Chandra observation of the TeV source HESS J1834-087

0.77 MB·English

by  Zdenka Misanovic

PreprinttypesetusingLATEXstyleemulateapjv.2/16/10
CHANDRA OBSERVATION OF THE TEV SOURCE HESS J1834–087
Zdenka Misanovic1, Oleg Kargaltsev2 and George G. Pavlov3,4
1SchoolofPhysics,MonashUniversity,Melbourne,3800VIC,Australia([email protected])
2Dept. ofAstronomy,UniversityofFlorida,Gainesville,FL32611-2055 ([email protected]fl.edu)
3Dept.ofAstronomyandAstrophysics,ThePennsylvaniaStateUniversity,
525DaveyLab.,UniversityPark,PA16802([email protected])
4St. PetersburgStatePolytechnical University,Polytechnicheskaya ul. 29,St.Petersburg195251, Russia
1 ABSTRACT
1 Chandra ACIS observedthe field of the extended TeV source HESS J1834–087for 47 ks. A previous XMM-Newton
0
EPICobservationofthesamefieldrevealedapoint-likesource(XMMUJ183435.3–084443)andanoffsetregionoffaint
2
extended emission. In the low-resolution, binned EPIC images the two appear to be connected. However, the high-

n resolution Chandra ACIS images do not support the alleged connection. Instead, in these images XMMU J183435.3–
a 084443isresolvedintoapointsource,CXOUJ183434.9–084443(L≃2.5×1033ergss−1,foradistanceof4kpc;photon
J index Γ ≃ 1.1), and a compact (. 20′′) nebula with an isotropic morphology and a softer spectrum (L ≃ 4.2×1033

7 ergs s−1, Γ ≃ 2.7). The nature of the nebula is uncertain. We discuss a dust scattering halo and a pulsar-wind
nebula as possible interpretations. Based on our analysis of the X-ray data, we re-evaluate the previously suggested

] interpretations of HESS J1834–087 and discuss a possible connection to the Fermi LAT source 1FGL J1834.3–0842c.
E
We also obtained an upper limit of 3× 10−14 ergs cm−2 s−1 on the unabsorbed flux of the SGR J1833–0832 (in
H
quiescence), which happened to be in the ACIS field of view.
h. Subjectheadings: gammarays: individual(HESSJ1834–087,1FGLJ1834.3–0842c)—stars: neutron—
p supernova remnants: individual (W41=G23.3–0.3)—X-rays: individual (CXOU
- J183434.9–084443,XMMU J183435.3–084443)
o
r
t
s 1. INTRODUCTION by energetic electrons injected/accelerated in the vicin-
[a Surveys of the Milky Way plane with modern Imag- ityofthepulsar(Lorentzfactor∼108),theVHEPWNe
ingAtmosphericCherenkovTelescopes(IACTs)havere- are powered by cooled, less-energetic electrons (Lorentz

1 vealed more than 70 sources of very high energy (VHE) factor ∼107), which have propagatedfurther awayfrom
v γ-rays1. While X-ray binaries (XRBs), young stel- thepulsar(e.g.,de Jager & Djannati-Ata¨ı2008). TheIC
2 lar clusters and background AGNs have all been pro- cooling time of these γ-ray-producingelectrons is longer
4 posed as possible sources of emission in the TeV en- than that of the electrons that produce X-rays, and
3 ergy range, the majority of the classified Galactic γ- hence, they trace the cumulative history of the PWN,
1 ray sources are thought to be associated with pul- similar to radio emission (the relic PWN model). If the
.
1 sar wind nebulae (PWNe) and supernova remnants pulsar contains a hadronic component, the VHE emis-
0 (SNRs)(Aharonian et al.2005,2006;Gallant et al.2008; sioncouldalsobe producedviaa hadronicprocesswhen
1 Kargaltsev & Pavlov 2010). However, about 30% of the relativistic PW hadrons encounter dense material such
1 cataloged VHE sources lack reliable classifications. Fur- as the host SNR shell. An alternative possibility is that
v: thermore, for many classified sources there remains an the SNR shock accelerates protons, which then interact
i uncertainty about the actual mechanism responsible for withthedenseISM(e.g.,amolecularcloud)andproduce
X theproductionoftheVHEγ-rays. Therefore,multiwave- VHE emission.
r length studies and modeling of the VHE source popula- The extended source HESS J1834–087 (18.7± 2.0×
a tion are of a great interest at the present time. 10−12 photons cm−2 s−1 at E >200 GeV, photon index
The two main mechanisms for the production of high- Γ = 2.5±0.2; Aharonian et al. 2005, 2006) is spatially
energy γ-rays are the inverse Compton (IC) scattering coincidentwiththe centerofthe radioshellSNRG23.3–
of low-energy photons, and the hadronic collisions, in 0.3 (W41; Kassim 1992), making the remnant a promis-
whichtheenergeticprotonscollidingwithambientparti- ingcandidateforgeneratingtheobservedγ-rayemission.
clesproduceneutralpions,whichdecayintohigh-energy However,thesmallersizeofthe HESSsource2 (diameter
photons (π0 → 2γ; see, e.g., Hinton & Hofmann 2009). ∼ 10′ compared to the SNR diameter of ∼ 30′) and the
Rotation-powered pulsars emit highly energetic pulsar projectedlocationwellwithintheSNRshellsuggestthat
wind (PW), which is responsible for the synchrotronra- the γ-ray emission does not come from the entire SNR
diation observed from radio up to MeV energies (see shell as it has been seen in other SNRs firmly associated
Kargaltsev & Pavlov 2008, 2010, for recent reviews). It withHESSsources(e.g.,Acero et al.2009). Albert et al.
is believed that the γ-ray PWN can be produced when (2006) detected the TeV source with the Major Atmo-
the ambient microwave, optical or infrared photons are spheric Gamma Imaging Cerenkov (MAGIC) telescope,
upscattered by relativistic PW electrons up to the GeV confirmeditsextension,andmeasuredthefluxconsistent
and TeV energies. While the X-ray PWNe are powered with the HESS measurement.
The region was also observed in radio with the Very
1 Seehttp://tevcat.uchicago.edu/
2 SeeTable3andSection5.1inAharonianetal.(2006).
2 Misanovic, Kargaltsev & Pavlov
Large Array (VLA), as part of the Galactic Plane sur- 2.1. Spatial analysis
vey in continuum and in CO molecular line emission, 2.1.1. Previous multiwavelength observations
and also in X-rays with XMM-Newton. Based on these
Figure1showstheVLA20-cmimageofW41(Whiteet
observations,Tian et al. (2007) estimated the remnant’s
age (104 −105 yrs) and the distance (4 kpc), and sug- al. 2005), revealing fine details of the highly structured
partial shell with a diameter of ∼30′, and also some ex-
gested that the observed γ-ray emission is produced by
tendedradioemissionnearthecenter. However,nopoint
a hadronic process when the particles accelerated by
radio sources are visible, and no detection of radio pul-
the SNR shock interact with a giant molecular cloud
sations from within the SNR shell has been reported so
(GMC) projected near the center of W41 (see also
far. The extended VHE source HESS J1834–087 is lo-
Yamazaki et al.2006). Anotherexplanationwasputfor-
cated near the SNR center. With an angular extent of
ward by Bartko & Bednarek (2008), who proposed the
′
′ 5.4 in radius (see Table 3 in Aharonian et al. 2006) the
nearby pulsar J1833–0827 (=B1830–08; 24 away from
TeV source appears to be significantly smaller than the
the center of the HESS 1834–078) as the source of rel-
SNR shell (Figure 1, left).
ativistic electrons. These electrons were ejected at the
The middle panel of Figure 1 shows the innermost
pulsar’s birth place, presumably near the SNR center,
region of the extended TeV source as it is seen in a
i.e.,withintheextentoftheHESSsource(arelicPWN).
17 ks exposure by XMM-Newton EPIC (for the anal-
Gaensler & Johnston (1995) argued that the association
ysis of XMM-Newton data see Tian et al. 2007, M+09).
of the PSR J1833–0827 and W41 is plausible, given the
M+09 detected a total of 16 point-like sources in the
pulsar’scharacteristicageof148kyrandthecorrespond-
ing projectedvelocity of ∼250km s−1, requiredto move FoV of the MOS detectors, with one of the bright-
est sources well within the extent of the TeV emis-
the pulsar from its birthplace to the current location.
sion. Furthermore, in the EPIC images this bright
However,ina recentanalysisofthe XMM-Newton data,
point-likesource(XMMU J183435.3–084443)appearsto
Schmitt et al. (2010, to be submitted to ApJ) rules out
be accompanied by extended X-ray emission, also lo-
theassociationofB1830–08andW41. Schmittetal. de-
cated within the extended TeV source. The spectral
tectedapossiblePWNsurroundingthepulsar,butsince
analysis of the XMM-Newton data indicates that both
the morphology of the nebular emission does not resem-
the point-like source and diffuse emission (designated
bleabow-shock,theyconcludedthatthepulsardoesnot
XMMU J183435.3–084443 and G23.234–0.317, respec-
movefastenoughto travelfromthe center ofW41 to its
′ tively; M+09) are highly absorbed and most probably
current position, approximately 26 away.
non-thermal, although the spectral parameters for the
Subsequent analysis of the XMM-Newton observa-
diffuseemissionarenotwellconstrainedduetothelarge
tions (Mukherjee et al. 2009, hereafter M+09) revealed
background contribution. GLIMPSE3 8µm image (Fig-
a bright point-like source XMMU J183435.3–084443 at
ure 1, bottom) shows no correlation between the diffuse
the center of W41 apparently connected to a region of
X-ray and IR emission.
faint extended emission also located within the extent
of the HESS source. M+09 have argued that the new 2.1.2. Images
pulsar/PWN candidate is the most likely counterpartof
OurChandraobservationresolvedXMMU J183435.3–
HESS J1834–087.
084443 into a point source, which we designate
To discriminate between these alternative scenarios,
CXOU 183434.9–084443 (centered at R.A.=18h 34m
we have observed the region of the TeV source with the 34s.944, Decl.=–08o 44′ 43′.′09), and a compact diffuse
Chandra ACIS detector. In Section 2 we present the emission extending up to ≃ 20′′ from the point source
multiwavelengthdataonHESSJ1834–087,which,inad-
(Figure2,top),withmoreorlessisotropicsurfacebright-
dition to the Chandra observation,also include archived
ness distribution, and an average surface brightness of
radio, infrared (IR) and γ-ray data. We further discuss 0.8 cts arcsec−2 in an annulus with radii 1′.′5 and 10′′
ourresultsinSection3andsummarizetheminSection4.
(S/N≈14). To separate the X-ray diffuse emission from
2. OBSERVATIONSANDRESULTS the PSF wings of the pulsar candidate image, we have
simulated a PSF, using MARX4 and assuming the mea-
We observed the region around HESS J1834–087 on
sured spectral properties of CXOU J183434.9–084443.
2009 June 7 with Chandra ACIS (ObsID 10126) for
The comparisonof the data and point source simulation
47.5 ks (46.5 useful scientific exposure) in timed expo-
(Figure 3) suggests a good agreement between the data
sure (TE) mode using very faint (VF) telemetry format.
and simulation in a small aperture (approximately up
The point source was imaged on the S3 chip approxi- ′′
′′ to 1 radius) while the extended emission is easily seen
mately 15 away from the nominal aimpoint, placing a
at larger radii. The 2MASS image of the same region
large part of the extended emission detected by XMM-
(Figure2,bottom)showsnonear-infrared(NIR)sources
Newton on the same chip. In addition, S1, S2, S4, I1 ′′
within 6 from CXOU J183434.9–084443position.
and I2 chips were also activated during the observation.
WesearchedforotherpointsourcesintheACISimage
We used CIAO 4.2 and CALDB 4.2.2 for our analysis.
in particular within the extent of the HESS source. We
Wehaveappliedadditionalparticlebackgroundcleaning
for the VF mode, using the pulse heights in the outer
3 The Galactic Legacy Infrared Mid-Plane Survey; see
16 pixels of the 5 × 5 event island to help distinguish
http://www.astro.wisc.edu/sirtf
between “good” X-ray events and “bad” events that are 4 For the MARX simulation we used the same off-axis and roll
most likely associated with cosmic rays. We also made angles,andthedetectorparametersasintherealobservation,but
extensiveuseofthearchivalmultiwavelengthdataofthis increased the exposure time by a factor of 100 (rescaling the re-
sultingimage)toreducethestatisticalerrors. Wefoundthatusing
regionobtainedwithXMM-Newton,VLAandSpitzerto theditherblurparameterof0′.′35givesthebestmatchbetweenthe
compare them with the Chandra images. dataandsimulation.
Chandra observation of J1834-087 3
found a faint (∼30 cts in 0.5–8 keV) point-like source
N at R.A.=18h 34m 42s.6, Decl.=–08o 45′ 01′.′95, approx-
′
imately 2 east of the pulsar candidate (Figure 4, top).
The source coincides with the region of the enhanced
HESS J1834-087 E
“tail” emission detected by M+09, although no point
sourcewasdetectedatthatpositionintheXMM-Newton
data. This could be due to the lower angular resolu-
tion and higher EPIC background or due to the tran-
sient nature of the source. The Chandra spectrum of
W41 this source is soft, with an absorption column almost
an order of magnitude lower than that of the candidate
20 cm
pulsar/PWN. Furthermore, the X-ray source coincides
with a 2MASS source, suggesting that it is a foreground
star. Most of the point-like X-ray sources detected in
the XMM-Newton data (Table 1 of M+09) are either
outside the FoV of the activated ACIS chips, far off-
axis,ornearthe chipedges,butwediddetectthe bright
sources11and15onI2andI3chips,respectively(follow-
ing the source designation by M+09). The X-ray spec-
MOS1+MOS2
tra of these sources are relatively soft, with absorption
2-10 keV columnsN =(1−9)×1021 cm−2,suggestingthatthey
H
are relatively nearby. Source 15 coincides with a faint
2MASSsource,alreadylistedinTable1ofM+09. Using
our more accurate Chandra position, we also find two
bright2MASSsourceswithintheChandraerrorcircleof
the source 11. We also found that the position of the
recently detected SGR J1833–0832(Gelbord et al. 2010;
G¨og˘u¨¸s et al.2010)iswithintheACISFoV,neartheedge
ofthe S1chip(Figure 4,bottom),butwedidnotdetect
2’
any source at this position. After correcting for the vi-
gnetting,andassumingthespectralparametersreported
byGelbord & Vetere(2010),we obtainedanupper limit
of3×10−14 ergscm−2 s−1 ontheunabsorbedfluxofthe
SGR in quiesence.
While looking for high-energy sources within the
ACIS FoV, we found Fermi LAT source 1FGL J1834.3–
GLIMPSE
0842c (0.11 Crab in 1–100 GeV, photon index Γ =
2.24±0.04; Abdo et al. 2010) located ∼5′ northwest of
8 micron
CXOU 183434.9–084443. The 95% error ellipse of the
1FGL J1834.3–0842c position also includes the north-
western part of the SNR shell (Figure 4, bottom).
The only X-ray source detected within the error ellipse
of the Fermi LAT source is a relatively faint source,
CXOU 183430.3–084142.8, with a bright optical coun-
terpart, located at the edge of the GeV emission. This
source is apparently a foreground star, most likely not
associated with the GeV emission.
Figure4(top)showsthelarge-scaleACISimage,which
covers the region of the extended HESS source. The
heavilysmoothedimageshowssomefaintextendedemis-
sion,whichmustbeapartofthelarge-scaleemissionseen
in the XMM-Newton images. Indeed, a closer compari-
son of the two images (see Figure 1) shows that the ex-
Fig.1.— Top: VLA 20-cm imageof the SNR G23.3–0.3 (W41) tendedemissionvisibleintheACISimagecoincideswith
showing its highly structured partial shell with a diameter of
∼30′ and some enhanced emission near the center (White et al. the brightest regions seen in the XMM-Newton data.
2005). The circle (radius 5.′4) shows the position and extent of However,the originofthis emissionremains elusive. We
HESSJ1834–087asmeasuredbyAharonianetal.(2006). Middle: seenoclearconnectionbetweenCXOU183434.9–084443
Broad-band(0.2–10keV)XMM-Newton17ksimageoftheinterior andtheextendedemission3′−4′ east-southeastofit. Its
of W41 combining MOS1 and MOS2 data (see also M+09). The
image with the pixel size of 8′′ is smoothed with a Gaussian of associationwith the SNR is also questionable, as it does
FWHM 40′′. The cross marks the position of CXOU 183434.9– notseemtocoincidespatiallywiththeradioshellorwith
084443 (see text). Higher-resolution images of the vicinity of the diffuse radio emission in the SNR interior (Figure 4,
CXOU183434.9–084443areshowninFigure2. Thetail-likeemis-
middle and bottom). To investigate the nature of this
sion detected in the XMM-Newton observation is marked in all
panels. Bottom: GLIMPSE (8 micron) image of the same region extended emission, its spectrum must be measured in a
oftheskyasinthemiddlepanel.
4 Misanovic, Kargaltsev & Pavlov
long observation. parameter when fitting the spectra simultaneously, with
XMMU J183435.3–084443 has also been detected by other parameters untied).
the Swift (Landi et al. 2006). M+09 included the three After obtaining the best-fit value for N , we fixed
H
archived Swift observations in their analysis, but could this parameter before fitting the ACIS spectra of the
not make any firm conclusions about intrinsic flux vari- point source and the nebula individually and calculat-
abilityofXMMUJ183435.3–084443(thecountratemea- ing the errors for other model parameters and for the
suredin the third observationshoweda 3σ increase over fluxes (Table 1). The unabsorbed luminosities are cal-
the first two, or 2σ increase over the mean rate). To culated for the 0.5–8keV band assuming a distance of
search for variability in the ACIS observations, we pro- 4 kpc (Tian et al. 2007). We note that the large ab-
duced light-curveswith a bin size rangingfrom ∼102 to sorption measured from the X-ray spectra is more than
103 seconds, but could not see any significantflux varia- an order of magnitude higher than the estimated N
HI
tions on the time-scales relevant for the ACIS data (i.e., value for the Galaxy in that direction (∼2×1022 cm−2;
considering the length and time resolution of the ACIS Dickey & Lockman 1990), implying a large amount of
observation). OurChandradataallowedustosearchfor molecular material in the line of sight (see Tian et al.
pulsations in the 6–100 s period range, but we did not 2007, and references therein).
find any periodic signal.
2.2. Spectral analysis
To minimize the contribution from the nebula, we ex-
tracted the X-ray spectrum of CXOU J183434.9–084443
′′
fromacircularregionwitharadiusof1 centeredonthe
point source (see Figure 3), while the background pho-
tons were extracted from an annulus with the inner and
′′ ′′
outer radii of 1.5 and 5.5, respectively. A total of 220
′′
counts were detected in the r = 1 aperture in the 0.5–
8 keV band. The comparison with the simulated point
sourcesuggeststhatapproximately30ofthesecountsare
from the nebula, while the contribution from the back-
ground is negligible (1–2 counts), implying an aperture-
corrected5 count rate of 4.80±0.04 cts ks−1. For spec-
tral fitting these counts were grouped to a minimum of
15 counts per bin.
To measure the nebular spectrum, we extracted a to-
′′ ′′
talof208photonsfromthe1.5–10 annuluscenteredon
CXOU J183434.9–084443, while the background contri-
′′
bution was measured from an annulus with 10 < r <
20′′. We estimate that there are ∼50 counts from the
true backgroundand ∼10 counts from the PSF wings in
the same aperture, implying a count rate of 3.2±1.1 cts
ks−1 for the background- and PSF-subtracted emission.
The nebula spectrum wasalso groupedto a minimum of
15 counts per bin.
The small number of counts and the lack of soft emis-
sion (the spectra are virtually cut off below 3 keV; see
Figure 5) did not allow us to constrain well the spec-
tralmodelparameters. Wefitthespectrawithabsorbed
power-law (PL) models and show the results in Table 1
and Figures 5 and 6. We first fit each spectrum inde-
pendently and noticed that the absorption columns for
the point sourceand for the diffuse emissionregionwere
similar (2.8+2.8 and 2.8+2.9 ×1023 cm−2, respectively),
−1.3 −1.5
suggesting a common distance. To determine the value
of the absorption column more accurately, we then per-
formed a simultaneous fit with the MOS1 and MOS2
spectra6 of XMMU J183435.3–084443, assuming again
a common absorption column density (i.e., linking this
5 From the radial profile of the simulated PSF and the data
(Figure 3) we estimated the point source contribution of 85% in
thisaperture.
6 The MOS spectra were only used to measure the absorption
column, because they contain not only photons from the pulsar
candidate but also a large fraction of the compact nebula that
couldnotberesolved.
Chandra observation of J1834-087 5
TABLE 1
Absorbed PLmodel fitsto the spectra of CXOUJ183434.9–084443andsurroundingextended emission.
NH Γ PLNorm. χ2ν/d.o.f AbsorbedFlux Luminosity
1023 cm−2 10−4cm−2s−1keV−1 10−13ergscm−2s−1 1033ergss−1
Pointsource 2.72(frozen) 1.09+−00..4402 0.95+−00..9427 0.81/13 2.56+−00..2222 2.3+−00..64
Nebula 2.72(frozen) 2.73+−00..5543 6.8+−93..19 0.64/13 0.98+−00..1132 4.1+−41..89
Note. — Theabsorptioncolumndensitywasfrozenatitsbest-fitvaluedeterminedinthejointfitwithXMM-Newtondata(MOS1andMOS2
spectraof the point-like sourceXMMU J183435.3–084443). The observed flux and unabsorbedluminosity, LX = 4πd2FXunabs, are calculatedfor
the3–8keVand0.5–8keVenergyband,respectively,forthedistanced=4kpc. Thelisteduncertaintiesareata1σconfidencelevel. Theobserved
fluxandunabsorbedluminosityofCXOUJ183434.9–084443areaperture-corrected.
30’’
Fig.3.— Theradialprofiles(in3–8keV)ofthesimulationand
datashowingthecontributionoftheextendedemissioninthevicin-
ityof the point source. The dashed linewith errorbarsshows the
PSFprofileofthepointsourcesimulatedusingMARX(seetext).
Thehistogramshowstheradialdistributionofthesurfacebright-
nessmeasuredinannularregionscenteredonthepointsource. The
backgroundlevelof0.15ctsarcsec−2hasbeensubtractedfromthe
data. The dashed linewithout errorbarsshows possiblecontribu-
tionofadusthalo(seeSection3.1.1andAppendix,equations A6
andA7).
TeV emission could arise either as a result of interac-
tion between the SNR shock and the GMC near W41
(Tian et al. 2007), or alternatively, it could be due to
Fig.2.— Top: Chandra ACIS 90′′×90′′ two-color (red: 0.5–2 a PWN powered by CXOU J183434.9–084443 (M+09).
keV; blue: 2–8 keV) image of the region around the pulsar and Below we discuss the constraints that our observation
compact PWN candidate coinciding with HESS 1834–078. The places on the nature of CXOU J183434.9–084443 and
imageisbinnedtoapixelsizeof1′′andsmoothedwithaGaussian HESS J1834–087. We also explore a possible link be-
ofFWHM∼4′′. Bottom: 2MASSimageofthe sameregion. The
tween these sources and the recently discovered Fermi
positionofthepulsarcandidateismarkedbyacirclewitharadius
of5′′. LAT source 1FGL J1834.3–0842c which is also located
within SNR W41.
3. DISCUSSION
3.1. CXOU J183434.9–084443 and its extended
Chandra resolved the compact extended emission sur-
emission
rounding CXOU J183434.9–084443, which is by far the
brightest X-ray source located within the extent of The point source appears to have a strongly absorbed
HESS J1834–087. It was previously suggested that the spectrum that fits hard PL. The flux, PL slope, and
6 Misanovic, Kargaltsev & Pavlov
2’ nts s keV−1 −1 102×10−3−3
u
o
C 0−4
1
×
5
2
χ 0
2
−
5
Energy, keV
2 - 8 keV
V−1 10−3
e
k
s −1
ounts 10−4
C
2
1
χ 0
1
−
2
−
2 5
Energy, keV
Fig.5.— Top: The X-ray spectrum of the pulsar candidate
extracted from the Chandra observation, and the corresponding
best-fit absorbed PL model. The absorption column is frozen to
its best fit-value (see Table 1) determined from the fit combining
20 cm
theChandraandXMM-Newtondata (seetext). Bottom: TheX-
rayspectrum of the PWN candidate extracted from the Chandra
dataandthecorrepondingabsorbedPLmodel,assumingthesame
absorptioncolumnasforthepulsarcandidate (Table1).
SGR 1833-0832
11
SGR J1833-0832
15
1FGL J1834.3-0842c
HESS J1834-087
Fig.6.—90%and99%confidencecontoursandtheunabsorbed
FWFHigM.∼4.—10′′,ToCph:andLraargAe-CscISaleim, asgmeo(o2t–h8edkeVw)ithencaomGpaaussssiniagnthoef aflnudxePsW(0N.5–c8ankdeiVd)atien.units of10−12 ergscm−2 s−1 forthe pulsar
region of the extended tail-like emission detected with XMM-
Newton (see Figure 1). The arrow shows the point source with
abright2MASScounterpart(seetext). Middle: VLAimageofthe the lack or any counterparts at lower frequencies make
sameregion. Theenhanced extended radioemissioninthe center CXOU J183434.9–084443 a plausible candidate for a
ofW41doesnotseemtocoincidewiththediffuseX-rayemission. young remote pulsar. However, other options, such as
Bottom: Atwo-colorimage(red: 20cm;blue: 0.5-8keV)showing
anAGNseenthroughthe Galacticplane,oraquiescent,
thewholeW41andalsotheprojectedACISchipsactivatedinthe
Chandraobservation. Theextent ofthe TeV source ismarkedby obscured XRB (similar to those discovered with INTE-
thebigcircle(radius5.′4)whiletheGeVsourceisshownbyits95% GRAL), remain viable. Although the latter interpreta-
positional error ellipse (6.′0×7.′2). Several X-ray sources detected
tioncouldexplainthe observedlargeN bytheintrinsic
in the FoV are marked, including the star candidate at the edge H
oftheGeVemission(smallcircle),andsources11and15asnum- absorption, we do not find any evidence of variability
bered in Table 1 by M+09. The position of the SGR J1833–0832
isalsomarkedwithasmallcircle.
Chandra observation of J1834-087 7
or an iron line which are commonly seen in obscured contribution of a dust halo. Although the dust halo in-
HMXBs (Walter et al. 2006). It seems rather plausible terpretation of the extended emission does not rule out
that a large amount of molecular material is generally theyoungpulsaroptionforthepointsource,itallowsfor
presentalongthisline ofsight;forinstance,G¨og˘u¨¸s et al. additional possibilities, such as a magnetar, an ANG, or
(2010)foundN ≃1023 cm−2 whilefittingthespectrum an XRB.
H
of the nearby SGR 1833–0832in the active state.
An important piece of evidence, which we will con- 3.1.2. A pulsar-wind nebula?
siderindetail,isacompact,r .20′′,extendedemission,
which has been clearly resolved from the point source Let us now assume that CXOU J183434.9–084443 is
(Figures 2 and 3). This emission appears to be approx- a pulsar, and a substantial fraction of the observed ex-
imately symmetric and exhibits a much softer spectrum tended emission around it is a PWN. The symmetric
compared to that of the point source (Γ = 2.7±0.5 vs. morphology of the candidate compact PWN would then
1.1±0.4). Two most likely interpretations are a PWN suggestthattheputativepulsardoesnotmoveveryfast.
arounda young pulsar in W41, or a dust scattering halo Even if the pulsar was born at the geometrical center of
around the strongly absorbed point source. A combina- the SNR (i.e., ∼ 2′ from its current position) 100 kyrs
tion of the two is also a possibility. In fact, a hint of a ago, its transverse speed would only be 22d4 km s−1,
plateau(ornon-monotonicbehavior)seeninFigure3be- which is in agreement with the nearly isotropic PWN
tween 3′′–10′′ may indicate at two possible components shape. Thus, the large-scale emission east of the pulsar
with different dependence on r. isnotexpectedtobeakintolongcollimatedtailsformed
behind supersonically moving pulsars, such as observed
3.1.1. A dust halo around a point source? by Kargaltsev et al. (2008). Indeed, the high-resolution
The very large absorbing column measured in X-rays X-rayimages ofthose tails show that the surface bright-
suggestsasignificantdustcolumnthatshouldleadtoan ness is usually the highest near the pulsar,and it gradu-
extended dust scattering halo (e.g., Predehl & Schmitt allydecreaseswith the distancefromthe pulsar,i.e., the
1995). The halo brightness grows with the increas- opposite of what we see in this case. It could be, how-
ing intervening dust column, which is usually as- ever, that the large-scale X-ray emission is akin to that
sumed to be proportional to the hydrogen absorp- in the Vela X PWN, where the X-ray emission is offset
tion column, N , measured from X-ray spectra. from the pulsar and thought to come form the relic pul-
H
More specifically, the scattering optical depth τ ≃ sar wind crushed and pushed aside by the reverse SNR
scat
S(N /1022 cm−2)(E/1 keV)−2 for S ≃ 0.5 found by shock (e.g., LaMassa et al. 2008).
H
Predehl & Schmitt (1995) gives τ < 1 at E > 3.7 The slope of of the extended emission spectrum, fit-
scat
keV for CXOU J183434.9–084443, whose spectrum vir- ted with the absorbed PL model, is rather steep, Γ =
tually cuts off below 3 keV (see Figure 5). Although 2.7±0.8,albeituncertain. ForaPWN,suchasteepslope
taking into account multiple scatterings would be more is unusual because the X-ray spectra of most7 PWNe
accurate,the largeuncertaintiesof the other parameters haveΓ=1−2(e.g.,seeFigure6inKargaltsev & Pavlov
(such as the dust grainproperties and dust distribution; 2008). The steep slope may, however, be indicative of
see Appendix) do not warrant this extra complication. strong synchrotron cooling. The ratio of the extended
Therefore, we have used the single scattering approxi- and point source luminosities, LPWN/LPSR ∼ 1.8, is
mation to calculate the halo profile by convolving the typical of PWNe (see Figure 5 in Kargaltsev & Pavlov
spectralintensity of a halo with the detector response in 2008). The spin-downpowerofthe putative pulsarpow-
the 3–8 keV energy range. We find that for the parame- ering the PWN can be estimated as E˙ = 4πr2cp ,
ters Θ = 360′′ and S ≃ 1, and for the dust distribution where r is the terminationshock radius and p s isamthbe
s amb
function f(x) defined in the Appendix (equation A8), ambient pressure. In the Sedov expansion phase, the
the dust halo model generally describes the observedra- pressure inside the SNR could be estimated from its ra-
dialprofile (see Figure 3). However,some deviations are dius by using a simple formula p = 3E/4πR3, where
amb
noticeable, hinting at a possibility of a second emission itisusuallyassumedthatthetotalSNRexplosionenergy
component. The dust distribution function f(x), used E =1051 ergs(e.g.,see O’C. Drury et al.2009). The es-
in the calculation of the halo model shown in Figure 3, timated pressure pamb,−9 = pamb/10−9 dyne cm−2 and
suggeststhatatleastsomedustmustbelocatednearthe the termination shock radius scaled to a plausible value
source,inagreementwiththeextremeabsorption,which r = r /1017 cm (corresponding to 2.7′′ at the W41
s,17 s
ckonuolwdnbteoaettxriisbtuitnedthtiostrheegilooncaolfmskoylec(uAllabrecrltoeutdsalt.h2a0t0a6r)e. distanceofd4 =4kpc)giveE˙ =4×1036rs2,17pamb,−9 erg
s−1 – a value typical for a young Vela-like pulsar. This
We should, however, point out that, in addition to
estimate is sensitiveto the value of r , whichcould be
the freedom of choosing the functional dependence of s,17
dustdistributionwiththedistance,theapproximatedust a factor of a few smaller then the scaling chosen8. The
scattering model we use has two other free parameters luminosityofthe compactPWN, ≃4.1×1033d24 ergs−1,
(S and Θ), which can attain values within rather broad could also be used to estimate E˙ from the L –E˙ corre-
X
ranges,dependingontheunknownpropertiesofthedust
grains along this particular line of sight (see Appendix). 7 A recently discovered PWN candidate in HESS J1632–408
Therefore, the mere fact that the model qualitatively (Balboetal.2010)alsoshowsasteepspectrumwithΓ=3.4+−00..68.
describes the observed radial profile does not guarantee However, the authors did not consider a possible contribution of
a dust halo, which may be quite substantial given the very large
thattheextendedemissionisindeedadusthalo. Onthe NH≃1.3×1023 cm−2 theymeasuredinX-rays.
other hand, the symmetric shape and the softer spec- 8 Toriradiiinbright,well-resolvedPWNeexceedrs byafactor
trum of the extended emission support at least partial of2onaverage(Bambaetal.2010).
8 Misanovic, Kargaltsev & Pavlov
lation(seeKargaltsev & Pavlov2008,forrecentresults). highly speculative.
Although this correlationshows a very largescatter, the On the other hand, a failure to find pulsations from
plausible range for E˙ ∼1036−1037 erg s−1 is in a good CXOUJ183434.9–084443,downto a restrictive limit for
agreement with the above estimated value. a young pulsar, would strongly suggest that it is an un-
related source. In this case, one would be forced to con-
3.2. The nature of HESS J1834–087, 1FGL clude that there is no plausible candidate for a young
J1834.3–0842c, and CXOU J183434.9–084443 pulsar within W41. This would imply that the TeV
emission is most likely produced by the interaction of
HESS J1834–087 belongs to the growing group of
protonsacceleratedintheSNRshockwithaparticularly
Galactic TeV sources that lack firm classifications or
densematerialinthenearbymolecularcloud(Tian et al.
identificationswithlower-energycounterparts. Thereare
2007; Albert et al. 2006). Obtaining a firm evidence of
currently about 25 TeV sources in this group, of which
this would be of a greatinterest because so far there are
only five can be considered truly “dark” sources (i.e.,
veryfew TeV sources (HESS J1745–303A,HESS J1800–
those without any detected counterparts) while the rest
240AB,HESSJ1848–018)wheretheinteractionwiththe
of the sources are coincident with one or more lower-
nearbymolecularcloudisconsideredasapossible(albeit
energy sources of a known nature (SNRs, massive open
notfirmly established) mechanismofthe VHE emission.
stellar clusters, GMCs, X-ray binaries), which could, in
Since HESS J1834–087 is located well within the W41
principle, power the TeV emission (although there are
shell and has a significantly smaller diameter, it could
no clear-cut cases). Spatial coincidence with W41 puts
only be associated with a part of the shell projected
HESS J1834–087 in the latter category. As such a good
within the W41 interior on the sky. The offset between
positional match between a TeV source and a bright ra-
the TeV and 1FGL sources would be even more difficult
dio SNR is very unlikely to happen by chance, HESS
to explain than in the PWN interpretation. Also, one
J1834–087 must have physical relation to W41. This
would have to conclude that the large-scale X-ray emis-
naturally leaves only two options for the origin of HESS
sion(seeSection2.1.2)mustbeassociatedwiththeSNR
J1834–087, a relic PWN or an SNR shock interacting
interior, which would require rather high temperatures
with the GMC, both of which have been previously dis-
to explain the measured hard spectrum (M+09).
cussed (M+09, Tian et al. 2007).
Keepinginmindtheuncertaintyintheinterpretations,
Our detection of extended emission around
weopttoplotthemultiwavelengthPWNspectralenergy
CXOU J183434.9–084443 does not provide a con-
distribution (Figure 7) but refrain from any modeling,
vincing argument in favor of either of these two possible
which would be premature due to the large number of
interpretations. Proving that CXOU J183434.9–084443
optionsaffordedbythe existingdata. Wenote,however,
is a pulsar (e.g., by detecting pulsations) would provide
that the GeV and TeV spectra generallyseem to match,
a very strong support to the relic PWN hypothesis.
which hints at their common origin and a possible inac-
Despite our highly significant detection of extended
curacy in the determination of the 1FGL J1834.3–0842c
emission around CXOU J183434.9–084443,it is difficult
position. Alternatively, the observed GeV emission co-
to unequivocally establish its origin because of the very
inciding with the part of the shell could be produced by
high absorption and a high likelihood of a significant
protons interacting with hadrons in the SNR shell and
dust scattering halo. It appears that a better way to
producingneutralpionsthatdecayemittingtheobserved
go about establishing the nature of CXOU J183434.9–
γ-rays. Therelativisticprotonsmaybeacceleratedeither
084443 would be a sensitive timing search in X-rays
in the PW (e.g., Arons 2009; Bednarek 2007) or by the
and GeV. However, even if CXOU J183434.9–084443
SNR shock interacting with the molecular clouds, which
turns out to be a pulsar, there will remain a question
has been also previously proposed as the source of the
about the origin of the nearby Fermi LAT source 1FGL
TeV emission (Tian et al. 2007; Albert et al. 2006).
J1834.3–0842c (see Figure 4), which is located within
W41 but offset from both the center of the TeV source 4. CONCLUSIONANDSUMMARY
and from CXOU J183434.9–084443, and even more
OurChandra observationresolvedtheextendedX-ray
offset from the patch of the large-scale X-ray emission,
emission in the immediate vicinity of CXOU J183434.9–
which could be a relic PWN akin to the Vela X PWN
084443, by far the brightest X-ray source within the
(see above). Unless the 1FGL source is not real9 or its
HESS J1834–087 extent. The source and the accompa-
position is inaccurate, a separate interpretation would
nyingextendedemissionarestronglyabsorbedinX-rays
be required. While awaiting for better Fermi image, we
leadingtosignificantuncertaintiesintheirspectra,which
can speculate that 1FGL J1834.3–0842c might be able
fit PL models with quite different slopes. The spectrum
to provide a long-soughtevidence of hadrons in a pulsar
of the extended emission is markedly softer suggesting
wind (e.g., Arons 2009; Bednarek 2007). Indeed, the
thattheemissionmightbeadustscatteringhalo. Alter-
positional coincidence with the SNR suggests that the
natively, the difference in the slopes could be due to the
GeV emission can be produced by pulsar-wind protons
strong synchrotron cooling, if the extended emission is
interacting with the dense material of the part of the
a PWN. Although our approximate dust scattering halo
SNRshell,whichintriguinglyturnsouttobethenearest
calculations allow us to obtain a reasonable fit to the
to the putative pulsar. We must admit, however, that
observed radial profile of the extended emission, large
until the pulsar nature of CXOU J183434.9–084443
uncertainties in the dust scattering model preclude us
is firmly established the above reasoning will remain
from definitive judgment on the nature of the extended
emission,leaving a PWN optionas a still viable alterna-
9 Itismarkedas“confused”inthe1FGLcatalog, whichmeans
tive. The previously reported faint, large-scaleextended
that it might be the result of an imperfect model of the diffuse
Galacticbackground(seeAbdoetal.2010,fordetails). emissionappearstobedisjoinedfromCXOUJ183434.9–
Chandra observation of J1834-087 9
complicates the interpretation of the GeV emission sug-
gesting that either the 1FGL J1834.3-0842c’sposition is
inaccurate or it could be an intriguing case of a pul-
sar wind interacting with the SNR shell via hadronic
mechanism. The X-ray timing and deeper imaging
observations are required to understand the nature of
CXOUJ183434.9–084443andofthe offsetlarge-scaleX-
ray emission. Together with improved GeV data, this
will make it possible to identify HESS J1834–087 and
1FGL J1834.3-0842c, and determine the nature of their
relationship with W41.
We also serendipitously observed the recently discov-
eredSGRJ1833–0832,which,however,wasnotdetected.
We placealimit ofof3×10−14 ergss−1 cm−2 onits un-
absorbed flux, which is a factor of 40 dimmer than was
measured by G¨og˘u¨¸s et al. (2010).
Support for this work was provided by the National
Fig.7.— Broad-bandspectrum of the PWN candidate showing
theX-rayspectrum(ACIS,black)andthecataloguedGeV(Fermi, Aeronautics and Space Administration through Chan-
red) and TeV (HESS, blue) emission. The TeV nebula spatially dra Award Number GO9-0076X issued by the Chan-
overlaps withthe X-rayPWN candidate, whilethe GeV emission dra X-ray Observatory Center, which is operated by
islocated between theTeV nebulaandpartoftheSNRshell(see
the Smithsonian Astrophysical Observatory for and on
Figure4).
behalf of the National Aeronautics Space Administra-
084443inabetterresolutionChandraimage,whichrules tion under contract NAS8-03060. The work was also
outapulsartailhypothesisandsuggestthatit mightei- partially supported by NASA grants NNX09AC84G
therbe apartofadisplacedrelic PWNora surprisingly and NNX09AC81G, and NSF grants No. 0908733 and
hot plasma in the SNR interior. 0908611. G. G. P was partly supported by the Ministry
Adjacent to the TeV source, and overlapping with a ofEducationandScienceoftheRussianFederation(con-
part of the SNR shell, there is extended GeV emission, tract 11.634.31.0001).
the Fermi LAT source 1FGL J1834.3-0842c. The off-
setfromHESS J1834–087andCXOUJ183434.9–084443
APPENDIX
DUSTHALO MODEL
Dust halos, often seen around bright point-like X-ray objects, are formed by scattering of source X-ray photons on
dust grains. Here we will only discuss the case of dust optically thin with respect to the photon scattering, τ .1,
scat
andconsideronlyazimuthallysymmetrichalos(whichimplies thatthedustdistributionacrosstheline ofsight(LOS)
is uniform within the interval of angles θ at which we see the halo). In this case the spectral halo intensity (phcm−2
s−1 keV−1 arcmin−2) is given by the equation
1 f(x) dσ (E,θ )
s s
I (θ,E)=F(E)N dx , (A1)
halo H x2 dΩ
Z0 s
where F(E) is the point source spectral flux (photonscm−2 s−1 keV−1), x=(D−d)/D is the dimensionless distance
from the X-ray source to the scatterer (D and d are the distances from the observer to the source and the scatterer,
respectively), θ ≃ θ/x (for small angles) is the scattering angle, f(x) is the dimensionless dust density distribution
s
1
alongthe LOS( f(x)dx=1),anddσ (E,θ )/dΩ isthe differentialscatteringcrosssectionper onehydrogenatom,
0 s s s
averagedover the dust grain distribution over sizes and other grain properties (see, e.g., Mathis & Lee 1991).
R
To understand the halo properties from simple analytical expressions, we will use the Rayleigh-Gans (RG) ap-
proximation, in which the total scattering cross section ∝ E−2; this approximation works better for higher energies,
E &0.5–2keV,depending onthe dust model. For some dustmodels, the averageddifferential crosssectionin the RG
can be approximated as (Draine 2003)
dσ (E,θ ) σ (E) 1
s s ≈ s , (A2)
dΩ πθ2 (1+θ2/θ2 )2
s s,50 s s,50
where
Θ S
θ ≈ and σ (E)≈ 10−22cm2 (A3)
s,50 E s E2
arethemedianscatteringangleandthetotalcrosssection,respectively;E istheenergyinkeV.TheconstantΘinfirst
′′
eq. (A3) depends on the dust model; Draine (2003) derived Θ = 360 from the dust model of Weingartner & Draine
′
(2001), while Bocchino et al. (2010) found Θ=7.4 for the model of Smith & Dwek (1998).
10 Misanovic, Kargaltsev & Pavlov
It follows from the second eq. (A3) that the scattering optical depth is
τ (E)≈SN E−2. (A4)
scat H,22
The factor S in the second eq. (A3) is a constant of the order of 1; e.g., S ≈ 1.3 from Figure 6 of Draine (2003),
while Predehl & Schmitt (1995) found a mean value S ≈ 0.49 for a number of halos observed with ROSAT (but the
scatter was very large), while Mathis & Lee (1991) discuss models with S =0.903, 1.09, and 0.47 (see their Table 1).
Costantini (2004; PhD thesis) estimated τ (1keV) for a number of halo sources observed with Chandra; the values
sca
of S derived from her results show a very strong scatter, S from 0.018 to 2.26. The scatter itself may be natural, as
the dust properties may be different for different sources.
It should be noted that the correlationof τ (1keV) with visual extinction A :
sca V
τ (1 keV)=(0.056±0.01)A (A5)
scat V
(Predehl & Schmitt 1995) is better than that with N , but A is rarely known for the objects of interest.
H V
Substituting (A2) and (A3) in (A1), we obtain the spectral intensity profile
S 1 f(x) θE 2 −2
I (θ,E)=F(E)N 1+ dx. (A6)
halo H,22πΘ2 x2 xΘ
Z0 " (cid:18) (cid:19) #
For comparison with the point source + halo profile observed in the energy range E < E < E , the sum of the
1 2
spectralintensities shouldbe convolvedwiththe detector response,withallowanceforthe imagespreadcausedby the
telescope and the detector. We have checked that the energy redistribution in the detector only slightly affects the
broadband radial profile for a smooth incident spectrum. Therefore, assuming the observable halo size to be much
larger than the PSF width, we obtain
E2 S 1 f(x) θE 2 −2
I (θ)= dEA (E)F(E) ψ(θ,E)+N dx 1+ , (A7)
obs eff  H,22πΘ2 x2 xΘ 
ZE1  Z0 " (cid:18) (cid:19) # 
whereA (E)isthedetector’seffectivearea,andψ(θ,E)isthenormalizedPSF,whichcanbetakenfromasimulation
eff  
(e.g., with MARX). The first term in Equation(A7) corresponds to the point source, while the second term describes
the halo.
The above equation can be integrated for a given set of halo parameters, and compared directly with the data. In
′′
particular, for the dust halo model shown in Figure 3 we picked Θ=360 , S =1, and the dust distribution function
x−1, x≤x
f(x)= o o (A8)
0, x>x
o
(cid:26)
with x =0.25. According to (A6), this distribution function corresponds to
o
F(E)N Sa a
I (θ,E)= H,22 arctana−
halo 2π(x Θ)2 1+a2
o (cid:18) (cid:19)
≈ F(E)NH,22Sa 1 a≫1 (A9)
4x2Θ2 (3π)−1a3 a≪1
o (cid:26)
where a=x Θ/θE.
o
REFERENCES
Abdo,A.A.,etal.2010, ApJS,188,405 Gaensler,B.M.,&Johnston, S.1995,MNRAS,275,L73
Acero,F.,Ballet,J.,Decourchelle, A.,Lemoine-Goumard,M., Gallant,Y.A.,etal.2008, inAmericanInstituteofPhysics
Ortega,M.,Giacani,E.,Dubner,G.,&Cassam-Chena¨ı,G. ConferenceSeries,Vol.983,40YearsofPulsars:Millisecond
2009, A&A,505,157 Pulsars,MagnetarsandMore,ed.C.Bassa,Z.Wang,
Aharonian,F.,etal.2005, Science,307,1938 A.Cumming,&V.M.Kaspi,195
Aharonian,F.,etal.2006, ApJ,636,777 Gelbord,J.M.,etal.2010, GRBCoordinates Network,Circular
Albert,J.,etal.2006,ApJ,643,L53 Service,10526,1,526,1
Arons,J.2009,inAstrophysicsandSpaceScienceLibrary,Vol. Gelbord,J.M.,&Vetere,L.2010,GRBCoordinatesNetwork,
357,AstrophysicsandSpaceScienceLibrary,ed.W.Becker, CircularService,10531, 1,531,1
373 G¨o˘gu¨¸s, E.,etal.2010,ApJ,718,331
Balbo,M.,Saouter,P.,Walter,R.,Pavan, L.,Tramacere,A., Hinton,J.A.,&Hofmann,W.2009,ARA&A,47,523
Pohl,M.,&Zurita-Heras,J.2010, A&A,520,A111 Kargaltsev,O.,Misanovic,Z.,Pavlov,G.G.,Wong,J.A.,&
Bamba,A.,Mori,K.,&Shibata, S.2010,ApJ,709,507 Garmire,G.P.2008, ApJ,684,542
Bartko,H.,&Bednarek,W.2008,MNRAS,385,1105 Kargaltsev,O.,&Pavlov,G.G.2008, inAmericanInstitute of
Bednarek,W.2007,Ap&SS,309,179 PhysicsConferenceSeries,Vol.983,40YearsofPulsars:
Bocchino,F.,Bandiera,R.,&Gelfand,J.2010,A&A,520,A71 MillisecondPulsars,Magnetars andMore,ed.C.Bassa,
deJager,O.C.,&Djannati-Ata¨ı,A.2008,ArXiv0803.0116 Z.Wang, A.Cumming,&V.M.Kaspi,171
Dickey,J.M.,&Lockman,F.J.1990,ARA&A,28,215
Draine,B.T.2003,ApJ,598,1026