Far-ultraviolet Observations of the North Ecliptic Pole with SPEAR

0.26 MB·English

by  Eric J. Korpela

Accepted forpublicationin ApJ Letters on 19December2005
PreprinttypesetusingLATEXstyleemulateapjv.6/22/04
FAR-ULTRAVIOLET OBSERVATIONS OF THE NORTH ECLIPTIC POLE WITH SPEAR
Eric J. Korpela1, Jerry Edelstein1, Julia Kregenow1, Kaori Nishikida1, Kyoung-Wook Min2, Dae-Hee Lee2,3,
Kwangsun Ryu2, Wonyong Han3, Uk-Won Nam3, Jang-Hyun Park3
Accepted for publication in ApJ Letterson 19 December 2005
ABSTRACT
We present SPEAR/FIMS far-ultraviolet observations near the North Ecliptic Pole. This area, at
◦
b∼30 andwithintermediateHicolumn,seemstobeafairlytypicallineofsightthatisrepresentative
of general processes in the diffuse ISM. We detect a surprising number of emission lines of many
elements at variousionizationstates representinggas phases from the warmneutral medium (WNM)
6
to the hotionizedmedium (HIM). We alsodetect fluorescencebands ofH , whichmaybe due to the
0 2
ubiquitous diffuse H previously observed in absorption.
0 2
2 Subject headings: ISM: general, ISM: lines and bands, ultraviolet: ISM

n
Ja 1. INTRODUCTION ss−−11..)(IBneccoamuspearsitsaorns,atratynpsiitcatlhpehointsotnrucmouennttrsalitteiins ∼<∼250
The North Ecliptic Pole (NEP) is a region of the sky
6 s during survey sweeps, this results in minimal loss of
with no obviously unusual features. Many prior space-
2
observing time. The area of sky lost due to a star is
missions (Einstein, ROSAT, COBE, IRAS, etc.) have
similar to the instrument imaging resolution across the
1 conducted deep surveys of this region. Several ground ′ ′
slit width (∼10 ×5). See Edelstein et al. (2005a) for a
v based surveys have also concentrated on this region as
description of the data reduction pipeline.
3 a complement to the space-based surveys (Labov et al.
To further remove stellar contamination from our
8 1989; Elvis et al. 1994). Its location at moderate galac-
5 tic latitude (b=+29◦) places it clear of the bright stars dataset, we excluded locally intense spatial bins (∼30%
1 of the galactic plane. In this region, the Galactic of the viewed area), defined to have >∼3× the median
0 N(HI) varies from about 2×1020 cm−2 to 8×1020 cm−2 count rate. This level corresponds with a slope change
6 in the histogram of the log count rate, indicative of a
which corresponds to an dust opacity of τ =0.6-
dust
0 2.3 at 1032˚A(Sasseenet al. 2002) or τ =0.3-1.4 at difference in distribution of scattered vs direct starlight.
/ dust Inclusion of unresolved starlight should not be a factor
h 1550˚A(Savage & Mathis 1979).
other than to raise emission-line determination errors.
p The Spectroscopy of Plasma Evolution from Astro-
Toobtainthehighestsensitivitypossible,wehavebinned
- physical Radiation (SPEAR) instrument, designed for
o the entire dataset into a single spectrum including a to-
observing emission lines from the diffuse ISM was
r tal of 3.5×105 counts over an effective full-slit exposure
t launched in late 2003. SPEAR is a dual-channel FUV
s of 13.5 ks.
a imaging spectrograph (short-λ (S) channel 900 - 1150
v: ˚A, long-λ (L) channel 1350 - 1◦750 ˚A) with λ◦/∆λ ∼550, 3. SPECTRALMODELING
withalargefieldofview(S:4.0 ×4.6’,L:7.5 ×4.3’)im-
i
X aged at 10′ resolution. See Edelstein et al. (2005a,b) for The raw SPEAR spectrum is composed of many com-
r an overview of the instrument, mission, and data anal- ponents:detectorbackground,dust-scatteredstellarcon-
a yses. SPEAR sky-survey observations consist of sweeps tinua,directandinstrumentallyscatteredairglow,andIS
atconstanteclipticlongitudefromtheNEPtothesouth atomicandmolecularemission. Tomodelaspectrumwe
ecliptic pole through the anti-solar point. The duration simultaneously fit all of the components to the observed
of each sweep varies between 900 and 1500 s. Each sur- spectrum.
vey scan includes the region near the north and south The spectral shape is distorted by the detector elec-
ecliptic poles resulting in large exposure times in these tronics and by the data processing pipeline. These
regions. WereportondeepSPEARobservationsofFUV distortions appear as fluctuations in a flat-field image.
emissionspectrafroma15◦radiusregioncenteredonthe Any broad-band component used in our model spec-
NEP. trum must be multiplied by this nonlinearity function.
(Korpela et al. 2003; Rhee et al. 2002).
2. OBSERVATIONSANDDATAREDUCTION BothSandLareaffectedbyinstrumentalbackground
The NEP data analyzed here include sky-survey and duetocosmicrays,radioactivedecaywithinthedetector,
pointed observationsthat occurredbetween 8 November and thermal charged particles entering the instrument.
2003 to 10 November 2004. To remove times of high Thesesourcesofbackgroundarerelativelyuniformacross
background and bright stars from the data set, we have the face of the detectors. Although the charged particle
excluded times during which the count rate exceeds 100 rate can change with time and orbital position, the po-
sitional distribution of these backgroundevents is domi-
1 SpaceSciences Laboratory,UniversityofCalifornia,Berkeley, nated by static fields within the instrument and is rela-
CA94720-7450; [email protected] tively constant. A sum of many shutter-closed dark ex-
2 KoreaAdvancedInstituteofScienceandTechnology,305-701, posuresmultiplied bya ratefactor canbe usedto model
Daejeon,Korea
3 Korea Astronomy and Space Science Institute, 305-348, Dae- this contribution to the spectrum.
jeon,Korea In L, the largest broad-band spectral component is
2 Korpela et al.
Fig. 1.— SchannelH2 andlineemissionmodels. Upperpanel: bestfitH2 fluorescencemodelisshownagainsttheresidualsfollowing
subtraction of the continuum/background model. The H2 features’ positions are well predicted, but their ratios are not. The H2 is
oversubtractednear1100˚AandundersubtractedshortwardofLyβ λ1026. Featuresatλ<1000˚Aareunderpredicted. LowerPanel: The
best fit atomiclinespectrum is plotted against the residualsfollowingsubtraction of the H2 spectrum. Probable H2 features and several
asyetunidentifiedfeaturesremainthroughouttheband.
dustscatteredstellarcontinua. Wemodeltheinputspec- pletely photoionized in the WNM (ionization potential
trum to the dust scattering process as a sum of spectra (IP) <13.6 eV). Forspecieswhicharelikelytobeabun-
for upper main sequence (UMS) stars. Each spectrum dantintheWNMandwarmionizedmedium(WIM),(IP
is weighted by using a power law UMS luminosity func- >13.6 eV) we assume that resonance lines are optically
∼
tion (dN/dMv ∝10αMv, where Mv is the absolute mag- thick (τ >10) and that the distances to the illuminating
nitude and dN/dMv is the volume density of stars per stars are large compared to the τ = 1 surface. In cases
unit magnitude) and a dilution factor dependent upon where the ground state is split, absorbed resonance line
the weighted mean distance of the stars. The spectral photons will be converted to the excited state, greatly
sum is multiplied by a dust opacity scaled to a variable changing the relative contributions in the multiplet. We
N(HI)andbyadustalbedofunction(Draine2003)with modelthisbyattenuatingthe groundstatetransitionby
a variable total dust surface area. The best fits for this a large factor. We will develop a more accurate way of
function typically occurvery nearto the measuredUMS treating these optical depth effects in future work.
luminosity function with α ∼ 0.2 (Reed 2001). Because
4. DISCUSSION
the stellar features in this spectrum cannot be directly
measured, there may be some bias in our measurements The spectrum of the bright Lyman-series geocoronal
ofthe overlyingresonancelines (Civ λλ 1548,1551,Ovi emission,easily visible in Fig. 1, can be modeled using a
λλ 1032,1038, Siivλλ 1394,1403, etc.) of order ∼ 1000 temperatureandtotallineflux. ThelargestSbroadband
ph s−1 cm−2 sr−1 (LU). We include this effect in our componentisscatteredgeocoronalLyαλ1216emission.
error analysis. Our best fit continuum, which could in- This line occurs in the bandpass gap between S and L.
clude unresolved stars, is 300±140 ph s−1 cm−2 sr−1- L is immune to Lyα contamination because it includes
50
˚A−1 (CU) in L. In S, due to its lower sensitivity and a CaF2 filter which is very nearly opaque to Lyα radi-
ation. We fit the S scattering with two components, an
highernon-astrophysicalbackground,weareonlyableto
exponentialscatter∝exp(−|λ−λ |/α)withα∼200˚A,
placeastatisticalupperlimitof1200CU.However,given 0
and an isotropic component. The scattering, summed
thewavelengthdependenceofthedustalbedo,weexpect
over the entire spectrum, is ∼ 0.08% of a ∼5 kRy Lyα
thatthetrueSastrophysicalcontinuumisnohigherthan
incident flux. Because the Lyα transition is optically
in L.
thick, and because the scattering fractions are indepen-
We firstattemptedtofitatomicline emissionby using
dent parameters,we fix the estimated incident Lyα flux
linear combinations of collisional ionization equilibrium
at 1100× the Lyβ flux which helps the solution to con-
(CIE) plasma models with a limited abundance gradi-
verge. Whilethefractionalscatterandexponentialwidth
ent vs T. We attribute this method’s lack of success to
of these components are free parameters in our spectral
photoexcitation of resonance lines, non-CIE effects and
model, the values we obtain vary little from one dataset
abundance variations that are not T related. To fit the
to the next.
line emission we have, instead, fit the spectrum of each
Our model includes spectral features of atmospheric
species separatelywith a per-species intensity andT pa-
Oi, Ci, and Ni that are much fainter than when seen in
rameter. The emission line spectra are convolved with
day-time observations (e.g. FUSE), as SPEAR only ob-
a λ-dependent Gaussian to model the instrument spec-
servesduringorbitaleclipse. Wedetectaspectralfeature
tral resolution. We use the atomic parameters of CHI-
near 990 ˚A that is consistent with the Oi λ 989 transi-
ANTI (Young et al. 2003) to determine the spectrum
tion. However, this feature may be blended with Niii λ
of the species vs T. This spectrum is then normalized
990. If the entirety of the 990 ˚A feature is due to Oi,
and multiplied by the intensity parameter. When CHI-
the lackof detectable Oiλ 1356emission wouldindicate
ANTI parameters are not available for a species, we use
T >103.6 K,whichisplausibleforinterstellarOi. Wedo
a spectrum calculated using the NIST Atomic Spectra
detect Oi λ 1356 in data that is not count rate filtered.
Database.
The other possibility is that the 990 ˚A is due to Oi at
We thenconsider the effects ofself-absorptionofspec-
T ∼ few hundred K whichseems unlikely evenfor atmo-
tralresonancelines. In the caseofhighly ionizedspecies
spheric Oi. The most likely explanation for the lack of
we assume all lines are optically thin. We also as-
any detection of Oi emission at λ6=990˚Ais that a large
sume this to be the case for species that would be com-
portion, if not all, of the 990 ˚A emission is due to Niii.
SPEAR observations of the NEP 3
Fig. 2.— Lchannel H2 andlineemissionmodels.
We detect relatively strong emission from H fluores-
2
cence bands even in this region of intermediate N(HI). TABLE 1
Observedwarm/hotISM emission lineintensities
This is not entirely surprising since absorption due to
H is a ubiquitous feature of FUV spectra taken by
2
FUSE. (Shull et al. 2000) We fit these features by using Species λ Iλ Statistical Model
Error Uncertainty
interpolation between of a sparse grid of H2 models (B. ˚A LUa LU LU
Draine 2004,private communication). The upper panels
ofFigs.1&2showthebestfitH modelplottedagainst Oiii] 1667 <500b – –
2 Oiv] 1400 1980 ±220 ±300
theresidualsofthespectrumfollowingbackground,con- Ovi 1032,1038 5724 ±570 ±1100
tinuum, and scattering subtraction. We detect H2 emis- Ciii 977 1260 ±270 –
sion features in both spectral bands. In S we detect the Civ 1548,1551 5820 ±280 +−01490
H2 λλ 965 band (to the left of the HI λ 972 line) and Siii∗ 1532 2430 ±200 +−6030
the H λλ 986 band (to the left of the Oi,Niii λλ 989 Siiv 1394,1403 1430 ±120 ±160
2
feature). The H λλ 1053 and λλ 1097 bands are not Nii 1085 8500 ±730 –
2 Niii 990 <2300 – –
detected atanysignificance. In L,H2 providescontribu- Niv 1485 1984 ±170 ±450
tions throughout the spectrum. We attribute the triple Alii 1671 5608 ±525 –
peaked band structure between 1350 and 1380 ˚A, the Heii 1640 1850 ±180 +−0500
peak near 1440 ˚A, the broad features λλ 1520-1530, λλ Nevi 999,1005 <4100 – –
1570-1590, and the peak near 1608˚A to H . Note that
2
these features are not well fit by our H model. It is Note. — aLU=ph s−1 cm−2 sr−1 bUpper limits are 90%
2
not surprising that our sparse linear model with a sin- confidence
gle dust absorption column does not fit these data since
the observed emission derives from many lines of sight
withvaryingillumination,gastemperaturesandabsorb- at 1533 ˚A to the Siii∗ λ 1533.4 transition into the up-
ing columns. In Figs. 1 & 2, lower panels we show the
per fine structure level of the ground state. The offset
residual spectrum following subtraction of the H com-
2 betweenthe model featureandthe observedline may be
ponent.
duetoanuncorrectedcontinuumfeatureunderlyingthis
We see intriguing indications of N fluorescent emis-
2 line (Fig. 3 of Edelstein et al. (2005a)). Since Siii is the
sion,mostnotablythefeaturesat1434˚A,1445˚A,1475˚A,
predominantstateofsiliconinboththeWNMandWIM,
and 1660˚A. FUSE has detected absorption due to N
2 the true Siii λ 1526.7ground state transition is likely to
along some sight-lines, albeit at shorter wavelengths
be optically thick and therefore undetectable. Our spec-
(Knauth et al. 2004). We cannot yet discount the pos-
tral models are improved by the addition of Feii transi-
sibility that these N features are atmospheric, as near
2 tions throughout band, most notably the Feii λλ 1564,
the magnetic poles, N+ can be transported to high alti-
2 whichfillsthegapbetweenCivλλ1548,1551andtheλλ
tude. Further study to determine whether the N emis-
2 1570-1590H band. Itisalsopossiblethatthereisacon-
sionvarieswithspacecraftorbitalpositionorwithgalac- 2
tribution to the H λλ 1600-1620 band due to the Feii
tic latitude of the observation should resolve this ques- 2
λλ 1608 resonance lines. These lines have a fairly low
tion. If confirmed as astrophysical, this would be the
oscillator strength and are expected to be optically thin
firstdetectionoffluorescentN emissionfromthediffuse
2 alongthisline ofsight. However,giventhelowoscillator
ISM.
strength, we expect these lines to be faint compared to
SPEAR has discovered some surprising spectral fea-
the Siii and Alii features.
tures, such as the resonance lines of the singly ionized
Emissionlinesofhighlyionizedgasexpectedtobegen-
species, Siii and Alii. The Alii feature at 1671 ˚A very
erated in a CIE plasma at HIM temperatures are shown
nearly overlies the Oiii] λλ 1665 which we do not de-
inTable1. Ofthesewedetect(>3σ)Ciiiλ977,Civλλ
tect in this dataset. In a lower resolution instrument
1548,1551,Siivλλ1394,1403andOviλλ1032,1038. The
thesespectralfeaturescouldbeblendedresultingincon-
Civλλ1549andSiivλλ1400emissionoverlieprominent
fusion of the Oiii] and Alii emission. We think it is
stellar absorption features which are likely to be present
likely that at least some of the prior claimed detections
in the dust-scattered stellar radiation. Since the depth
(Martin & Bowyer 1990) of Oiii] were in fact misidenti-
of these features is unknown we set the lower range of
fied detections of Alii λλ 1671. We attribute the feature
our model uncertainty to be the intensity that would be
4 Korpela et al.
calculated if there were no continuum feature present. at these temperatures. Some portion of the difference
Many of these emission features are resonance lines could be due to abundance variations (with ISM deple-
(R. J. Reynolds 2005, private communication) and, be- tions of O and C higher than those of Ne, He, and Si),
cause of the substantial IS radiation field at these wave- although the suggested depletions are much higher than
lengths, further work will be required to determine
whether these lines arise due to collisional excitation in
theHIMorduetoresonantscatteringoftheISradiation
field. We believe that the short wavelengthlines are less
likely to be resonant scatter because the relative inten-
sity of the scattered stellar continuum in the S band is
much lower (Korpela et al. 1998).
Our Civ λλ 1548,1551 intensity of 5820±280 LU is
marginally higher than values for the diffuse ISM re-
ported by Martin & Bowyer (1990). However Martin
andBowyerassumedafeaturelesscontinuumratherthan
a continuum with stellar features. If we make the same
assumption our value drops to 4330±210 LU which is
more within the range of values they report.
Our Ovi λλ 1032,1038 intensity of 5724±1100 LU is Fig. 3.— Solar abundance emission measure (EM) vs T for
somewhat higher than the typical value measured by detectedemissionlinesandsignificantupperlimits. Thelinesrep-
FUSE (Shelton 2003). This could be a systematic ef- resentthelociofEMnecessarytoproducetheobservedintensities.
fect due to the uncertain calibration of the S band. In
facttheratiooftheH featureintensitiesbetweenLand
2
S suggests that we have underestimated the sensitivity
the accepted ISM values. In future work we will derive
of S. Since this is an average over a very large region, it
constraints to non-CIE cooling models and abundances
is also possible that the Ovi emission is due to several
using these line ratios.
small regions with very high intensity, such as SNR.
The double peaked feature near 1400 ˚A is a blend of 5. CONCLUSIONS
Siiv λλ 1393.8,1402.8and Oiv] λλ 1400.7,1407.4. The ◦
We have presented the FUV spectra of a 30 region
Oiv] emission is of great interest because, as a semi-
around the NEP. We detect a variety of emission lines
forbidden line, it is not strongly affected by the radi-
from high and low ionization gas. The high ionization
ation field, but is a direct indicator of collisional pro-
lines, taken on a per species basis, are consistent with
cesses. Assuminga 2:1ratiofor the Siiv doubletcompo- EMfrom0.001to0.005cm−6pc throughoutT =104.5to
nents, our model determines the Oiv] doublet intensity 105.5 Krange. Forthosefewspecieswhichhavemultiple
to be 1980±300LU and the Siiv doublet intensity to be
lines in the band, the calculated temperatures tend to
1430±160LU.
fallneartheCIEpeakabundanceforthespecies. Despite
We detect high stage ions of two noble gases in our
this,alinearcombinationofCIEmodelswitharestricted
spectra,Heiiλλ 1640andNeviλλ 999,1005. Although
abundance vs T gradient does not provide a good fit
the features corresponding to the Nevi lines are statisti-
to the spectrum. This is likely due to non-CIE effects
callysignificant,thepoorfittotheunderlyingH bands
2 and abundance variations are not directly T related. It
prevents us from placing more than an upper limit to
is also likely that many lines include some resonantly
thesefeatures. BecauseofthelackofdetectionofanyOi
scattered stellar radiation. Modeling this scattering will
λ 1356,we do notexpect any contaminationofthe Heii
be a significant portion of our continued research.
linewithOiλ 1641. TheHeiilineismuchbrighterthan
We have discerned line emission from species that are
would be expected in a solar abundance CIE model (in
expected to exist in the WNM and WIM. We interpret
ratiotoCiiiandOiii). Thiscouldbe duetodepletionof
these lines as resonantly scattered stellar radiation. In
C and O to <0.1 solar. A depletion of this magnitude
∼ the future, we will investigate if spatial variation of the
would be significantly higher than the accepted values
Alii and Siii features can indicate the effect of grainfor-
for the ISM.
mation and destruction on the gas phase abundances of
InFig.3 weshowequivalentsolarabundanceemission
these elements.
measure(EM)andlimits for severalspecies. The results
DespiterelativelylowgalacticN(HI)inthisdirection,
are not entirely consistent between species. For exam-
weseeasubstantialamountofH fluorescence. Further
ple higher EM is necessary to explain the Siiv and Heii 2
modeling of the H fluorescence is necessary to fully
emission than are compatible with the amount of Ciii 2
understandfaintspectralfeatureswhichcouldberelated
emission. The questionable Nevi detections would indi-
to H or due to unrelated atomic lines.
cate EM 10× that suggested by the Ovi measurements. 2
The bulk of the inconsistency in Nevi is likely to be due
to undersubtraction of the underlying H . Much of the SPEAR/FIMS is a joint project of KASSI & KAIST
2
inconsistency in other lines is likely due to non-CIE ef- (Korea)andU.C., Berkeley(USA), fundedbythe Korea
fects which are present in gas being heated or cooling MOST and NASA Grant NAG5-5355.
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