tour plots for this simultaneous power­law fit are shown in
Figure 5. Not surprisingly (since A dominates the total
counts), we find similar results when using for B the best­fit
column from A alone.
Rather than accepting the power­law di#erence as intrin­
sic, we might entertain the possibility that B su#ers from
strong absorption, since this can mimic a flat power­law
slope in low S/N X­ray spectra. If we assume that the sys­
tem is lensed, then the underlying X­ray power­law contin­
uum slopes should be indistinguishable, even though their
line­of­sight absorption may vary. We thus fit model 3 by
assuming a single power­law slope, but allowing for di#er­
ent absorbing columns between the two components. The
best­fit continuum slope is # ¼ 2:14 ×0:34
#0:28 , and the absorbing
columns are quite di#erent: NH;A ¼ 3:4 ×3:5
#2:9 # 10 20 cm #2 and
NH;B ¼ 43:9 ×29
#18 # 10 20 cm #2 : Contour plots for this simul­
taneous absorption column fit are shown in Figure 5.
If we also allow the absorption columns to di#er between
A and B (model 4), the di#erence between the best­fit
power­law slopes becomes insignificant. This is because, as
expected, the small number of counts in B are insu#cient to
constrain both parameters. While the parameters for A
change little, for B, we find NH;B ¼ 23:6 ×28:1 #19:2 # 10 20 cm #2
(consistent with no absorption) and # B ¼ 1:37 ×0:79
#0:66 .
The observed range of continuum slopes in radio­quiet
quasars is 1:5 < # 2#10keV < 3 with a mean of # ¼ 2:0 # 0:25
(1 # dispersion; George et al. 2000). Therefore, the very flat
slope measured here for B might be unrealistic, possibly
a#ected by line­of­sight absorption. Higher S/N X­ray spec­
TABLE 3
Spectral Fit Parameters
Model QSO C
N H
(10 20 cm #2 ) # 2 (dof) a
1................ A 2.19 # 0.15 3.8 42.54 (40)
B 0.79 # 0.4 . . . 25.09 (40)
2................ A 2:30 ×0:36
#0:30 5.3 # 3.1 67.9 (79)
B 0:83 ×0:49
#0:44 . . . . . .
3................ A 2:14 ×0:34
#0:28 3:4 ×3:5
#2:9 65.4 (79)
B . . . 43:9 ×29
#18 . . .
4................ A 2:23 ×0:33
#0:30 4:3 # 3:1 42.77 (39)
B 1:37 ×0:79
#0:66 23:6 ×28:1
#19:2 21.32 (39)
Note.---Fit parameters based on simultaneous fitting of spectra
using Primini statistics in Sherpa (Freeman et al 2001). Uncertainties
are 90% confidence limits. Where no uncertainties are shown, the
parameter was frozen at the value displayed. Where only QSO A
shows a value, the parameter was fitted simultaneously to both com­
ponents A and B. Models are described in the text.
a Values of # 2 are based on spectra binned to 10 counts per bin,
using given fit parameters.
728 GREEN ET AL. Vol. 571

troscopy would reveal whether it is the absorption, the
intrinsic continua, or both that di#er between the two
images. Only repeated observations, a daunting investment,
can definitively reveal whether the observed di#erences are
due to variations in a single lensed QSO whose images su#er
di#ering time delays. Since the evidence for absorption is
weak, if Q2345+007A,B are indeed lensed, this would pre­
dict that a program to monitor a suitable sample of single
luminous QSOs over a period of a few years would result in
the discovery of large temporal variations in the X­ray con­
tinuum slopes of individual objects. 9 Su#ciently large
changes in either the column or continuum slope in single
QSOs have not been observed to date (e.g., Lawson &
Turner 1997).
5.2. Intrinsic or Intervening Absorbers
Could a lens+absorption interpretation explain other
data? From our MOSAIC imaging, or from the colors pub­
lished by Pello et al. (1996), the di#erence between g 0 # r 0
and r 0 # i 0 for the two images is negligible (#0:032 # 0:02
and 0:037 # 0:02, respectively). The optical colors show lit­
tle evidence for reddening relative to the main stellar locus,
and indeed are quite blue in g 0 # r 0 , as expected for QSOs at
this redshift (Richards et al. 2001). Thus any putative
absorber is likely to be warm (ionized). Warm X­ray
absorbers are usually accompanied by a significant decrease
in relative X­ray strength, as might be in evidence in the
larger # ox ¼ 1:7 of Q2345+007B, but also by evident
absorption in the rest­frame ultraviolet spectrum, particu­
larly in the blue wing of C iv (Brandt, Laor, & Wills 2000).
High S/N rest­frame UV spectra of Q2345+007A,B
(SS91) show little di#erence between the C iv profiles, and
no strong intrinsic absorption. Indeed, absorption lines of
intervening C iv systems at z abs ¼ 1:798, 1.799, and 1.983
are all stronger in A than in B. Near the C iii] emission line,
B shows stronger absorption lines in the blue wing, while A
shows an overall lower profile. However, this absorption is
most likely due to intervening (not intrinsic) absorption at
z ¼ 1:7717 (SS91). Intervening absorbers such as these in
galactic halos or disks have UV columns and ionization lev­
els that are far too small to cause detectable absorption of
the X­ray emission (O'Flaherty & Jakobsen 1997). Interest­
ingly, an intervening C iv absorber at z # 0:75 (same as the
putative cluster C1) would not show strong lines within the
wavelength range of the existing spectra. SS91 find evidence
for a weak (W # # 1) Mg ii absorption doublet in
Q2345+007B at this redshift. The problem is that an ionized
absorber at z ¼ 0:75 would be optically bright (R < 22:5;
Cohen et al. 2000) and easily detected in the deep optical
imaging.
A B ¼ 25 galaxy that might cause such intervening
adsorption is detected in the wings of the PSF of
Q2345+007B in a deep ground­based image (Fischer et al.
1994), but no such candidates are seen in a HST/NICH­
MOS image to H > 22:5 (Munoz et al. 1998). Mg ii absorb­
ers are usually associated with galaxies of Hubble types E
through Sb (Steidel 1993). Even a normal L* Sc galaxy
would be 2 mag brighter than the H magnitude limit at this
redshift. An early­type L* galaxy would be at least 4 mag
brighter (see the closely related discussion for
Q1634+267A,B in Peng et al. 1999). So no reasonable can­
didates exist for a line­of­sight absorber to account for the
X­ray properties of Q2345+007B.
5.3. Optical/X­Ray Flux Ratios of the WSQP
For comparison to the optical and infrared bands, the
epoch 1990 flux ratios in B; R; I ; J, and K bands from Pello
et al. (1996) are 3.1, 3.0, 2.9, 4.2, and 4.0 respectively. The
di#erence between these flux ratios and those in the X­ray
band (6.55, corresponding to 2.0 on a magnitude scale)
seems significant enough to rule out the lens interpretation.
Weir & Djorgovski (1991) found evidence for variation in
the magnitude di#erence between the two components, and
Gopal­Krishna et al. (1993) argued that detection of a time
delay is crucial to proving that the WSQP is indeed a lens. If
the pair is lensed, then the expected time delay is about a
year. A comparison with optical flux ratios more nearly con­
temporaneous to the X­ray observation is desired. The dates
of our own optical imaging do not match the Chandra
observations, but begin less than 2 months after. We com­
pile in Table 2 the magnitude di#erences between
Q2345+007A,B in several optical filters. Flux ratios
extracted from our (epoch 2000.75) MOSAIC image using
SExtractor (Bertin & Arnouts 1996) in Sloan g 0 , r 0 , and i 0 fil­
Fig. 5.---Contour plots of confidence levels for simultaneous fits to ACIS spectra of Q2345+007A,B. The plot at left shows the confidence levels for C A vs.
C B in model 2, where the absorption is assumed to be common. The plot at right shows confidence levels for NH;A vs. NH;B in model 3, where the power­law
slope is assumed to be common.
9 Analogous predictions for optical spectroscopic variability by SS91
were later confirmed by Small et al. (1997).
No. 2, 2002 QSO PAIR Q2345+007 729

ters are 5.6, 5.5, and 5.3, respectively. With errors at most
#0.1 for each ratio, these all agree with each other and sug­
gest no strong di#erential reddening of a lensed pair. How­
ever, these optical ratios still di#er from the broadband X­
ray flux ratio at the #2 # level. Our epoch 2001.63 images
show ratios of 4.5, 4.6, 4.4, which di#er at greater than 3 #
from the X­ray flux ratio. 10 The X­ray (0.3--8 keV) flux
ratios we measure during the separate Chandra observations
of 2000 May and June are 8:2 # 2:1 and 6:2 # 1:0, respec­
tively. Given the evidence of variability, we cannot rule out
the possibility that at the epoch of the X­ray observations
the optical ratios were identical to the X­ray flux ratio.
Rusin (2002) notes that quasar pairs with flux ratios greater
than 3 : 1 are statistical outliers in the expected distribution
of dark lenses and are therefore likely a priori to be binary
quasars.
KFM99 developed a statistical proof for comparing the
optical and radio properties of the WSQP population to
demonstrate that most WSQPs must be binary quasars
rather than lenses, and that their incidence was explained by
the triggering of quasar activity during galaxy mergers.
KFM99 classified the quasar pairs as both radio­faint
(denoted O 2 ), both radio­bright (O 2 R 2 ), or radio­mixed
(O 2 R). The radio­mixed pairs have wildly discrepant opti­
cal/radio flux ratios (by factors of 50 or more) and are
clearly binary quasars rather than lenses. The ratio of radio/
optical power spans 10 4 in optically selected quasars, and
the radio­loud fraction is PR # 10% (e.g., Hooper et al.
1995). Hence the existence of each O 2 R pair implies the exis­
tence of 1=2P R # 5 O 2 pairs, and KFM99 could show statis­
tically that at most 8% (22%) of the WSQP pairs could be
gravitational lenses at a 1 # (2 #) limit. These limits still
admit a population of optically dark clusters equal to the
normal cluster population, so the possibility of optically
dark clusters still remains.
By imaging an O 2 quasar pair with Chandra we can see if
the X­ray and optical flux ratios are discrepant---an
`` O 2 X '' pair is not a lens. The only disadvantage of using
quasar X­ray emission is that it is more common than radio
emission and is even considered the signature of an active
nucleus. Nonetheless, X­ray--to--optical flux ratios span a
range of #400 (Pickering, Impey, & Foltz 1994), so it is
probable that binary quasars will show discrepancies in
their # ox values.
Assuming a power­law continuum slope of # ¼ 2:4 from
spectral fitting of A and the Milky Way gas column, the
observed count rate corresponds 11 to an (absorbed) X­ray
flux in this band for QSO A of FX;A ¼ 2:52 # 10 #14 ergs
cm #2 s #1 , while for B the flux is 0.386 in the same units.
From the unabsorbed values (20% larger for this slope and
N H ), the inferred monochromatic rest­frame luminosities at
2 keV are logL 2 keV ¼ 26:84 and 26.02 (in ergs s #1 Hz #1 ), or
45.11 and 44.30 broadband in ergs s #1 .
Using the optical B magnitudes from Pello et al. (1996),
the logarithm of the monochromatic 2500 A š optical lumi­
nosities (in units ergs s #1 ) are 31 and 30.5, for A and B,
respectively. These correspond to M A
B ¼ #26:08 and
M B
B ¼ #24:84 in absolute magnitudes at 4400 A š (Schmidt &
Green 1983). We assume an optical spectral index of
# ¼ #0:5, consistent with that found in SDSS QSOs (with a
spread of 0.65; Richards et al. 2001), with specific optical
normalization from Marshall et al. (1984). The resulting # ox
values 12 are 1.60 and 1.72. The value for A is quite normal
for optically selected QSOs (e.g., Green et al. 1995; Yuan et
al. 1998), while the value of 1.72 for B is more similar to the
values found for absorbed QSOs (Green et al. 2001; Brandt
et al. 2000). If there is strong absorption along the sight line
to QSO B, it is not detectable in an X­ray spectrum contain­
ing only 55 counts, but more importantly, not evident in its
optical colors or rest­frame UV spectroscopy, as discussed
above.
5.4. Optical Spectra
The exceptionally detailed intrinsic similarity in the opti­
cal spectra of the quasar images (SS91) remains the obsti­
nate core of a long­running mystery. The di#erences seen
are primarily in intervening absorbers, and in the pairs' C
iii] emission­line strength. The intervening absorbers are
expected to di#er in either the lens or binary scenario, but
more so in the latter. But accounting for di#erences in the
emission­line profiles is more problematic. Microlensing by
stars in the lens could account for #10% di#erences at most
(Schneider & Wambsganss 1990). 13 However, in a lens sce­
nario, the di#erences are most likely due to intrinsic profile
variations in a single quasar that occur on a characteristic
timescale less than the time delay between the two images,
roughly a year (9 months for z lens ¼ 1:5; #1.8 yr for
z lens ¼ 0:75). Small et al. (1997) confirmed that the observed
di#erences between emission lines in the Q2345+007 images
are consistent with the spectral di#erences seen on 1--1.5 yr
timescales in individual QSOs. So the few observed spectral
di#erences do not rule out the lens interpretation.
Given that the current study refutes the lens hypothesis, is
it likely for two distinct QSOs to have spectra that are as
similar as observed? Quasar broad emission line strengths
and profiles seem to be loosely correlated with other QSO
properties, such asluminosity, # ox , and narrow emission line
strength (e.g., Baldwin 1977; Laor et al. 1997; Green et al.
2001). The great similarity in emission lines in
Q2345+007A,B argues that, if they are distinct objects,
their luminosity, the geometry of their emission regions, and
their inclination angle should all be similar. Detailed empiri­
cal tests on the probability of achieving such similar spectra
at random among a sample of unrelated quasars require
large databases of quasar spectra, uniformly analyzed, and
with similar S/N and resolution. Such studies are beyond
the scope of this paper, but adequate measurement data sets
are now becoming available (Forster et al. 2000).
6. SUMMARY
We conclude that Q2345+007A,B is a quasar binary.
First, we find no evidence for a halo in the field su#ciently
massive to produce the image separation, provided that
halo contains at least 1/3 of the normal cluster baryon frac­
10 The mean change in magnitude di#erence for stars we find in the same
filter at the two di#erent epochs is 0 # 0:12 mag.
11 We use the Chandra Portable, Interactive Multi­Mission Simulator
(PIMMS; originally Mukai 1993).
12 # ox is the slope of a hypothetical power law from 2500 A š to 2 keV;
# ox ¼ 0:384 log L 2500 =L 2 keV . Use of the nominal best­fit slope of # B ¼ 0:9
for B does not significantly a#ect the # ox values (<1% change).
13 The modeling of microlensing e#ects to date has assumed Keplerian
cloud orbits in the broad­line region, which yield greater predicted di#eren­
ces than infall models. More recent outflow disk wind models (Murray &
Chiang 1997; deKool & Begelman 1995; Elvis 2000) are highly aspect
dependent and may yield larger di#erences in microlensing models.
730 GREEN ET AL. Vol. 571

tion. Mechanisms may exist to suppress star formation in
galaxies (e.g., Elmegreen & Parravano 1994) or to produce
only exceptionally low surface brightness galaxies (e.g.,
Sprayberry et al. 1995), thereby diminishing the optical
luminosity of cluster members. There is no such mechanism
known that could significantly reduce the baryon fraction in
a cluster potential well. The energy injection from star for­
mation or active galactic nuclei (AGNs) would need to be
adequate not only to heat and eject gas from the galaxies,
but also from the massive cluster potential. Even if the ener­
getics were plausible, we are unable to find any of the con­
comitant optical or infrared emission (e.g., McLeod, Rieke,
&Weedman 1994).
Second, the two quasars have X­ray properties inconsis­
tent with the lens hypothesis. The X­ray flux ratios are
inconsistent with the optical flux ratios. Their X­ray spectral
properties also di#er significantly, and we find no evidence
for either extinction or absorption that might explain the
inconsistency in terms of the lens hypothesis.
Significantly deeper spatially resolved X­ray spectroscopy
could show more definitively that the pair is a binary. The
best available hope to prove the pair is lensed would be
detection of a time delay in correlated variability. Optical or
near­IR detection of host galaxies in the QSO images could
also provide strong arguments for or against the lens
hypothesis.
The X­ray evidence in the field of Q2345+007 does not
support the lens hypothesis. This WSQP is thus likely to be
a binary quasar, and it currently stands as the example with
the highest redshift, the largest separation, and the most
detailed agreement between optical/UV spectra of its com­
ponents. We may be studying a pair of luminous QSOs
whose hosts, separated by #60 h 50 kpc, have a history of
dynamical interaction. Plausible models hold that quasar
activity is triggered by tidal interactions in a galactic merger
but that activation of the galactic nuclei occurs late in the
interaction, when the nuclei are within 80 # 30 kpc of each
other (Mortlock, Webster, & Francis 1999). Simple dynami­
cal friction models reproduce the observed distribution of
projected separations of pairs but predict that binary qua­
sars are only observable as such in the early stages of galac­
tic collisions, after which the supermassive black holes
would orbit within the merger remnant. We speculate that
Q2345+007 may represent the highest redshift example
known of an interaction­triggered but as­yet unmerged
luminous AGN.
We thank Frank Valdes, Lindsey Davis, and the IRAF
team for writing and helping with the MSCRED package.
Thanks also to Brian MacLean at STScI for early help with
the GSC2.2. Thanks to Ani Thakar at Johns Hopkins for
help retrieving SDSS data. Funding for the creation and dis­
tribution of the SDSS Archive has been provided by the
Alfred P. Sloan Foundation, the participating institutions,
the National Aeronautics and Space Administration, the
National Science Foundation, the US Department of
Energy, the Japanese Monbukagakusho, and the Max
Planck Society. The SDSS Web site is http://www.sdss.org.
This work was supported by Chandra grant GO 0­1161X
and NASA grant NAS 8­39073. A. D., P. J. G., D. K., M.
M., and A. S. acknowledge support through NASA con­
tract NAS 8­39073 (CXC). C. S. K. is supported by NASA
grants NAG 5­8831 and NAG 5­9265. We are grateful to
Tony Tyson for providing deep images of the field and for
his comments on the text.
REFERENCES
Aldcroft, T. L., Karovska, M., Cresitello­Dittmar, M. L., Cameron, R. A.,
& Markevitch, M. L. 2000, Proc. SPIE, 4012, 650
Baldwin, J. A. 1977, ApJ, 214, 679
Bertin, E., & Arnouts, S. 1996, A&AS, 117, 393
Bonnet, H., Fort, B., Kneib, J.­P., Mellier, Y., & Soucail, G. 1993, A&A,
280, L7
Brandt, W. N., Laor, A., & Wills, B. J. 2000, ApJ, 528, 637
Chartas, G., Bautz, M., Garmire, G., Jones, C., & Schneider, D. P. 2001,
ApJ, 550, L163
Cohen, J. G., et al. 2000, ApJ, 538, 29
de Kool, M., & Begelman, M. C. 1995, ApJ, 455, 448
Ebeling, H., Edge, A. C., Allen, S. W., Crawford, C. S., Fabian, A. C., &
Huchra, J. P. 2000, MNRAS, 318, 333
Elmegreen, B. G., & Parravano, A. 1994, ApJ, 435, 121
Elvis, M. 2000, ApJ, 545, 63
Erben, T., van Waerbeke, L., Mellier, Y., Schneider, P., Cuillandre, J.­C.,
Castander, F. J., & Dantel­Fort, M. 2000, A&A, 355, 23
Fischer, P., Tyson, J., A., Bernstein, G. M., & Guhathakurta, P. 1994, ApJ,
431, L71
Forster, K., Green, P. J., Aldcroft, T. L., Vestergaard, M., Foltz, C. B., &
Hewett, P. C. 2001, ApJS, 134, 35
Freeman, P. E., Kashyap, V., Rosner, R., & Lamb, D. Q. 2002, ApJS, 138,
185
George, I. M., Turner, T. J., Yaqoob, T., Netzer, H., Laor, A., Mushotzky,
R. F., Nandra, K., & Takahashi, T. 2000, ApJ, 531, 52
Gopal­Krishna, Y., M., Wiita, P. J., Smette, A., Pati, A., & Altieri, B. 1993,
A&A, 280, 360
Green, P. J., Forster, K., & Kuraszkewiecz, J. 2001, ApJ, 556, 727
Green, P., J., et al. 1995, ApJ, 450, 51
---------. 2000, BAAS, 31, 1493
Holder, G. P., Mohr, J. J., Carlstrom, J. E., Evrard, A. E., Leitch, & Erik
M. 2000, ApJ, 544, 629
Hooper, E. J., Impey, C. D., Foltz, C. B., & Hewett, P. C. 1995, ApJ, 445,
62
Hwang, U., Mushotzky, R. F., Burns, J. O., Fukazawa, Y., & White, R. A.
1999, ApJ, 516, 604
Jones, C., & Forman, W. 1984, ApJ, 276, 38
Kearns, K., Primini, F., & Alexander, D. 1995, in ASP Conf. Ser. 77, ed.
R. A. Shaw, H. E. Payne, & J. J. E. Hayes (San Francisco: ASP), 331
Keeton, C. R., et al. 2000, ApJ, 542, 74
Kim, D.­W., et al. 2000, AASHEAD Meeting, 32, 2604
Kochanek, C. S., Falco, E. E., Impey, C. D., Lehar, J., McLeod, B. A., Rix,
H.­W., Keeton, C. R., & Munoz, J. A. 2000, ApJ, 543, 131
Kochanek, C. S., Falco, E., & Munoz, J. A. 1999, ApJ, 510, 590 (KFM99)
Laor, A., Fiore, F., Elvis, M., Wilkes, B. J., & McDowell, J. C. 1997, ApJ,
477, 93
Lawson, A. J., & Turner, M. J. L. 1997, MNRAS, 288, 920
Markevitch, M. 1998, ApJ, 504, 27
Marshall, H. L., Avni, Y., Braccesi, A., Huchra, J. P., Tananbaum, H.,
Zamorani, G., & Zitelli, V. 1984, ApJ, 283, 50
McLeod, B., Rieke, & M., Weedman, D. 1994, ApJ, 43, 528
Mohr, J. J., Mathiesen, B., & Evrard, A. E. 1999, ApJ, 517, 627
Morrison, R., &McCammon, D. 1983, ApJ, 270, 119
Mortlock, D. J., & Webster, R. L. 2000, MNRAS, 319, 872
Mortlock, D. J., Webster, R. L., & Francis, P. J. 1999, MNRAS, 309, 836
Mukai, K. 1993, Legacy, 3, 21
Mulchaey, J. S., & Zabludo#, A. I. 1998, ApJ, 496, 73
Munoz, J. A., Falco, E., Kochanek, C. S., Lehar, J., McLeod, B. A., Impey,
C. D., Rix, H.­W., & Peng, C. Y. 1998, Ap&SS, 263, 31
Murray, N., & Chiang, J. 1997, ApJ, 454, L105
Nieto, J.­L., Roques, S., Llebaria, A., Vanderriest, C., Lelievre, G.,
Alighieri, S., Macchetto, F. D., & Perryman, M. A. C. 1988, ApJ, 325,
644
O'Flaherty, K. S., & Jakobsen, P. 1997, ApJ, 479, 673
Osterbrock, D. E. 1993, ApJ, 404, 551
Patnaik, A. R., Schneider, P., & Narayan, R. 1996, MNRAS, 281, 17
Pello, R., Miralles, J. M., Le Borgne, J.­F., Picat, J.­P., Soucail, G., &
Bruzual, G. 1996, A&A, 314, 73
Peng, C. Y., et al. 1999, ApJ, 524, 572
Pickering, T. E., Impey, C. D., & Foltz, C. B. 1994, AJ, 108, 1542
Postman, M., Lubin, L. M., Gunn, J. E., Oke, J. B., Hoessel, J. G.,
Schneider, D, P., & Christensen, J. A. 1996, AJ, 111, 615
Richards, G. T., et al. 2001, AJ, 121, 2308
Rusin, D. 2002, ApJ, in press (astro­ph=0202360)
Schmidt, M., & Green, R. F. 1983, ApJ, 269, 352
Schneider, P., Ehlers, J., & Falco, E. E. 1992, Gravitational Lenses XIV
(Berlin: Springer)
Schneider, P., & Wambsganss, J. 1990, A&A, 237, 42
No. 2, 2002 QSO PAIR Q2345+007 731

Small, T. A., Sargent, W. L. W., & Steidel, C. C. 1997, AJ, 114, 2254
Sol, H., Vanderriest, C., Schneider, J., Lelievre, G., & Pedersen, H. 1984,
A&A, 132, 105
Sprayberry, D., Bernstein, G. M., Impey, C. D., & Bothun, G. D. 1995,
ApJ, 438, 72
Steidel, C. C., & Sargent, W. L. W. 1991, AJ, 102, 1610 (SS91)
---------. 1993, in Proc. 3d Teton Astronomy Conf., The Environment and
Evolution of Galaxies, ed. J. M. Shull & H. A. Thronson (Dordrecht:
Kluwer), 263
Tody D., et al. 1986, IRAF Users' Manual (Tucson: NOAO)
Tyson, J. A., Seitzer, P., Weymann, R. J., & Foltz, C. 1986, AJ, 91, 1274
Umetsu, K., & Futamase, T. 2000, ApJ, 539, L5
Valdes, F. 1998, in ASP Conf. Ser. 145, Astronomical Data Analysis Soft­
ware and Systems VII, ed. R. Albrecht, R. N. Hook, & R. A. Bushouse
(San Francisco: ASP), 53
Valdes, F., & Tody, G. 1998, Proc. SPIE, 3355, 497
Vikhlinin, A., Forman, W., & Jones, C. 1999, ApJ, 525, 47
Weir, N., & Djorgovski, S. 1991, AJ, 101, 66
Wittman, D., Tyson, J. A., Margoniner, V. E., Cohen, J. G., &
Dell'Antonio, I. P. 2001, ApJ, 557, L89
Xanthopoulos, E., et al. 1998, MNRAS, 300, 649
Yuan, W., Brinkmann, W., Siebert, J., & Voges, W. 1998, A&A, 330, 108
Note added in proof.---As this article went to press, M. R. Meteger and T. A. Small (2002, private communication) commu­
nicated their finding of a significant broad absorption line (BAL) on the blue wing of C iv in component B from a 1998 Keck
spectrum that reached further to the blue than the previous (SS91) spectrum. Since the BAL QSOs are X­ray weak (P. J.
Green, T. L. Aldcroft, S. Mathur, B. J. Wilkes, &M. Elvis, ApJ, 558, 109 [2001]), this finding of strong intrinsic absorption in
B can account for its X­ray weakness and adds further evidence against the lensing hypothesis.
732 GREEN ET AL.