Äîêóìåíò âçÿò èç êýøà ïîèñêîâîé ìàøèíû. Àäðåñ îðèãèíàëüíîãî äîêóìåíòà : http://hea-www.harvard.edu/~bradw/cv/papers/MBM12.ps.gz Äàòà èçìåíåíèÿ: Thu Apr 7 18:40:27 2005 Äàòà èíäåêñèðîâàíèÿ: Sat Dec 22 03:10:57 2007 Êîäèðîâêà: Ïîèñêîâûå ñëîâà: superbubble

CHANDRA OBSERVATIONS OF MBM 12 AND MODELS OF THE LOCAL BUBBLE
R. K. Smith
NASA Goddard Space Flight Center, Code 662, Greenbelt, MD 20771; and Department of Physics and Astronomy,
The Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218; rsmith@milkyway.gsfc.nasa.gov
R. J. Edgar, P. P. Plucinsky, and B. J. Wargelin
Harvard­Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138
P. E. Freeman
Department of Statistics, Carnegie Mellon University, 5000 Forbes Avenue, Pittsburgh, PA 15213
and
B. A. Biller
Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721
Received 2004 September 24; accepted 2005 January 3
ABSTRACT
Chandra observations toward the nearby molecular cloud MBM 12 show unexpectedly strong and nearly equal
foreground O viii and O vii emission. As the observed portion of MBM 12 is optically thick at these energies, the
emission lines must be formed nearby, coming from either the Local Bubble ( LB) or charge exchange with ions from
the Sun. Equilibrium models for the LB predict stronger O vii than O viii, so these results suggest that the LB is far
from equilibrium or that a substantial portion of O viii is from another source, such as charge exchange within the solar
system. Despite the likely contamination, we can combine our results with other EUVand X­ray observations to reject
LB models that posit a cool recombining plasma as the source of LB X­rays.
Subject headinggs: ISM: bubbles --- plasmas --- supernova remnants --- X­rays: ISM
Online material: color figure
1. INTRODUCTION
After decades of observation, the nature of the diffuse soft
(# 1
4 keV ) X­ray background is still mysterious. Early observa­
tions of soft X­ray emission toward many sight lines ( Bowyer
et al. 1968; Bunner et al. 1969; Davidsen et al. 1972) showed that
there is substantial diffuse emission (see also reviews by Tanaka
& Bleeker 1977; McCammon & Sanders 1990). Broadband
spectral data from proportional counter observations fit a three­
component model: an unabsorbed 10 6 K thermal component, an
absorbed 2 ; 10 6 K thermal component, and an absorbed power
law. The latter two components contribute mostly at higher en­
ergies, while most of the 1
4 keV band comes from the local (i.e.,
unabsorbed) 10 6 K thermal component.
Williamson et al. (1974) excluded on physical grounds all
then­proposed sources for the 1
4
keV band emission other than a
hot (#10 6 K) ionized plasma. Current theories to explain the
origin of this emission still require an ionized plasma, and in­
clude (1) a local young supernova explosion in a cavity (Cox &
Anderson 1982; Edgar & Cox 1993), (2) a series of supernovae
( Innes & Hartquist 1984; Smith & Cox 2001), or (3) an over­
ionized plasma slowly recombining after substantial adiabatic cool­
ing (Breitschwerdt &Schmutzler 1994; Frisch 1995; Breitschwerdt
et al. 1996). One additional possibility was first suggested by
Cox (1998), who pointed out that charge exchange from the solar
wind might create at least part of the 1
4 keV band emission.
Independent of the observed 1
4 keV band emission, absorption­
line measurements to many nearby stars (e.g., Welsh et al. 1990,
1991) show that we are surrounded by a irregularly shaped
``cavity'' with very low density. The standard model for the
Local Bubble ( LB) combines these two observations into the
``displacement'' model (Sanders et al. 1977; Snowden et al.
1990). In this picture, the diffuse X­rays come from an elongated
bubble of hot gas with average radius #100 pc, with the Sun near
the center. If the LB is filled with hot (#10 6 K) gas in collisional
ionization equilibrium (CIE), then the models of the resulting
emission fit the observed 1
4 keV band soft X­ray spectrum. How­
ever, this phenomenological model explains neither the origin
of the hot gas nor the low­density region.
Early LB models that attempted to explain both the hot gas
and the low­density region (Cox & Anderson 1982; Edgar &
Cox 1993) modeled it as an #10 5 yr old supernova remnant. How­
ever, this model predicts a large column density of O vi that is
simply not observed along many sight lines (Shelton &Cox1994;
Oegerle et al. 2005). Since every oxygen atom passing through
the blast wave needs to pass through this ionization state, the
model predictions are quite robust and are simply not observed.
In addition, the total thermal energy required by the phenom­
enological models is between 0:37 ; 10 51 and 1:1 ;10 51 ergs---
nearly the entire kinetic energy of a supernova. Smith & Cox
(2001) considered models with multiple supernovae that are
allowed by the O vi data and energy considerations and that
roughly fit the observed broadband emission. The recombining
plasma model, described in detail in Breitschwerdt et al. (1996),
has also been able to qualitatively fit existing observations. How­
ever, as discussed below, we are now able to reject the basic
recombining plasma model by combining our results with other
EUV and X­ray observations.
2. X­RAY EMISSION FROM THE LOCAL BUBBLE
The difficulty in measuring the soft X­ray spectrum has lim­
ited further analysis. A 10 6 K plasma with solar abundances, in
equilibrium or not, is line dominated in the range 0.1--1.0 keV,
but the first observation able to even partially resolve these lines
was done only recently with the Diffuse X­Ray Spectrometer
A
225
The Astrophysical Journal, 623:225--234, 2005 April 10
# 2005. The American Astronomical Society. All rights reserved. Printed in U.S.A.

(DXS; Sanders et al. 2001). Most of the lines are from L­shell ions
of elements from the third row of the periodic chart: Si, S, Mg
(and their neighbor Ne), and M­shell ions of Fe around 72 eV.
There should also be some X­ray emission from the K­shell lines
of carbon, nitrogen, and oxygen.
Fitting the spectrum requires a spectral emission code for
a collisional plasma, such as described in Smith et al. (2001) or
Kaastra & Jansen (1993). These codes model the thousands of
lines emitted by the cosmically abundant elements because of
collisional excitation. However, the lines in the observed spec­
trum are numerous, and the atomic data for the line emission are
incomplete and sometimes inaccurate. In any event, no calcu­
lated spectrum matches the high­resolution observations of the
soft X­ray spectrum, such as those from DXS, and it is not clear
whether the problem lies in the physical models or in the atomic
data (Sanders et al. 2001). To avoid these problems, unambig­
uous and strong emission lines are needed. In a 10 6 K plasma, the
most abundant line­emitting element is oxygen. In equilibrium at
10 6 K the dominant stage is helium­like O +6 , with trace amounts
of O +7 and O +5 . However, the hot gas in the LB need not be in
equilibrium. If (somehow) a recent (P10 5 yr) supernova created
the LB, then the gas would still be ionizing. Conversely, if the LB
is old (k10 6 yr), then it should be recombining. In either case, the
ratios of the O +6 and O +7 ion populations will not be at their
equilibrium values.
The strongest O vii and O viii emission lines are in the soft
X­ray range, between 0.5 and 0.8 keV. O vii has a triplet of lines
from n ¼ 2 ! 1, the so­called forbidden ( F ) line at 0.561 keV,
the resonance ( R) line at 0.574 keV, and the intercombination ( I )
line (actually two lines) at 0.569 keV, while O viii has a Ly#
transition at 0.654 keV. These lines are regularly used as plasma
diagnostics in other situations: Acton & Brown (1978) discussed
the nonequilibrium ionization effects on O vii lines in solar flares.
Vedder et al. (1986), using data from the Einstein Focal Plane
Crystal Spectrometer ( FPCS), measured the ionization state in
the Cygnus Loop using these lines. Gabriel et al. (1991) discussed
the general use of O vii and O viii emission lines in hot plasmas,
specifically applied to Einstein FPCS observations of the Puppis
remnant. The goal of our work is to apply these methods, already
used to model supernova remnants, to the specific case of the
LB.
Earlier observations suggested that oxygen emission lines
are strong in the LB. Inoue et al. (1979) detected the O vii F+
I+R emission lines with a gas scintillation proportional counter.
Schnopper et al. (1982) and Rocchia et al. (1984), using solid­
state Si(Li) detectors, detected O vii as well as lines from other
ions. More recently, a sounding rocket flight of the X­Ray Quan­
tum Calorimeter ( XQC) observed a 1 sr region of the sky at high
Galactic latitude between 0.1 and 1 keV with an energy resolu­
tion of 5--12 eV, and detected both O vii and O viii. The O vii flux
was 4:8 # 0:8 photons cm #2 s #1 sr #1 ( hereafter line units, LU ),
and the O viii flux was 1:6 # 0:4 LU ( McCammon et al. 2002).
XQC's large field of view means that the source of the photons
cannot be determined, but this does represent a useful upper limit
to the LB emission in this direction.
Measuring the emission coming solely from the LB re­
quires observations of clouds that shadow the external emission.
Snowden et al. (1993, hereafter SMV93) used ROSAT observa­
tions of the cloud MBM 12 to measure the 3
4 keVemission in the
LB and found a 2 # upper limit of 270 counts s #1 sr #1 in the
ROSAT 3
4 keV band. 1 This band includes both the O vii triplet
and the O viii Ly# line. They also fitted a ``standard'' 10 6 K CIE
LB model assuming a path length of #65 pc and found a good
match to the observed 1
4
keVemission with an emission measure
of 0.0024 cm #6 pc. This would generate only #47 counts s #1
sr #1 in the 3
4
keV band, well within the 2 # limit. For compari­
son, this model predicts 0.28 LU of O vii line emission, which is
the dominant contribution to the 3
4 keV band emission.
The ROSAT PSPC had insufficient spectral resolution to
separate the O vii and O viii emission lines from each other, the
background continuum, and possible Fe L line emission. We
therefore used the Chandra ACIS instrument to redo the SMV93
observations with higher spectral and spatial resolution. Our re­
sults were unfortunately affected by a large solar flare during part
of the observation, the charge transfer inefficiency degradation
experienced by the ACIS detectors early in the Chandra mission,
and the somewhat higher than expected background. Despite
these issues, we were able to clearly detect strong O vii and O viii
emission lines.
3. OBSERVATIONS AND DATA ANALYSIS
MBM 12 is a nearby high­latitude molecular cloud (l; b ¼
159N2; #34 # ). The distance to MBM 12 is somewhat controver­
sial. Hobbs et al. (1986) placed it at #65 pc based on absorption­
line studies; this is revised to 60 # 30 pc using the new Hipparcos
distances for their stars. However, Luhman (2001), using infra­
red photometry and extinction techniques, found a substantially
higher distance of 275 # 65 pc. Andersson et al. (2002) found a
similar distance (360 # 30 pc) for most of the extincting mate­
rial, but also found evidence for some material at #80 pc. As part
of a much larger survey of Na i absorption toward nearby stars,
Lallement et al. (2003) found foreground dense gas toward stars
90--150 pc distant in the direction of MBM 12. They also sug­
gested that the distance discrepancy could be resolved if the
molecular cloud MBM 12 is itself behind a nearby dense H i
cloud. However, since we use it as an optically thick shadowing
target, so long as MBM 12 is nearer than the nonlocal Galactic
sources of 3
4 keV band X­rays that contribute to the diffuse back­
ground, the precise distance is unimportant. Indeed, the greater
distance is interesting because it would constrain possible sources
of diffuse 3
4 keV band X­rays.
Our Chandra observations of MBM 12 were separated into
two nearly equal sections. The first observation ( hereafter Obi0)
was performed on 2000 July 9--10. This observation was cut
short by a severe solar flare that led to an automatic shutdown.
Preshutdown, the flare apparently also created a systematically
higher background in the ACIS during the entire observation,
which unfortunately meant the entire observation had to be ex­
cluded (see x 3.1). The second observation ( hereafter Obi1), was
performed on 2000 August 17 for #56 ks during a time of lower
solar activity. The pointing direction for both observations was
02 h 55 m 50. s 2, 19 # 30 0 14B0 (see Fig. 1). The primary CCD, ACIS­
S3, was placed on the peak of the 100 #m IRAS emission, which
we assume corresponds to the densest part of the cloud.
MBM 12 is a relatively thin molecular cloud, and so we must
consider whether it is optically thick to background X­rays even
at the relatively soft energies of O vii (#0.57 keV ) and O viii
(0.654 keV ). Assuming solar abundances, the cross sections at
these energies are 9:4 ; 10 #22 and 6:7 ; 10 #22 cm #2 per H atom,
respectively ( Balucinska­Church & McCammon 1992). Mea­
suring the column density over the face of the entire cloud is dif­
ficult. Absorption­line measurements show the column density
at one position; Luhman (2001) found that most stars in MBM 12
have A V < 2, with background stars in the range A V ¼ 3 8.
Converting this to an equivalent hydrogen column density via
1 For consistency we present all surface brightnesses in units of steradians;
note that 1 sr ¼ 1:18 ; 10 7 arcmin 2 .
SMITH ET AL.
226 Vol. 623

NH /A V ¼ 1:9 ; 10 21 cm #2 mag #1 (Seward 2000) gives NH <
3:8 ; 10 21 cm #2 for much of the cloud, with a maximum value
in the range NH ¼ (6 15) ; 10 21 cm #2 . SMV93 also measured
the column density by cross­correlating the ROSAT 3
4 keV and
IRAS 100 #m flux, and derived a value of N H /F(100 #m) ¼
1:3 ×1:1
#0:1 ; 10 20 cm #2 ( MJy sr #1 ) #1 . For these observations, the
central ACIS­S3 detector was positioned on the brightest in­
frared position of MBM 12, where F(100 #m) ¼ 30 35 MJy
sr #1 , so using the SMV93 ratio, NH ¼ (3:6 8:4) ; 10 21 cm #2 .
Since we are observing in the direction of the densest part of
MBM 12, we believe that NH ¼ 4 ; 10 21 cm #2 is a conservative
value. The opacity of MBM 12 to background line emission is
then #(O vii) ¼ 3:76 and #(O viii) ¼ 2:68, and any distant emis­
sion will be reduced by 98% at O vii and 93% at O viii. Only an
extremely bright background source could affect our results, and
this is unlikely. The ROSAT PSPC had a large 2 # field of view,
and although SMV93 found some 3
4
keV emission off­cloud, it
was only #3.3 times brighter than their 2 # upper limit for the
3
4
keV toward the cloud.
3.1. Data Processing
The data reduction process was somewhat unusual, since
emission from the LB completely fills the field of view, and the
features of interest are extremely weak. Our original plan was to
use the front­illuminated ( FI ) CCDs to measure the individual
lines, and the back­illuminated ( BI ) CCDs (which have more
effective area but lower resolution and higher background ) to
confirm the result. Based on the Gendreau et al. (1995) measure­
ment of the soft X­ray background, we expected prelaunch to
obtain at most #0.004 counts s #1 from O vii and #0.001 counts
s #1 from O viii lines from both the four ( FI ) ACIS­I CCDs and
the ( BI ) ACIS S3 CCD. However, after the proton damage to the
FI CCDs early in the mission, the highly row dependent response
of the FI CCDs means that these lines would be very difficult to
extract robustly. Therefore, we focused on data from the BI CCD
ACIS­S3.
We used CIAO, version 3.1, tools to process the observations,
along with CALDB, version 2.27. Our result depends crucially
on the background measurement, which must be done indirectly
since the source fills the field of view. The two most important
steps are removing flares and point sources. We followed a pro­
cedure similar to that described in Markevitch et al. (2003), al­
though as the data were taken in FAINT mode, VFAINT filtering
was not possible. For the purposes of source­finding only, we
merged the two observations using the -->align _ evt routine. 2 We
then ran -->
celldetect, requiring a signal­to­noise ratio (S / N) > 3,
which found 16 sources (including the well­known source XY
Ari). We excluded all these sources, using circles of 15 00 radius
(except XY Ari, where we used a 30 00 radius circle). This proce­
dure eliminated 15% of all events, with XYAri alone accounting
for 11%, and removed 5.0% of the area, leaving a total field of
view of 67 arcmin 2 . We then made a light curve of the ACIS­S3
events between 2.5 and 7 keV for each observation, as shown
in Figure 2. We constructed a histogram of the observed rates
and fitted it with a Gaussian, representing the quiescent rate, plus
a constant to roughly model the flares. We obtained good fits
in both cases, which showed that the quiescent level in the first
observation was 0.22 counts s #1 , while in the second it was
0.16 counts s #1 . The quiescent rate in the first observation is
significantly above those shown in Markevitch et al. (2003),
which ranged from 0.1 to 0.16 counts s #1 . These fluctuations in
the background counting rate are likely due to protons that are
trapped in the Earth's radiation belts or directly from solar flares.
Following Markevitch et al. (2003), we removed all times
when the count rate was not within 20% of the average rate for
the second observation. The maximum allowed rate is therefore
0.192 counts s #1 , which limits the first observation to 2.2 ks of
``good time,'' too little data to be useful. This exclusion was not
capricious. Although the first MBM 12 observation was #40%
brighter than the average background rate for ACIS­S3, we ex­
pended substantial effort attempting to model it. All our attempts
required adding what amounted to arbitrary terms at the same
Fig. 2.---Left: Light curve from the first MBM 12 observation between 2.5 and 7 keV, with a dashed line showing our best­fit quiescent rate of 0.22 counts s #1 .
Right: Same, from the second observation, with a quiescent rate of 0.16 counts s #1 .
Fig. 1.---IRAS 100 #m image of MBM 12, with the Chandra observation
coordinates overlaid. [See the electronic edition of the Journal for a color version
of this figure.]
2 See http://asc.harvard.edu /cont­soft /software/align _ evt.1.6.html.
OBSERVATIONS OF MBM 12 AND LOCAL BUBBLE MODELS 227
No. 1, 2005

energies as our emission lines, so no conclusive results were pos­
sible. Therefore, we reluctantly excluded it from the remaining
analysis. After filtering the second observation, we were left with
38.2 ks of usable data.
To cross­check these results, we followed Markevitch et al.
(2003) and compared the total counts in selected energy bands to
the total counts seen with pulse­height amplitude ( PHA) values
between 2500 and 3000. PHA values in this range correspond to
energies above 10 keV, which are outside Chandra's range and
thus effectively measure the ``particle background'' rate. Table 1
shows the results for our MBM 12 observations and one of our
background data sets, ObsID 62850 (see x 3.2). In Table 1 we
show the results from the filtered Obi1 data as well as the 2.2 ks
of ``good'' data from Obi0 and values from all of Obi0. Ratios
of the number of events in these bands to the number of events
with PHA values between 2500 and 3000 are given as well.
The unfiltered Obi0 data show very high ratios, suggesting con­
tamination from some source besides the normal particle back­
ground. Conversely, the Obi1 observations of MBM 12 show
slightly higher ratios compared to the ObsID 62850 ratios, con­
sistent with a weak source such as the diffuse X­ray background.
The ratios seen for ObsID 62850 are similar to the average values
observed by Markevitch et al. (2003) for the EHM observations.
3.2. Background
Since the LB fills the field of view in all normal Chandra ob­
servations, determining the true background is nontrivial. There
are three types of Chandra observations that contain only in­
strumental or near­Earth background: the Dark Moon observa­
tion described in Markevitch et al. (2003), the event histogram
mode ( EHM ) data taken by ACIS during HRC­I observations
and also described in Markevitch et al. (2003), and the ``stowed
ACIS'' observations first described in Wargelin et al. (2004).
These latter observations (to date, ObsID 4286, 62846, 62848,
and 62850) were done with ACIS clocking the CCDs in ``timed
exposure'' mode and reporting event data in the VFAINT te­
lemetry format, and with ACIS in a stowed position at which
ACIS receives negligible flux from the telescope and the on­
board calibration source. We used data from ObsIDs 62848,
62848, and 62850 as our background data sets. Together these
data sets represent 144 ks of observations, although they were
treated independently in our fits. ObsID 4286 is less than 10 ks
and was not used. Between our observations and the last of these
observations (in 2003 December) there was little change in the
overall background rate. 3
We fitted the MBM 12 data from Obi1 simultaneously with
stowed ACIS data using both a foreground and background
model and -->Sherpa's CSTAT statistic. The fits were restricted to
the range 0.4--6 keV, since above 6 keV the background ap­
pears to rise because of the near­constant cosmic­ray background
combined with the falling mirror response, while below 0.4 keV
the background rises and falls in a highly variable fashion. Our
foreground model consisted of two unabsorbed delta functions
for the O vii and O viii lines, plus an absorbed power law and
thermal component to represent the cosmic X­ray background
and the distant hot Galactic component. We did not include a
thermal component with T #10 6 K to represent the LB emis­
sion, as nearly all of that emission would be below 0.4 keV,
except for the oxygen lines. The absorption was allowed to vary
between NH ¼ 3:6 and 8:4 ; 10 21 cm 2 for both components;
the best­fit value was 6 ; 10 21 cm 2 , although our results were
not particularly sensitive to this value. We used the values from
Lumb et al. (2002) for the power­law component ( # ¼ 1:42 #
0:03, normalized to 8:44 ×2:55
#0:23 photons cm #2 s #1 keV #1 at 1keV ),
as well as for the temperature of the thermal component (0:2 #
0:01 keV ). The normalization on the thermal component was
allowed to vary, since significant variation has been seen for this
absorbed hot gas ( Kuntz & Snowden 2000). The O viii Ly# line
position was set at 0.654 keV. However, since the O vii emission
is from a triplet of lines whose dominant member is unknown, we
allowed the centroid of the line complex to float within a range
of O vii line positions between the forbidden line at 0.561 keV
and the resonance line at 0.574 keV. We found that the sharp rise
seen below 0.5 keV in both the source and background could be
fitted using a low­energy Lorentzian, and we also included delta
functions at 1.78 keV (Si­K) and 2.15 keV (Au­M) to fit the
particle­induced fluorescence seen at these energies. Finally, the
particle­induced continuum background was modeled as a line
with slope and offset, which was not folded through the instru­
mental response.
4. RESULTS
Since we used the standard model of the diffuse X­ray back­
ground, we were not surprised to find a good fit to the data. The
source model had only three significant parameters: the O vii and
O viii line fluxes, and the normalization on the absorbed Galactic
thermal component. The total absorbed flux from this last com­
ponent was only FX (0:4 6 keV ) ¼ 0:066 photons cm #2 s #1
sr #1 , significantly less flux than from either of the two oxygen
lines. Our results for the line emission are shown in Figure 3 and
Table 2. Table 2 also lists O vii and O viii fluxes measured at high
Galactic latitude with ASCA (Gendreau et al. 1995) and XQC
(McCammon et al. 2002). In addition, line fluxes due to helio­
spheric charge exchange (Snowden et al. 2004, hereafter SCK04)
and geocoronal charge exchange ( Wargelin et al. 2004, labeled
``Dark Moon''; see x 5.1) are listed, along with the 2 # upper
limits from the ROSAT observations of MBM 12 (SMV93). The
listed errors are 1 #, except for the Wargelin et al. (2004) data,
where the 90% likelihood range is shown. The SMV93 limits
TABLE 1
Total Counts and Ratios in Selected Bands
ObsID 62850 MBM 12 Obi1 ``Good'' Obi0 ``All'' Obi0
Band Counts Ratio Counts Ratio Counts Ratio Counts Ratio
PHA (2500­3000) ......... 36,197 1.000 21,735 1.000 1407 1.000 75,925 1.000
0.5--2.0 a ......................... 4762 0.132 4015 0.185 314 0.223 46,387 0.611
2.0--7.0 a ......................... 8545 0.236 7492 0.344 513 0.365 63,693 0.839
5.0--10.0 a ....................... 22,742 0.628 14,877 0.684 978 0.695 96,311 1.269
a In units of keV.
3 See http://cxc.harvard.edu /ccw/proceedings/index.html /presentations/
markevitch.
SMITH ET AL.
228 Vol. 623

assume all the emission is from that one line, since ROSAT could
not resolve O vii from O viii. When we allowed the O vii line
to float between the forbidden and resonance line energies, we
found that any value would lead to acceptable fits. The statis­
tics and CCD resolution could not distinguish between either
the low­ or high­energy end, so our results are shown for an as­
sumed ``average'' line position of 0.57 keV. Our best­fit O vii
flux is consistent with the high­latitude Gendreau et al. (1995)
results. However, the O viii emission is significantly larger than
the results of Gendreau et al. (1995) or McCammon et al. (2002).
The strength of the O viii is not strongly dependent on the details
of background subtraction, since it is far from the rapidly rising
background found at low energies on ACIS­S3. We conclude
that this measurement is correct, but contaminated with solar
system emission from charge exchange. Other measurements in
Table 2 show that charge exchange can easily swamp any LB
emission. For example, the Wargelin et al. (2004) ``bright'' re­
sults from the dark Moon (from their Table 7) are significantly
larger than either our results or those of Gendreau et al. (1995),
and the same is true of the SCK04 results.
5. DISCUSSION
Our original goal was to determine the age of the LB by mea­
suring the ratio of the O vii and O viii emission lines. In equi­
librium, the O viii emission would be negligible. A number of
nonequilibrium models, however, predict detectable O viii. Smith
& Cox (2001) considered a range of models involving a series of
supernova explosions. If the LB is still evolving 3--7 Myr after
the last explosion, then their Figure 13b suggests that the O viii
would still be detectable relative to the O vii emission. 4 Alter­
natively, Breitschwerdt & Schmutzler (1994) and Breitschwerdt
(2001) would also imply a detectable amount of O viii emission
(see x 5.3). Solar flares, the variable response of the ACIS­I, and
complicated background of the ACIS­S3 CCD have proved sig­
nificant but not insurmountable obstacles. However, the strong
O viii line is a much larger problem, since it is significantly larger
than has been measured at high latitudes that include both the
LB and Galactic halo emission. As stated above, it seems likely
that our results have been contaminated by charge exchange with
the solar wind.
5.1. Charge Exchange
Charge exchange, as first noted by Cox (1998), is a serious
complication for diffuse soft X­ray background studies. Charge
exchange creates X­rays when electrons jump from neutral ma­
terial (usually hydrogen or helium atoms) to excited levels of
highly ionized atoms. Charge exchange between neutral H and
O +7 , for example, can create O vii emission. The highly charged
ions come from the solar wind and coronal mass ejections (CMEs).
The neutral material can come from either the geocorona or the
interstellar medium as it flows into the heliosphere; see Cravens
(2000) for an overview of this mechanism. Cravens et al. (2001)
show that the so­called long­term enhancements (LTEs) observed
during the ROSAT sky survey, which were apparent brightenings
of the diffuse X­ray sky lasting days to weeks, can be explained
by X­ray emission from charge exchange between the solar wind
and either the geocorona or the interstellar medium, or both.
Wargelin et al. (2004) detected a strongly time­dependent
O vii line along with weaker O viii emission from Chandra ob­
servations of the dark Moon. The source of these X­rays is al­
most certainly charge exchange between solar wind ions and the
geocorona. The time variability of the solar wind makes the time
dependence of the X­ray emission understandable. Detailed spec­
tral models of charge exchange emission have been developed
(e.g., Kharchenko & Dalgarno 2001). Wargelin et al. (2004) used
these calculations to model the O vii and O viii geocoronal emis­
sion as a function of the solar wind oxygen ion flux (which can be
Fig. 3.---Top left: Best­fit spectrum for MBM 12 observation. Bottom left: Best­fit background spectrum from ObsID 62850; note the fluorescence lines at 1.78
and 2.15 keV. Top right: Best­fit spectrum from MBM 12 between 0.4 and 1 keV, showing the strength of the oxygen lines. Bottom right: Same, showing the
background fit.
TABLE 2
Oxygen Line Emission
Ion
This Work
( LU )
ASCA
( LU )
XQC
( LU )
SCK04
( LU )
Dark Moon
( LU )
SMV93
( LU )
O vii .................. 1.79 # 0.55 2.3 # 0.3 4.8 # 0.8 7.39 # 0.79 6.5--13.6 <7.1
O viii.................. 2.34 # 0.36 0.6 # 0.15 1.6 # 0.4 6.54 # 0.34 2.7--6.1 <3.6
4 Note that in Smith & Cox (2001) Figs. 11, 12, 19, and 20 are incorrect; they
should be reduced by a factor of 4# because of an error by the author.
OBSERVATIONS OF MBM 12 AND LOCAL BUBBLE MODELS 229
No. 1, 2005

estimated from ACE measurements). Their model estimates the
surface brightness in each line as
L S ¼
v c n p f O y il # i n n0
4#
5R E
10R E
r min
# # 2
; Ï1÷
where v c is the solar wind velocity, n p is the proton density, f O is
the relative abundance of oxygen, y il is the line yield per charge
exchange, # i is the charge exchange cross section, and n n0 is the
neutral particle density at 10 Earth radii (R E ); r min is the geo­
centric distance to the edge of the magnetosphere or the space­
craft position, whichever is farther. Some of these parameters
are measured by ACE. 5 This model predicts that the O viii /O vii
ratio due solely to charge exchange is 1:36(n O ×8 /n O ×7 ) × 0:14,
where n I is the density of ion I. This result uses values from
Tables 5 and 6 of Wargelin et al. (2004) and including the O vii
K# line in the O viii emissivity, since they would be blended. In
a typical slow solar wind, n O ×8 /n O ×7 # 0:35, implying a ratio of
0.616 (SCK04). Unfortunately, ACE does not yet routinely pro­
vide measurements of n O ×8 /n O ×7 , as it does with n O ×7 /n O ×6 . Dur­
ing Obi1, n O ×7 /n O ×6 was #0.7, approximately midway between
the average value for the slow solar wind (#0.3) and the value of
#1.4 measured during the brightest dark Moon observations.
Going beyond the Earth­Moon system, the time­variable
effects of charge exchange in the broader heliosphere were mea­
sured by SCK04. They used a series of four XMM­Newton obser­
vations of the Hubble Deep Field--North (a high­latitude patch
of sky devoid of bright X­ray sources) to measure O vii and O viii
emission and correlated their results with solar wind observa­
tions from ACE. They found that three out of four observations
(as well as part of the fourth) agreed with the standard diffuse
X­ray background model, and that during these times the pro­
ton flux and n O ×7 /n O ×6 ratios in the solar wind were (1:5 3:2) ;
10 8 cm #2 s #1 and 0.15--0.46, respectively. However, part of
the fourth observation showed substantially brighter O vii and
O viii, as well as sharp increases in both the proton flux (to 8:0 ;
10 8 cm #2 s #1 ) and the n O ×7 /n O ×6 ratio (to 0.99). Their measured
heliospheric surface brightnesses are given in Table 2.
Our viewing geometry for MBM 12 during Obi1 is such that
we would expect an effect due to charge exchange, especially
for O vii. The heliospheric ``downstream'' direction, with respect
to the Sun's motion through the Local Cloud is centered around
ecliptic coordinates 74N5, #6 # ( Lallement & Bertin 1992), which
corresponds to December 5 for the Earth's orbital position.
MBM 12 (47N6, +3 # in ecliptic coordinates) is not far from this
direction, so when Obi1 was done on August 17 it should have
been viewing a sparse column of neutral heliospheric gas---
mostly He, since the H would be largely ionized as it flowed by
the Sun. From Table 9 of Wargelin et al. (2004), we therefore
expect, for a typical slow solar wind, a total ROSAT surface
brightness of #470 counts s #1 sr #1 . The geocoronal component
should be roughly one­fourth of this, as the viewing direction is
(roughly) through the flank of the magnetopause, again from
Table 9 of and Figure 7 of Wargelin et al. (2004). We therefore
expect that the total surface brightnesses will be #5 times the
geocoronal model values in Table 8 of Wargelin et al. (2004), or
1.4 and 0.56 photons cm #2 s #1 sr #1 for O vii and O viii, respec­
tively. Given the uncertainties, and the expectation that charge
exchange emission does not account for the entire soft X­ray
background observed by ROSAT ( but see Lallement 2004), the
O vii prediction matches our result well, but the O viii is still
anomalously high for a typical slow solar wind alone.
However, in addition to the solar wind, it should be noted that
there were a number of halo CME events during 2000 July. 6
These would have the effect of mixing into the Earth's side of the
interplanetary medium some CME plasma that may well have
a substantially higher ionization state. These vary from event to
event, and there are only a few events that are suitable for study
both in the near­Sun environment with SOHO and -->in situ with a
solar wind ion charge state spectrometer.
For example, one such event was observed in 2002 November
and December ( Poletto et al. 2004). The UVCS data measured
remotely by SOHO at 1.7 R # above the solar limb show lines of
Fe xviii, and the SWICS data measured -->in situ by Ulysses also
show high ions of iron ( Fe +16 being prominent). This leads these
authors to conclude that, for this CME event, the ionization state
is characteristic of a ``freezing in'' temperature in excess of
6 MK. At such temperatures, oxygen is fully ionized.
Cravens (2000) argues that solar wind charge exchange with
the neutral interstellar medium ( ISM ) mostly occurs at distances
of a few AU from the Sun because of the depletion of neutral gas
near the Sun. Since 100 km s #1 translates to about 1.7 AU per
month, and CME fronts range from about 20 to over 2000 km s #1
( Webb 2002), we are not concerned with a single event but an
average over the week or two prior to the observation.
We therefore conclude that the large O viii /O vii ratio we
observe could be expected from the charge exchange of CME
plasma with the neutral ISM, but that the details are sufficiently
complex (and the data sufficiently sparse) that prediction of this
ratio is difficult to impossible. We can say, however, that the ob­
served O vii and O viii provides a (weak) upper limit to the emis­
sion from the LB.
5.2. Collisional Models
What do these observations imply if the observed line emis­
sion is due to electron/ion interactions rather than charge ex­
change? In this case it is relatively easy to calculate the expected
contribution from each atomic process. O vii emission from any
of the triplet of lines is created by direct excitation (where we
include cascades from excitation to higher levels), electron re­
combination onto O +7 , or inner­shell ionization of O +5 . O viii
lines come from direct excitation or recombination from bare
oxygen ions; K­shell ionization of O +6 does not have a large
cross section for emitting O viii. If the plasma is at or near ion­
ization equilibrium, then we can ignore blending from satellite
lines or other emission lines in the region of interest, as they
should contribute only a small fraction of the total emission; we
consider the far from equilibrium case in x 5.3. For example, at
CCD resolution the O vii 1s3l ! 1s 2 line will blend with the
O viii Ly# line. However, it contributes at most 6% of the flux of
the O viii Ly# line and so can be ignored to a first approximation.
With these assumptions, the observed line surface brightness
(in LU ) along a particular line of sight can be written as
L S (O vii) ¼ 1
4#
Z dR
n 2
e (r)
1:2 h n ×5 (r)# IS
O vii (T )
× n×6 (r)# DE
O vii (T ) × n×7 (r)# RC
O vii (T ) i ; Ï2÷
L S (O viii) ¼
1
4#
Z dR
n 2
e (r)
1:2 h n×7 (r)# DE
O viii (T )
× n×8 (r)# RC
O viii (T ) i ; Ï3÷
5 Data available at http://www.srl.caltech.edu/ACE.
6 See catalog at ftp:// lasco6.nascom.nasa.gov/pub/ lasco/status / LASCO _
CME _ List _ 2000.
SMITH ET AL.
230 Vol. 623

where L S is in LU, n e (r) is the electron density at distance r,
and n×n (r) is the ion density of O ×n at r. We also assume that
the hydrogen and helium are fully ionized, so n e # 1:2n H ;
# IS
I (T ), # DE
I (T ), # RC
I (T ) are the rate coefficients (in cm 3 s #1 )
for creating an emission line via inner shell, direct excitation,
or recombination for ion I, and are plotted in Figure 4.
Since four different ions may be involved in creating O vii and
O viii emission lines, we cannot uniquely solve for the density,
temperature, and ionization state of the system. However, we can
easily show that these line fluxes are not commensurate with an
equilibrium plasma. The O viii flux provides the strongest limit,
as this line is not expected in an equilibrium model for the LB.
Figure 4 shows that the direct excitation rate for this line rises
rapidly with temperature, so a higher temperature will produce
more O viii. Based on the Kuntz & Snowden (2000) survey of LB
results (see their Table 1), we assume that the allowed electron
temperature in the LB is <2 ; 10 6 K; in equilibrium at this tem­
perature, n O ×6 /n O ×7 > 1. Figure 4 clearly shows that the direct
excitation rate for O vii is more than triple the rate for O viii ex­
citation, which would predict O vii /O viii > 3, in strong contrast
to our result. Therefore, either the O viii line emission is not com­
ing from the LB, or the plasma is recombining, with significantly
more O +7 or O +8 ions than O +6 ions. We consider some aspects
of this possibility below.
5.3. Recombining Plasma Models
Breitschwerdt & Schmutzler (1994) proposed a radical
new model for the LB suggesting that it was not in fact ``hot,''
but rather the X­ray emission was created by recombination
of highly ionized atoms with cool (10 4 to 3 ; 10 5 K) electrons.
Their model assumes that the LB was a preexisting cavity, prob­
ably formed by earlier supernova explosions and possibly as­
sociated with the Sco­Cen superbubble. Then, inside this cavity a
few million years ago, a massive star inside a dense cloud ex­
ploded. After the supernova shock­heated and ionized the entire
cloud, the resulting hot ions expanded into the low­density cav­
ity and cooled adiabatically. This model was partially inspired by
the relatively high average density of 0.024 cm #3 measured from
a nearby (130 pc distant) pulsar's dispersion measure (Reynolds
1990). If this density is representative of the LB as a whole,
then an #10 6 K plasma would create far more soft X­rays than
are observed. Breitschwerdt & Schmutzler (1994) showed that
a low­temperature, high­ionization state plasma with this den­
sity could generate the observed 1
4
keV X­ray emission. Later pa­
pers ( Breitschwerdt et al. 1996; Breitschwerdt 2001) have de­
scribed the model in more detail and made predictions of the
flux in various bandpasses. This model also naturally explained
the existence of the ``Local Fluff,'' neutral H i clouds (e.g.,
Bochkarev 1987; Frisch 1995), which are hard to understand in
the context of hot bubble models.
Although more complicated than the equilibrium isothermal
LB model, a strength of this model is that it does not require
a particular fine­tuning or a complicated mix of hydrodynamic
and plasma models. As noted by Breitschwerdt (2001), only some
basic input parameters---cavity radius R, electron density n e ,
electron temperature T e , and evolution time t ---are needed, and a
fairly wide range of these values matched the available data at the
time.
However, FUSE observations (Shelton 2003) put a surpris­
ingly low 2 # upper limit on the O vi kk1032, 1038 emission
from the LB of 500 and 530 LU, respectively, and 800 LU for
both lines combined. Shelton (2003) noted that this result alone
strongly constrains any recombining plasma model and, when
combined with other discrepancies in the C iii emission and the
N(O vi) column density, ``practically eliminate[s] this class of
models.'' However, this argument rests on discrepancies in two
unrelated ions, O vi and C iii, leaving open the possibility that
these issues could be finessed with the correctly tuned input pa­
rameters. Recombining plasma models are still being invoked
in recent papers (e.g., Freyberg et al. 2004; Hurwitz et al. 2005).
The O vi upper limits are quite stringent, however, and combin­
ing them with our measurements allows a rigorous test of the
model using all relevant oxygen ions as they recombine from
fully stripped O +8 to O +5 and beyond. Any acceptable recom­
bining plasma model would have to match the upper limits found
for each ion's emission lines, while also emitting at least some
of the observed 1
4 keV band X­rays.
Interestingly, in the case of a purely recombining plasma, the
ionization state of the plasma at any point in its evolution can be
calculated analytically. If ionization can be ignored, then the
population of any ion state can be written as
dp i
dt ¼ #R i n e p i
× R i×1 n e p i×1 ; Ï4÷
where p i is the population of the ion with i electrons and R i is the
recombination rate (in cm 3 s #1 ) from the ith ion state to the i × 1
state. In the case of the fully stripped ion ( p 0 ), the second term
in equation (4) is zero, as there is no higher ion. Equation (4)
then leads to a linked series of first­order differential equations
that can be easily solved. Assuming that initially the population
is fully stripped [ p 0 (0) # 1], the population of any ion can be
written as
p i
¼ X i#1
j¼0
R i#1 f i#1
j
R i # R j
exp (#R j n e t) # exp (#R i n e t)
# #
( ) ; Ï5÷
where f i
j equals 1 when i ¼ j ¼ 0 and is otherwise defined by
the recursion relation:
f i
j ¼
R i#1 f i#1
j
R i # R j
if j < i;
# P i#1
j¼0
R i#1 f i#1
j
R i # R j
if j ¼ i:
8 > > > < > > > :
Ï6÷
Fig. 4.---Shows # curves as a function of electron temperature, using APEC
and data from ATOMDB, ver. 1.3.1. The total emission from each process is
n e n I #.
OBSERVATIONS OF MBM 12 AND LOCAL BUBBLE MODELS 231
No. 1, 2005

To calculate R i we used radiative recombination rates from
Verner & Ferland (1996) and dielectronic rates from Romanik
(1988, 1996). Combining this with Figure 4 we can calculate the
predicted O viii Ly# and O vii triplet flux due to recombination
for any temperature and density. We calculated the O vi flux from
electron collisions (recombination from O +6 was negligible) us­
ing the collision strengths from Griffin et al. (2000).
However, the upper limits on the O viii, O vii, and O vi emis­
sion and the N(O vi) column density were not sufficient, as a suf­
ficiently low density would always be allowed (see Fig. 5, bottom
right). The recombining ion model, however, was initially de­
veloped with a relatively high density in mind and tested against
the observed 1
4 keV band emission. We have developed an ap­
proximate but robust model of the X­ray spectrum from a re­
combining plasma to compare to the ROSAT R12­band emission.
This model considered only emission from dielectronic and ra­
diative recombination, omitting direct excitation. In a plasma
with low electron temperature, the radiative recombination con­
tinuum appears as a nearly linelike (width #kT ) feature at the
binding energy of the recombined level. We therefore assumed
that each radiative recombination led to a photon at the binding
energy of the ion; this ignores cascades, which would modify
somewhat the exact distribution of the photons but is reasonably
accurate for our purposes. The dielectronic satellite lines were
calculated using data from Romanik (1988, 1996). We then
folded the resulting spectrum between 50 and 1000 eV through
the ROSAT R12­band response to calculate the expected emis­
sion in this band.
For the O vii and O viii limits we used the 2 # upper limits from
the MBM 12 observation, despite our strong suspicion ( based on
the ACE data and the odd ratio of the two lines) that these are
already contaminated by charge exchange. In addition, in the case
of O viii our 2 # upper limit is larger than the value observed by
XQC ( McCammon et al. 2002) over a 1 sr field of view, which
necessarily includes both LB and more distant emission. There­
fore, we feel that this is a very conservative upper limit to the
local contribution from O vii and O viii.
Defining an ``upper limit'' to the column density of O vi in the
LB is difficult, and so we attempted to determine a reasonable
value by indirect methods. Shelton & Cox (1994) found through
a statistical analysis that the most likely value for N(O vi) in the
LB was 1:6 ; 10 13 cm #2 . The Copernicus data this was based on
(Jenkins 1978) also shows that no O or B star within 250 pc has a
column density greater than this value. In addition, FUSE results
toward 25 nearby white dwarfs (Oegerle et al. 2005) show at
most 1:7 ; 10 13 cm #2 , with an average value of #7 ; 10 12 cm #2 .
We therefore chose N(O vi) < 1:7 ; 10 13 cm #2 as our upper
limit to the column density of O vi through the LB.
Fig. 5.---Top left: O viii Ly#, O vii triplet, and O vi doublet emission and observed upper limits for two different models as a function of fluence. Thick curves
show a model with R ¼ 114 pc, n e ¼ 0:024 cm #3 , and T e ¼ 3 ; 10 5 K; thin lines are for a model with R ¼ 114 pc, n e ¼ 0:01 cm #3 , and T e ¼ 2 ; 10 4 K. Top right:
O vi column density in the LB for the same models, along with the observed upper limit from the LB. Bottom left: ROSAT R12 emission from both models, with
a line showing our estimated lower limit. Bottom right: Maximum allowed electron density for a 114 pc bubble in either model, based only on the O viii, O vii, and
O vi limits. The effects of each oxygen ion's limit are visible. Except at large fluences (when the R12 emission is negligible), the maximum allowed density is less
than the assumed model density, showing why these models are excluded.
SMITH ET AL.
232 Vol. 623

We wanted to choose a very conservative value for the min­
imum R12­band emission from the LB required by the model.
Breitschwerdt & Schmutzler (1994) assumed that some of the
observed soft X­ray emission came from distant superbubbles,
since ROSAT had recently seen clouds in absorption. Nonethe­
less, they also required some local emission, although the re­
quired amount was somewhat uncertain. Since then, Kuntz et al.
(1997) studied the foreground X­ray emission by analyzing the
shadows seen by ROSAT toward nearby molecular clouds. They
measured the foreground R12 emission seen toward nine clouds,
finding a range of 3790--6190 counts s #1 sr #1 . This was followed
by the exhaustive study of X­ray shadows by Snowden et al.
(2000), who examined 378 clouds seen with ROSAT. Their re­
sults showed substantial variation from position to position, with
a mean and standard deviation of 5810 # 1960 counts s #1 sr #1 .
Based on these results, we decided to require at least 910 counts
s #1 sr #1 in the R12 band, far below the value found by Kuntz et al.
(1997) and 2.5 # below the Snowden et al. (2000) mean value.
In Figure 5 we show the predicted O viii, O vii, and O vi
emission, the O vi column density, the R12 emission for two
recombining plasma models, and the maximum allowed elec­
tron density based solely on the oxygen ion upper limits for both
models. The oxygen, neon, and argon abundances were taken
from Asplund et al. (2004), with all other abundances taken from
Anders & Grevesse (1989). The plots are shown as functions
of the fluence (# n e t), a convenient variable as all the rates are
proportional to the electron density. The thick curves show the
results for the model described in Breitschwerdt (2001), with a
cavity radius of 114 pc, an electron density of 0.024 cm #3 , and
an electron temperature of 3 ; 10 5 K. Note how the R12 emis­
sion matches the observations over a wide range offluences, sug­
gesting that our 1
4 keV emission model agrees with that used in
Breitschwerdt (2001). Whenever the R12 emission reaches and
exceeds the lower limit in this model, however, the O vi emission
exceeds its upper limit. The thin curves show a second model,
with the same radius, an electron density of 0.01 cm #3 , and an
electron temperature of 2 ; 10 4 K. In this case there is a small
region with fluence #10 10 cm #3 s, where the R12 and O vi
emission pass their lower and upper limits, respectively. How­
ever, in this range the predicted O vii and O viii substantially
exceed the upper limits toward MBM 12. At later fluences the
predicted O vii and O viii emission drops, but the N(O vi) column
density exceeds the observed value. In sum, there is no point that
simultaneously meets all the upper and lower limits.
The R12 emission calculation is necessary, since it pro­
vides the only lower limit, but it is also quite complex to calcu­
late, and it is useful to consider what can be derived purely from
the simpler oxygen ions. Figure 5 (bottom right) plots the limit
on the electron density available purely from the oxygen emis­
sion and absorption upper limits. For example, the O viii emis­
sion is given by
O viii ¼ n e n ×8
R
4#
# RC
O viii (T ); Ï7÷
which can easily be converted into an upper limit on the elec­
tron density, since the remaining terms are determined by re­
combining ion model. As can be seen in the model, the limits
from the oxygen ions alone require the electron density to be
below the assumed density until very late times or fluences.
We have examined a wide range of input parameters for the
recombining plasma model. We considered electron densities up
to 0.03 cm #3 , electron temperatures in the range 10 4 to 5 ; 10 5 K,
and all cavity radii between 30 and 300 pc, and followed each set
of values for 10 7 yr, the maximum plausible lifetime of the LB.
Although the results vary substantially with the input parame­
ters, as shown in Figure 5, none of our input parameters lead to
predicted values that can simultaneously match these observa­
tional limits. In essence, the R12 lower limit requires a minimum
density of highly ionized ions, while the tight observational
limits on O vi, O vii, and O viii simultaneously limit the density of
these ions. We can with confidence dispose of at least the static
recombining plasma model.
6. CONCLUSIONS
We have observed the nearby molecular cloud MBM 12 with
Chandra and measured the foreground O vii and O viii surface
brightness toward the cloud, which should absorb any distant
emission. Our observed values are higher than expected, and it
appears likely, based on the ACE data, that charge exchange in
the heliosphere or geocorona has contaminated our results. We
measure a unexpectedly low ratio of O vii /O viii ¼ 0:76 # 0:26,
which is hard to understand in the context of other measurements
or most LBmodels. In addition, circumstantial evidence also sug­
gests that our results could be contaminated by charge exchange.
Despite this limitation, we are able to combine our upper lim­
its with results from FUSE on O vi emission and absorption and
the 1
4 keV ROSATR12 emission to reject the constant temperature
recombining plasma model originally proposed by Breitschwerdt
& Schmutzler (1994) over a wide range of input parameters.
Since the sole moderate­resolution spectrum of the LB taken by
DXS showed that the constant temperature equilibrium model
could also be rejected (Sanders et al. 2001), this result implies
that more complex models of the soft X­ray emission from the
LB are required.
A number of more complex models have been proposed,
however (e.g., Smith & Cox 2001; de Avillez & Breitschwerdt
2003), and selecting among the various possibilities will require
high­resolution observations. If, for example, we could resolve
the O vii triplet, then we would be able to determine whether the
plasma is ionizing, recombining, or both in different regions.
Although Astro­E2 will have the necessary resolution, its effec­
tive area--solid angle product is only #330 cm 2 arcmin 2 , com­
pared to #20,000 cm 2 arcmin 2 for ACIS­S3. Constellation­X,
however, will be a powerful instrument for understanding the
recent history of our local environment.
We are grateful to Don Cox, Mike Juda, John Raymond,
Robin Shelton, Steve Snowden, and Shanil Virani for helpful
discussions. This work was supported by the Chandra X­Ray
Science Center ( NASA contract NAS8­39073) and NASA
Chandra observation grant GO0­1097X.
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