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CHANDRA SPECTRA OF THE SOFT X­RAY DIFFUSE BACKGROUND
M. Markevitch, M. W. Bautz, 1 B. Biller, Y. Butt, R. Edgar, T. Gaetz, G. Garmire, 2 C. E. Grant, 1
P. Green, M. Juda, P. P. Plucinsky, D. Schwartz, R. Smith, A. Vikhlinin, S. Virani,
B. J. Wargelin, and S. Wolk
Harvard­Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138; maxim@head­cfa.harvard.edu
Received 2002 June 25; accepted 2002 October 2
ABSTRACT
We present an exploratory Chandra ACIS­S3 study of the di#use component of the cosmic X­ray
background (CXB) in the 0.3--7 keV band for four directions at high Galactic latitudes, with emphasis on
details of the ACIS instrumental background modeling. Observations of the dark Moon are used to model
the detector background. A comparison of the Moon data and the data obtained with ACIS stowed outside
the focal area showed that the dark Moon does not emit significantly in our band. Point sources down to
3 # 10 #16 ergs s #1 cm #2 in the 0.5--2 keV band are excluded in our two deepest observations. We estimate the
contribution of fainter, undetected sources to be less than 20% of the remaining CXB flux in this band in all
four pointings. In the 0.3--1 keV band, the di#use signal varies strongly from field to field and contributes
between 55% and 90% of the total CXB signal. It is dominated by emission lines that can be modeled by a
kT ¼ 0:1 0:4 keV plasma. In particular, the two fields located away from bright Galactic features show a
prominent line blend at E # 580 eV (O vii+O viii) and a possible line feature at E # 300 eV. The two
pointings toward the North Polar Spur exhibit a brighter O blend and additional bright lines at 730--830 eV
(Fe xvii). We measure the total 1--2 keV flux of 1:0 1:2 # 0:2
Ï ÷ # 10 #15 ergs s #1 cm #2 arcmin #2 (mostly
resolved) and the 2--7 keV flux of 4:0 4:5 # 1:5
Ï ÷ # 10 #15 ergs s #1 cm #2 arcmin #2 . At E > 2 keV, the di#use
emission is consistent with zero, to an accuracy limited by the short Moon exposure and systematic
uncertainties of the S3 background. Assuming Galactic or local origin of the line emission, we put an upper
limit of #3 # 10 #15 ergs s #1 cm #2 arcmin #2 on the 0.3--1 keV extragalactic di#use flux.
Subject headings: intergalactic medium --- ISM: general --- methods: data analysis ---
X­rays: di#use background --- X­rays: ISM
1. INTRODUCTION
The existence of a cosmic X­ray background (CXB) was
one of the first discoveries of extrasolar X­ray astronomy
(Giacconi et al. 1962). In the intervening four decades,
observations with improving angular and spectral resolu­
tion have enhanced our understanding of the components
that make up this background. Several broadband, all­sky
surveys have been performed using proportional counter
detectors (Marshall et al. 1980; McCammon et al. 1983;
Marshall & Clark 1984; Garmire et al. 1992; Snowden et
al. 1995, 1997; for a review of pre­ROSAT results, see
McCammon & Sanders 1990). These surveys form a consis­
tent picture of the angular distribution of X­ray emission in
the various bands. Above 2 keV, the emission is highly iso­
tropic on large angular scales and has an extragalactic ori­
gin. Below 2 keV, the X­ray background is a mixture of
Galactic di#use emission (see, e.g., Kuntz & Snowden 2000
and references therein), heliospheric and geocoronal di#use
components (see, e.g., Cravens 2000), and extragalactic flux
from point sources and, possibly, from intergalactic warm
gas that may contain the bulk of the present­day baryons
(see, e.g., Cen &Ostriker 1999).
The earliest observations could provide only limited spec­
tral information on the background. Marshall et al. (1980)
found that the spectrum in the 3--50 keV range was well fit­
ted by a thermal bremsstrahlung model with kT # 40 keV.
In the 3--10 keV band, this can be approximated by a power
law with a photon index of #1.4. At energies below 1 keV,
the background surface brightness exceeds the extrapola­
tion of this power law (Bunner et al. 1969). Later observa­
tions with gas scintillation proportional counters and solid­
state detectors (Inoue et al. 1979; Schnopper et al. 1982;
Rocchia et al. 1984) suggested emission lines in the 0.5--
1.0 keV band, most likely from oxygen. The evidence for
emission lines in this band has become more convincing in
recent observations using CCDs (Gendreau et al. 1995;
Mendenhall & Burrows 2001) and in a high­resolution spec­
trum obtained by McCammon et al. (2002) in a microcalori­
metric experiment. A definitive demonstration of spectral
lines in the 0.15--0.3 keV band (at low Galactic latitude) was
obtained by Sanders et al. (2001), using a Bragg crystal spec­
trometer. Observations that combine high spectral and
angular resolution are essential for disentangling the many
CXB soft emission components.
Chandra and XMM should soon provide a wealth of
new information on the CXB. Several works have already
taken advantage of the Chandra's arcsecond resolution to
study the point­source component of the CXB (see, e.g.,
Mushotzky et al. 2000; Brandt et al. 2001; Rosati et al.
2002). The first XMM results are starting to appear as well
(De Luca & Molendi 2002; Warwick 2002; Lumb et al.
2002), utilizing the large e#ective area of that observatory.
In this paper, we present a Chandra ACIS study of the dif­
fuse CXB at high Galactic latitudes. The main advantage of
Chandra over all other instruments is its ability to resolve
point sources down to very low fluxes and probe the true dif­
1 Center for Space Research, Massachusetts Institute of Technology,
Cambridge, MA 02139.
2 Department of Astronomy and Astrophysics, Pennsylvania State
University, 525 Davey Laboratory, University Park, PA 16802.
The Astrophysical Journal, 583:70--84, 2003 January 20
# 2003. The American Astronomical Society. All rights reserved. Printed in U.S.A.
70

fuse background. In addition to that, compared to ROSAT
PSPC (which had lower detector background), ACIS has
energy resolution su#cient to identify spectral lines. Com­
pared to XMM EPIC, ACIS appears to be less a#ected by
instrumental background flares (although during quiescent
periods, the ACIS detector background per unit sky signal
is higher). To the extent that results can be compared,
we confirm many of the recent findings made with other
instruments.
Technical aspects of our study, especially the ACIS
instrumental background modeling, are quite complex, and
we discuss them here in detail. Much of our analysis proce­
dure may be useful for studies of extended sources such as
clusters of galaxies. Uncertainties are 1 # unless specified
otherwise.
2. DATA SET
For this exploratory study of the di#use CXB, we selected
four Chandra ACIS­S observations at high Galactic lati­
tudes, listed in Table 1. Our main focus is two relatively deep
(90--100 ks) observations, observation IDs (ObsIDs) 3013
and 3419, obtained at positions away from any bright
Galactic features seen in the RASS (ROSAT All­Sky
Survey) R4--R5 ( 3
4 keV) band (Snowden et al. 1997). For
comparison, we also analyze two shorter archival observa­
tions toward an edge (ObsID 869) and the middle (ObsID
930) of the North Polar Spur, 3 which exhibits bright emis­
sion in the ROSAT R4--R5 band. In Figure 1, positions of
the four observations are overlaid on the RASS R4--R5
map. None of these observations' original goals were related
to CXB, and there are no nearby cataloged extended sour­
ces, except an irregular galaxy that was a target of ObsID
869 and is spatially excluded from our analysis.
3. ACIS INSTRUMENTAL BACKGROUND
A critical part of this study is modeling of the ACIS
instrumental background. This background is caused by
cosmic charged particles and consists of a slowly changing
quiescent component and at least two species of highly vari­
able background flares, whose spectra are very di#erent
from that of the quiescent component. Below we describe in
detail how these components were dealt with from a practi­
cal perspective; their exact physical nature is beyond the
scope of this paper.
3.1. Background Flare Filtering
We use data from the ACIS back­side--illuminated (BI)
chip S3. Compared to the ACIS front­side--illuminated (FI)
chips, S3 has a higher sensitivity at low energies. However,
its sensitivity to low­energy X­rays also renders it more sen­
sitive to particle events, which results in more frequent
background flares than in FI chips (Plucinsky & Virani
2000; Markevitch 2001 4 ). The quiescent background is
stable and predictable; therefore, when the accuracy of the
background subtraction is critical, it is best to exclude flare
periods from the analysis. The spectra of the flaring and qui­
escent background components (discussed below) are such
that the best energy band in which to look for flares in chip
S3 is approximately 2.5--7 keV. Figure 2 shows light curves
in this energy band for our four observations (from the
whole S3 chip, excluding celestial sources). The time bin size
(#1 ks) is chosen to limit the statistical scatter while provid­
ing a reasonably detailed light curve.
In observations 3419, 869, and 930, the quiescent rate is
easily identifiable and very close to that in most other obser­
vations performed during 2000--2001. To limit the back­
ground modeling uncertainty, we exclude from further
analysis all time bins above and below a factor of 1.2 of this
rate (the rate can be lower, for example, because of occa­
sional short intervals of missing telemetry and bad aspect,
etc.). The resulting clean exposures are given in Table 1.
3.1.1. Anomalous Background in Field 3013
Observation 3013 is unusual in that the apparent `` quies­
cent '' rate between the numerous flare intervals (Fig. 2) is
about 30% higher than in other observations. It is also more
variable than the quiescent rate usually is. It appears that, in
fact, all of this exposure is a#ected by a long flare, which
requires special treatment (and, unfortunately, will add to
the systematic uncertainty of the results).
Again, we exclude the time periods above a factor of 1.2
of the apparent quiescent rate. Assuming that the back­
ground excess in the rest of the exposure is indeed a flare, we
can try to model it by taking advantage of the empirical
finding, based on a number of observations, that the spec­
tral shape of the most frequent, `` soft '' species of the BI
flares stays the same even while their flux varies strongly in
time. To illustrate this, we derive spectra of the background
flare components from the rejected, high count rate periods
of observations 3013 and 930. We first check the light curves
of FI chips also used in these observations, in order to
exclude flares of a di#erent, `` hard '' species that is less fre­
quent, a#ects both the BI and FI chips, and has a di#erent
spectrum. There are 1--2 ks of such flares within the already
excluded time periods in each observation. After excluding
those, we extract the spectra from the remaining high­rate
periods and subtract from them the spectra of the quiescent
periods of the same observations, normalizing them by
the respective exposure ratios (thereby subtracting all
3 The North Polar Spur, part of the Loop I supershell of emission in the
RASS (Egger & Aschenbach 1995), is thought to be the collision of explo­
sive remnants with the Local Bubble.
4 For Markevitch (2001), see http://asc.harvard.edu/cal, sections
`` ACIS,'' `` Background,'' `` General discussion.''
Fig. 1.---ROSAT PSPC all­sky map of CXB in the R4--R5, or 3
4 keV,
band (Snowden et al. 1997), in Galactic coordinates. Positions of our obser­
vations are marked; labels give ObsIDs.
DIFFUSE SOFT X­RAY BACKGROUND 71

time­independent sky and instrumental background com­
ponents). The resulting flare spectra are shown in Figure 3.
In a similar manner, we also extracted flare spectra from
several other archival observations (ObsIDs 766, 326, 2206,
2076, 2213, and 1934) spanning a period 2000 February--
2001 September and a range of flare intensities. Flares in all
those observations can be described by a model consisting
of a power law with a photon index of #0.15 and an expo­
nential cuto# at 5.6 keV, without the application of the tele­
scope e#ective area and the CCD quantum e#ciency
(commands arf none or model/b in XSPEC). In Figure 3,
this model is overplotted on the flare spectra from our CXB
fields 930 and 3013, fixing the spectral shape and fitting only
the normalization. As the figure shows, flares in both our
observations have nearly the same spectral shape, despite
their time di#erence of 1.5 yr and very di#erent intensities.
3013 3419
869 930
Fig. 2.---Light curves of the four CXB observations in the 2.5--7 keV band, where the contribution of flares is most easily detected. ObsIDs are marked in
each panel. For ObsID 869, only 72% of the chip area is used. Very Faint mode filtering was applied (x 3.3). Shaded bins are above or below a factor of 1.2 of
the quiescent level and are excluded.
TABLE 1
Data Summary
Parameter ObsID 3013 ObsID 3419 ObsID 869 ObsID 930 Moon
l; b
Ï ÷ (deg).......................................................................... (259.6, +56.9) (187.1, #31.0) (36.6, +53.0) (358.7, +64.8) . . .
Galactic N H (#10 20 cm #2 ) .................................................. 4.1 11.4 4.3 1.8 . . .
ROSATR4--R5 flux (#10 #6 counts s #1 arcmin #2 ) .............. 160 90 200--250 a 400 . . .
Observation date................................................................ 2001 Dec 13 2002 Jan 8 2000 Jun 24 2000 Apr 19 2001 Jul 26
Total (uncleaned) exposure (ks) ......................................... 112 98 57 40 16
Exposure for source detection (ks) ..................................... 101 92 52 28 . . .
Exposure for CXB spectra (ks) ........................................... 69 86 52 20 11
Field solid angle (arcmin 2 ).................................................. 69 69 51 69 70
a A#ected by an artifact in the ROSAT all­sky map.
72 MARKEVITCH ET AL. Vol. 583

Freeing the spectral shape parameters, we obtained a
photon index of #0.10 # 0.07 and a cuto# at 5:2 # 0:7 keV
for observation 930 and #0.1 # 0.3 and 7:2 ×1
#3:4 keV for
3013, consistent with the above fit for the composite flare
spectrum.
Therefore, if the background excess a#ecting the useful
period of observation 3013 is indeed a residual flare, we can
expect that it has the same spectral shape, only a still lower
normalization. As is seen in x 6, its spectrum is indeed con­
sistent with this assumption at the energies for which the
comparison is possible. To try modeling the residual back­
ground excess in 3013, we chose to fix the shape parameters
to the best­fit values from this particular observation (given
above), even though they are less strongly constrained than
those from the composite spectrum, to account for any pos­
sible slow evolution of the spectrum, which cannot be ruled
out with the data at hand. This choice has no significant
e#ect on our results. The normalization of this flare model is
determined in x 6, after subtraction of the quiescent back­
ground component and removal of point sources.
For comparison, Figure 3 also shows a quiescent instru­
mental background spectrum from the dark­Moon observa­
tions (discussed below). We note that the BI soft flare
spectrum at high energies is much softer than the quiescent
background; unless the flare is very strong, its contribution
above 10 keV is unnoticeable (this is not true for the other,
hard flare species mentioned above). This fact is used below.
3.2. Quiescent Background
3.2.1. Dark­Moon Observations
To separate the CXB and instrumental components of
the ACIS background, Chandra observed the dark Moon in
2001 July in a series of six short pointings (ObsIDs 2469,
2487, 2488, 2489, 2490, and 2493), tracking the Moon for a
total of about 15 ks. ACIS chips S2, S3, I2, and I3 were on
and telemetered data in Very Faint (VF) mode. A second
installment of Moon observations in 2001 September
exposed chips I2 and I3. Technical di#culties encountered
in these two runs, related mostly to the fact that the Chandra
aspect camera cannot be used near the Moon, prevented
further dark­Moon observations. Here we use only the S3
data from 2001 July. During that run, optical flux from the 1
3
of the lunar disk that was illuminated was imaged onto to
the ACIS focal plane. The ACIS optical blocking filters
were not designed to reject visible light from the sunlit
Moon, and a detectable o#set signal (bias error) was pro­
duced. The e#ect was most severe in chip I2, but very small
in chip S3. A correction to each event's pulse­height ampli­
tude (PHA) was calculated individually by averaging the
lowest 16 pixels of the 5 # 5 pixel VF mode event island.
This o#set in S3 was well below the threshold of a#ecting
the event grades, so this problem did not result in any loss of
events due to the onboard grade rejection (as was the prob­
lem for chip I2 in this data set). The average correction to
the energy for events in chip S3 was within a few eV, negli­
gible for our purposes.
A 2.5--7 keV S3 light curve for all 2001 July dark­Moon
observations is shown in Figure 4. The end of the exposure
was a#ected by an apparent faint flare, which was filtered
using the same factor of 1.2 threshold as in x 3.1 (we had to
use smaller time bins, which resulted in some statistical devi­
ations that were also excluded for consistency). The result­
ing clean Moon exposure for S3 is 11,400 s. At high
energies, where the CXB contribution is negligible, the
Moon quiescent background rate was within a few percent
of that in other recent observations. Scientific results from
the Moon observations will be discussed by C. E. Grant et
al. (2003, in preparation).
3.2.2. Event Histogram Mode Data
An independent approach to calibrating the ACIS instru­
mental background utilizes the event histogram mode
(EHM) data (Biller, Plucinsky, & Edgar 2002). 5 These data
are collected during science observations by the HRC­I
detector while ACIS is stowed inside the detector support
structure. This structure blocks celestial X­rays but does not
a#ect the particle rate significantly (as is seen below). In this
mode, the telemetry capacity available for ACIS is small, so
the only information transmitted is a PHA histogram for
the events from a predefined region of the chip. For this
reason, the usual exclusion of bad pixels and the position­
dependent gain correction cannot be applied. At this
location, ACIS is also faintly illuminated by the internal
Fig. 3.---Spectra of the excluded background flares (see text) in ObsIDs
930 (red ) and 3013 (black), compared to the quiescent spectrum from the
dark Moon (blue), for the whole S3 chip. Red histograms show a model that
was fitted to a combination of other observations with flares and renormal­
ized, without a change of shape, to match the ObsID 930 and 3013 spectra.
Fig. 4.---Same as Fig. 2, but for the 2.5--7 keV light curve for the dark­
Moon observations. The rate is shown before VFmode filtering (x 3.3).
5 See http://asc.harvard.edu/cal, sections `` ACIS,'' `` Background,''
`` Event Histogram mode.''
No. 1, 2003 DIFFUSE SOFT X­RAY BACKGROUND 73

calibration line source. The flare component discussed in
x 3.1 is never observed in the stowed position.
An EHM spectrum from the whole S3 chip accumulated
over the 2001 July--October period (straddling the date of
the Moon observation) is shown in Figure 5 (see also Biller
et al. 2002). For comparison, we overlaid a dark­Moon
spectrum. In order for it to be directly comparable to the
EHM spectrum, in deriving it we did not exclude bad CCD
pixels, apply gain corrections (that is, we used PHA rather
than pulse height invariant, or PI, values), or apply the addi­
tional VFmode filtering (x 3.3).
Figure 5 shows that away from the calibration source
lines, the agreement between the spectra is quite remarkable
over the entire energy range, within the statistical accuracy
of the Moon data set. One might expect both the dark­
Moon and the EHM spectra to exhibit emission above the
non--cosmic background level seen in ordinary observa­
tions; e.g., ROSAT and ASCA Moon data suggested emis­
sion at low energies (Schmitt et al. 1991; Kamata et al.
1999), and the detector support structure may be radioac­
tive. In principle, one could also imagine a component of
the quiescent particle background that can be blocked by
that structure. While the coincidence of the two spectra in
Figure 5 does not rule out a conspiracy of these possibilities,
it makes each of them very unlikely. It also supports the
conclusion by Freyberg (1998) that the ROSAT emission in
the direction of the dark Moon was actually fluorescent
emission from a region around the Earth (i.e., below the
Chandra orbit). Therefore, we assume that both the dark­
Moon and the EHM data give the true quiescent back­
ground, and we use the Moon data set as an instrumental
background model for the sky data below.
The EHM data were also used directly as the background
model for a CXB study by R. J. Edgar et al. (2003, in prepa­
ration). Another measure of the ACIS background was
obtained in 1999 August, just prior to opening the Chandra
telescope door. That data set is analyzed in Bagano# (1999),
which may be consulted for background line identifications
and other qualitative information. 6 Unfortunately, those
data cannot be used directly for recent observations,
because both the background and the ACIS detector have
evolved significantly since then. Forthcoming calibration
observations with ACIS stowed and working in the full
imaging mode (as opposed to EHM), could be used for the
most recent observations.
3.2.3. Background Time Dependence
From the analysis of a large number of ACIS observa­
tions and monitoring of the rate of ACIS events rejected on
board (Grant, Bautz, & Virani 2002), we know that the
high­energy quiescent background slowly declined from
launch until around the end of 2000 and has been relatively
constant during 2001 (Fig. 6 shows this behavior in the 5--10
keV band, where the instrumental background dominates),
in anticorrelation with the solar cycle. At low energies, the
qualitative behavior is similar, but it is di#cult to tell exactly
because of the di#ering sky signal. In addition to the slow
evolution, there are small variations of the quiescent rate on
short timescales. Since the background can change between
the observations of our fields and the Moon, the Moon
background may need a correction, and the background
uncertainty must be included in the final results.
The EHM data are free from the sky signal and flares and
can be used for checking the quiescent background time
dependence, especially in the soft band most important for
the present study. From the comparison of EHM (as well as
blank­field) data in di#erent energy bands, it appears that
during the short­term quiescent background fluctuations,
such as those in Figure 6, its spectral shape does not change
significantly and only the normalization varies. This behav­
Fig. 5.---Spectra of dark­Moon (black) and EHM data (red ). This Moon spectrum is extracted in PHA channels in a special manner to be directly
comparable to the EHM data (see text). The energy scale is approximate (PHA values multiplied by 4.7 eV per channel). Apart from the faint internal
calibration source lines present in the EHM data (most notably at 5.9 and 1.5 keV), there is very good agreement between the spectra.
6 See http://asc.harvard.edu/cal, sections `` ACIS,'' `` Background,''
`` Measurement with closed mirror cover.''
74 MARKEVITCH ET AL. Vol. 583

ior enables us to account for these variations by normalizing
the model instrumental background by the ratio of rates at
high energies, e.g., 10--12 keV, where the contribution from
celestial sources and possible faint undetected flares (x 3.1.1)
is negligible.
Figure 7 shows EHM rates from the full S3 chip in the
0.5--2, 2--7, and 5--10 keV bands divided by the rates in the
PHA interval of 2500--3000 ADU (approximately 10--12
keV; PHA is preferred over energy because many science
observations use an onboard cuto# at 3000 ADU). The
period from 2001 July is shown because earlier EHM data
were collected from parts of the chip. The scatter of these
ratios is quite small; in all energy bands, the intrinsic scatter
around the mean required in addition to the Poissonian
scatter is 1.5%--2% (1 #) or less, a reduction from 3%--4%
for fluxes not normalized by the high­energy rate. One also
notices that neither the rates nor the spectral shape changed
systematically between the dates of the Moon observation
and our deep observations 3013 and 3419.
We use such high­energy rate matching in our CXB anal­
ysis by normalizing the Moon spectrum by the ratio of the
respective 2500--3000 ADU rates. For ObsIDs 3013, 3419,
930, and 869, such normalizations are 1.06, 0.97, 1.01, and
0.93 times the ratio of the Moon exposure to the respective
exposures, so this correction is, as expected, small. We
adopt a systematic uncertainty of the resulting quiescent
background normalization of 2% (1 #), as derived above.
We note that over longer periods, the spectral shape of
the detector background can change---the spectra of the
empty­field observations from 2000 and 2001 di#er by #5%
at energies where the noncosmic component dominates.
This means that the Moon spectrum with a simple normal­
ization adjustment might not be a good model for our
earlier observations 930 and 869. However, as seen below,
the soft di#use signal in those observations is so strong that
the background uncertainty does not matter.
3.3. VF Mode Background Filtering
In VF ACIS telemetry mode, the detector background
can be reduced significantly by rejecting events with signal
above the split threshold in any of the outer pixels of the
5 # 5 pixel event island, after an approximate correction for
the charge transfer ine#ciency. Details of the method can
be found in Vikhlinin (2001). 7 Figure 8 shows the e#ect of
such filtering on the spectra of the background (the dark
Moon) and real X­ray events from an extended celestial
source una#ected by photon pileup. It results in a significant
reduction of the background rate, especially at the lowest
and highest energies, while rejecting only about 2% of the
real events. The background reduction is stronger for FI
chips (not used in this work). All four of our CXB observa­
tions, as well as all Moon observations, were telemetered in
VFmode and filtered in this manner.
4. POINT SOURCES
We now detect and exclude point sources in our CXB
fields. All of the S3 chip is within 7 0 of the optical axis, and
the point­spread function (PSF) is narrow over the whole
field, so point­source detection is photon­limited and back­
ground is relatively unimportant for it. It is therefore advan­
tageous to include periods of moderately high background
if it significantly increases the exposure. For observations
3013, 3419, and 930, we applied a less restrictive light­curve
filtering (using the 0.3--10 keV band, and for 930 a higher
threshold factor of 2, instead of 1.2) to increase the expo­
sures for source detection (see Table 1).
Fig. 6.---Time dependence of the S3 chip 5--10 keV quiescent background
rate in various blank­field observations since launch (see Markevitch 2001).
Statistical errors are comparable to the symbol size. Red symbols denote
observations used in this work (ObsID 3013 is shown before the residual
flare correction).
Fig. 7.---Ratios of the S3 background rates in the 0.5--2, 2--7, and 5--
10 keV bands to the 2500--3000 ADU (#10--12 keV) rate, for EHM obser­
vations that used the full chip. Horizontal lines show average values.
7 See http://asc.harvard.edu/cal, sections `` ACIS,'' `` Background,''
`` VFmode.''
No. 1, 2003 DIFFUSE SOFT X­RAY BACKGROUND 75

In the three fields with pointlike original targets (ObsIDs
3013, 3419, and 930), we excluded r ¼ 30 00 circles around the
target from the analysis. In ObsID 869, whose original tar­
get is a galaxy, we masked that source liberally and used the
remaining 72% of the chip area.
Source detection was performed in two bands, 0.3--2 and
2--7 keV. The source candidates were identified in a standard
manner, by applying wavelet filtering to the image to reduce
statistical noise and then searching for local brightness max­
ima using the code of Vikhlinin et al. (1998). The source
positions were then refined by photon centroiding and the
source fluxes calculated within the 90% PSF--encircled
energy radii r 90 (and then divided by 0.9), using as the local
background a wavelet decomposition component contain­
ing details on the largest angular scale. For simplicity, we
set a relatively high, spatially uniform lower limit on the
source flux that corresponds approximately to 8--10 photons
anywhere in the field in each observation. This ensures that
all detected sources are real and that the background contri­
buti