星形成ゼミ 2014/1/17 Newsletter #252 p21-27 10本
1.Probing the turbulent ambipolar diffusion scale in molecular clouds with spectroscopy
Talayeh Hezareh1, Timea Csengeri1, Martin Houde2, Fabrice Herpin3 and Sylvain
Bontemps4
Massive dense coreのCygX-N03, CygX-N53に対して、Turbulent ambipolar diffusion scaleとMagnetic field strengthを PdBIのH13CN,H13CO+ (J=1-0)のline profileを用いて求めた。
結果として、turbulent ambipolar diffusion scale length が3.8±0.1 mpc、21.2±0.4 mpcで、
magnetic field strengthが0.33mG, 0.76mGを得た。
CygX-N53よりN03の方がSNが悪いデータであり、inclinationの効果が含まれていることもあるけれど、磁場強度の違いは２つのコアでfragmentationの効率の違いを意味しているだろう。
Probing the Turbulent Ambipolar Diffusion Scale in Molecular Clouds with Spectroscopy 12
11

T. Hezareh, T. Csengeri, M. Houde, F. Herpin, and S. Bontemps

CygX−N030.8

0.8
σ2(km/s)2

σ2(km/s)2

0.6
0.4
0.2
H13CN
H13CO+

0.0
3.0

CygX−N53

1.0

3.5

4.0

4.5

5.0

5.5

0.6
0.4
0.2
0.0
3.5

6.0

H13CN
H13CO+

4.0

4.5

5.0

5.5

6.0

0.30

0.8

0.20

0.6
σ2(km/s)2

σ2(km/s)2

0.25

0.15
0.10

.....

0.05
0.00
3.0

σ2n = 0.09L 0.52−−− σ2i = 0.09L 0.52 − 0.05

3.5

4.0
4.5
5.0
Length scale (arcsec)

Figure 4. Top: Plot of σ2 as a function of length scale in CygXN03. The measurement uncertainties for the data points are notshown to avoid clutter in the graph. Bottom: The lower envelopes of the H13 CN and H13 CO+ data fitted for Kolmogorovtype power laws.

5.5

0.4

0.2
6.0

..... σ2n = 0.24L 0.52−−− σ2i = 0.24L 0.52 − 0.34

0.0
3.5

4.0

4.5
5.0
Length scale (arcsec)

5.5

6.0

Figure 5. Same as Figure 4, but for Cyg-N53.

c 0000 RAS, MNRAS 000, 000–000!

c 0000 RAS, MNRAS 000, 000–000!

2.Imaging the circumstellar environment of the young T Tauri star SU Aurigae
S.V. Jeffers1,2, M. Min3,2, H. Canovas4, M. Rodenhuis2, and C.U. Keller2
William Herschel telescope(4.2m)の偏光撮像装置ExPoで、SU Aurのimageを取得した。
SU Aurigaeは、距離143pcにあるT tauriとHerbig starの境界にあるG2型星である。
今回の結果で得られたdiskサイズは500AU(先行研究で言われていた2-300AU)。このサイズはHerbigのdiskのmedian値で、T tauriの最大である。また、西の方向に広がった成分が確認された。すでに報告されているが今回の結果からこれまで言われていたjetのcavity
wallではなく、母体分子雲コアの残存物であると考えられる。

V. Jeffers et al.: Imaging SU Aur at the T Tauri/Herbig transition

marise the current knowldescribe the observationalth the polarimetric image
The observations are interlling in Section 5, and the
We present our conclusions

Table 1. SU Aur’s stellar parameters

Parameter
Reference
Coordinatesa RA
04 55 59.39
DEC+30 34 01.52
Spectral type
G2 III
1
MV
9.16 mag
2
B-V
+0.90
2
Effective temperature
5860±100K
3
8
S. V. Jeffers et al.: Imaging SU Aur at the T Tauri/Herbig transition
Mass
1.88±0.1 M&
3Radius
3.5 R&
4auri star (Herbig
1952;
Table
3. Values
v sini for the model parameters.
66 km s−151993) and is a member of
6
on, with a distance ofParameter143
Age (log t)
6.8±0.08yr
7 Value
SymbolLuminosity
9.292.26
3s classified as G2 subgiant
Input parameters−1.65 L&∼9.7 L&2b
sin i = 66 km s−1 and isone (star)
MassM
1.88 M!
!
Inclination
62 +4
8
rs with V=+9.16 magTemperatureand−8 degrees(effective)
effDistance
143Tpc
9 5860 K
acterised by strong photoLuminosity(star)
L!
Av
0.9 mag
3 6.3 L!
rspots and bright accretion
Distance(1) Herbig (1952);(2) DeWarf et al.
D (2003); (3) Bertout
144 parsec
et al. (2007);
Surface (4)densityturnoverR0
500
Muzerolle
et al. point(2003); (5) Johns-Krull(1996);(6)auUnruh et al.
many studies at all wave(2004);(7)
Bertout et al. (2007);(8) Akeson 0.05
et al. µm(2002); (9)Minimum
grain
size
a
min
imarily focused on underBertout& Genova (2006). a The fstellar coordinates120
are taken from
Gas-to-dust
ratio
g/d
s large-scale circumstellar
Simbad (http://simbad.u-strasbg.fr). b The value of the luminosity ofInnerradius
of
the
gas
disk
R
2
R
!
res of a circumstellar disk DeWarf et al. (2003) comprises thegas,in
luminosity of thestar= 6.3±1.6
rate
the cloudof1 the disk Ṁ
10−7 M! /yr
et al. L& andd UV excess (BertoutMass-infall
infalltheofluminosity
∼3.4
L& .

arcseconds

arcseconds]
Polar opening of the cloudµ0
0.8
nd in spectroscopic obserm circumstellar and/or
ac- parameters
FittedFigure 4. Fitted SED of SU AUR using the model paramepolarimetricis achieved by
employingthe
Inner radiusTheof thedust disk sensitivity
Rdust,in
0.17
AU
ters described in Section 5. Data are taken from Cutri & et al.
et al. 1993).
beam-exchange
ervations of SU Aur’sDust
cir- massin the disk method (Semel
Mdust
1.2 This· 10−4method
M! resystematic effects by simultaneously detecting orthogo- (2012) (WISE), Høg et al. (2000) (the Tycho-2 Catalogue),sing long-baseline interferSurface duces
density
power-law
γ
2.4
S. V.
Jeffers
et al.: Imagingthe
SU Aurat the T Tauri/Herbig
transition (1991) (Johnson UBV),
3 Morel & Magnenat (1978)
S. V. Jeffers et al.: Imaging SU Aur at the T Tauri/Herbig transition
7
Mermilliod
linear polarisation states
andexchanging
polarisation
et al. nal grain
erferometer (AkesonMaximum
size
amax
66 µm
most sensitive to the hot- states between the two beams. ExPo uses a polarising cube (Johnson UBVRIJKLM), Ishihara et al. (2010) (AKARI/IRC
Power law of the size distribution p
4.8switch- mid-IR all-sky Survey - L18W+S9W), Ducati (2002)
tar, resulting in a reliable beamsplitter and a ferro-electric liquid crystal (FLC)−3
Turbulence parameter 3 synchronised
α with the 35
9 ·Hz
10 frame rateius (0.13/0.18 au depend- able half-waveplate(Johnson VJKLMN), Ofek (2008) (Johnsonrate
3 · comprise
10−7 M! /yr
6 JKH), ESA
camera. ExPoṀobservations
long seet al. of the EM-CCD
or inclined (AkesonMass-accretion
P=15%
P= 3% Droege et al. (2006) (Cousins
Inferredparameters
quencesofmodulated
short
exposures
for
each
FLC
position, (1997) (Tycho-2.BT+.VT),
spectropolarimetric obserDisk inclination(approximate)i
50 degrees
each lasting
28.82ms . The EM-CCD
technology
enables the I, Johnson V), Monet et al. (2003) (USNO-B1.0.B1+.R1),
al. 2005) who also suggest
4all-sky Survey th
ISAS/JAXA(2010)(AKARI/IRC
mid-IRobservationofcomparatively
fainttargets(down
to13
magniDisk
position
angle
of
short
axis
10-20
degreescally thick scattering disk.
005) determined SU Aur’s tude in the V-band) at these short exposure times without being WIDEL+WIDES+N160+N60), Kharchenko (2001) (Johnson◦
1
1 fromdominated
detectorreadout noise.
2 and Hauck et al.
with the value of ≈60
deAkesonet al. (2005).
The
mass-infall
rate byis taken
BV), Skrutskie et al. (2006) (2MASS JKsH)
For a given position of the FLC polarisation modulator, themographic Doppler imag(1990) (Johnson UBV), Andrews & Williams (2005) (450 andA summary of SU Aur’s instrument will be sensitive to polarisation in two orthogonal
850µm
filters).
◦00 Q),isbutdominatedStokesnot in the 45bydirection(Stokes1.
plies that directions
the size (e.g.
distribution
submicronU). To measure polarisation along those axes, the FLC is rosized particles with almost
no large dust grains and the upper
tated by 22.5◦. For redundancy, data are taken at four FLC an-2
lysis
grain size gles
is poorly
constrained.
The◦◦ disk extends between 0.16
(0◦ , 22.5,-145◦ and 67.5).
6.1. SU Aur’s disk
au and 500 au
the total
is computed
Thewhereobservationsof SUdust
Aur mass
were obtained
withouttothebeuse
nical setup and summarise
1.2 · 10−4 M
.
The
mass
accretion
rate
of
disk
material
onto
theof!an adaptive optics system. Our data analysis method (de- The derived disk parameters were computed-4by fitting the SED−7
scribedmakesof the numerous
shortαexposures-2
star is 3 · 10
M!below)/yr witha a useturbulence
parameter= 9 · 10−3by and comparing the resulting image with the ExPo images. This
including a custom
shift-and-add
combination
approach
and a surface-density
power
law (as a image
functionof radius)
γ = is consequently much more accurate than fitting-6 the SED alone.
(Canovas et al. 2011), which reduces the effect of atmospheric
2.4 (ref Equation 5).
Additionally, because SU Aur is known to be actively accreting-3
eter that was designed and seeing.
Figure 2 we show the resulting model at fixed inclination (Kurosawa et al. 2005), an optically thick gas disk was added
visitor instrument atIntheangles.
to 2 -2
Herschel Telescope.
ExPoFrom these models we conclude that the model with an inside the dust-sublimation radius that extended inwards-6 -4
0 2 4 63.2. Observations-3 -2-1 polarimetric0
1 R .2Excluding
3 this additional luminosity source resulted
inclination
angle of 50◦ reproduces
the observed
wavelength range of
500 toin a arcseconds!
arcseconds
y in linear polarisation
SU dark
Aur wereasthe
partcentre
of a larger
imageofbestThe(refobservations
Figure 1).ofThelane secured
throughof very poor fit to the SED at 1.5 to 6 microns. The resulting fit
6
Figure 3. Full-field ExPo image of the circumstellar environment of SU Aur in polarised intensity, taken without filter. The
with respect to theat observing
campaign
to is
image
the circumstellar
environments
thestarpolarised
intensity
image
a resultof the subtractionof intensity
to the (PSED
corresponds
to intensity
a simulated
imagethat closely
re- of magnitude and north is up and east is to the left. The polarisation vectors plotted
Figure
1. ExPo imageof the circumstellar
environmentof SU Aur in polarised
without a filter. The
intensity
scale is logarithmic
over two orders
I ), taken
denhuis et al. 2012). ExPo of
young
from
December2010 toJanuary
2010.
To represents a lower limit to the true polarisation on the image indicate the fraction of polarised light (PI /I). At the very centre of the image the vectors are unrealistically small
scale
is logarithmic
over 27
twocomponent,
orders of magnitudeand4 thelengthofthe tovectors
the polarisation
ofstars
the central
which
was
made
of the ExPo image.
because the total intensity of the star dominates the polarised light. As explained in the text, our derived polarisation degree is as in circular polarisation.
measuretheiscomplete
signal (Stokes Q and sembles the spatial extent and brightness
degree. North
up and east linear
is to thepolarisation
left.

correct any residual instrumental polarisation. A side effect of The mass accretion rate of 3 · 10−7 Mlower
limit agrees
to the real value.
well with the! /yrthis method is that we also subtracted the polarisation originat- value computed by Calvet et al. (2004) of 5 − 6 · 10−7 M! /yr.
U)foursubexposuresof
4000
frames
each
were
taken
per
FLC
Approximately
1000
dome
flats
were
combinedtoproduceone−4
5.3.
Extendedcircumstellar
ing at the innermost edge of the disk, resulting in the artificial Additionally,
the determined dust disk mass of 1.2 · 10environmentM!
position. The observations of SU Aur are summarised in Table masterflat per FLC position (0◦ , 22.5◦ , 45◦ and 67.5◦ ).
The large-scale ExPo image (15!! × 15!!
et al.
dark lane seen
the imagealsoinMin(assuming a gas-to-dust ratio of 120,
ref. Table 3) is quite
2. Theinpolarisedfeatures(seedescribed
this paperare 2012).
present inSince weboththedata same
sets, though
the data set from
1 January isonused
used exactly
dataonlyreduction
techniquesthe3.3.
obsertypicalandisapproximately
1
percentofthestellarmass
Data reductionin this analysis as the atmospheric seeing was much lower.
vations andthe modelled data, the dark lane is also present in (Scholz et al. 2006; Williams & Cieza 2011).
For ExPo, the FLC modulates the orientation of the lin- The data-reduction software has specifically been developed◦
the model ear
imagesthatstate
wereconvolved
withsimulated
ExPo with the aim
minimisingatmosphericinstru- layers of SU Aur’s disk arepolarisationof theincoming lightby 90theevery
0.028s forExPo
Theof grain
sizes
in the andsurface
The aim is to separate the linearly polarised
switching
2). between successive A and B frames, resulting in a to- mental effects.mostly
PSF (Figurein the formof smallgrainsof the unpolarisedstarlight.
The first
step is that follow a size distribution
tal of 2000 A + 2000 B frames = 4000 frames. For each frame signal from thatof individual
frames,
then
the
four
subframes(A or B) the light is separated using a beamsplitter into two the processingwith
a power-law exponent of 4.8. This implies that most of theorthogonal beams (left and right beams), which are simultane- (Aleft , Aright , Bleft , Bright ) are combined in sequence to form a pomass is in the form of small grains, and because there are very
ously imaged on the EM-CCD, a concept that is the basis of larised intensity image. This image is then combined at four
6. Discussion
few
large datagrains,maximum
calibration
to formtheira calibrated
polarised size cannot be clearly cona dual-beam polarimeter. For every pair (A and B frames) of FLC angles withobservations, this results in four different images of SU Aur intensity image. The data analysis is described in more detail
We have forthe first time directly imaged the circumstellar en- strained. The grain sizes were obtained from fitting the domi(with a total of 8000 for each FLC angle as described above), by Canovas et al. (2011) and is briefly summarised here.
vironmentfor
ofinstancethe classical
T
Tauri
star
SU
Aur.
Using
interprenantsilicatefeatureat∼10µm in SU Aur’s Spitzer spectrum(A+B): Aleft , Aright , Bleft , Bright .
exposure timeeach subexposure
is 4000*
tive models, The
we totaldeterminedtheof inner
radius ofthe dust
disk,
simultaneouslywith SU Aur’s SED (see Figure 4). Supporting
3.3.1.
Individual
frame processing
0.028 s (total=112 s). At the end of every FLC sequence a sethe mass ofriesthe
dustframes
in thethetarget,
surface-density
power
smallforgrainsin the outer layers of SU Aur is
Eachlaw,
image isthe
firstpresence
individually ofcorrected
dark, flatfieldof darkare disk,taken percomprising approxiidentified and removed
8000
the beginningand end of everythenightturbulencea and bias effects.
the power mately
law oftheframes.
dustAtgrain
size distribution,
theCosmic
workraysof areKessler-Silacci
et al.by(2006), who analysed Spitzer
set of dome flat-fields were taken with a polarisation calibrator. comparing the two simultaneous left and right images of eachparameter, and the mass-accretion rate.
spectra of disks around T Tauri stars. Kessler-Silacci et al.

Classical T Tauri stars: magnetic fields, coronae, and star-disc interactions(a) AA Tau 20093

16

Johnstone et al.

Figure 11. Potential field extrapolations for AA Tau (upper panels), the TW Hya 2010 field (middle panels), and V2247 Oph(lower panels) showing closed and open field lines (left-hand panels) and accreting field lines (right-hand panels). Open field lines
are deep blue, closed field lines are red, and field lines that are able to support accretion are light blue. These three stars have beenchosen to represent the three kinds of accretion geometries that are typical in the sample of CTTSs. The dominantly dipolar field of
AA Tau leads to accretion footpoints at high latitudes, while the material accreting onto TW Hya is deflected towards mid-latitudes due
to its strong octupole field component that is approximately anti-parallel to its weak dipole component. The complex field of V2247 Ophleads to a complex pattern of mid-to-low latitude accretion.
c 0000 RAS, MNRAS 000, 000–000

A. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with Herschel

4.Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observedwith Herschel
A. Karska1,2, F. Herpin3,4, S. Bruderer1, J.R. Goicoechea5, G.J. Herczeg6, E.F. van
Dishoeck1,2, I. San Jos ́e-Garc ́ıa2, A. Contursi1, H. Feuchtgruber1, D. Fedele1, A.
Baudry3,4, J. Braine3,4, L. Chavarr ́ıa5, J. Cernicharo5, F.F.S. van der Tak7,8 and F.
Wyrowski9
ハーシェル遠赤外のatom/molecular linesを用いて、High mass YSO 10個について調べた
結果を、low- とintermediate の結果と比較した。
lineとしては、CO, H2O, OH, [O I]。周波数範囲は、55-190μm(Herschel/PACS)。
結果として、・High mass YSOsのFar-IRのcoolingはCOが担っている(15-85%でmedianは75%)。[O I]
はちょっと(20% median)。H2OとOHは1%。low massでは、CO,H2O,OH,[O I]の寄与はほぼ同じ。H2Oのcoolingの効果が下がっているのは吸収による。
A. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Regions observed with HerschelFig. 4. Relative contributions of [O i] (dark blue), CO (yellow),
H2 O (orange), and OH (red) cooling to the total far-IR gas cooling at central position are shown from left to right horizontallyfor each source. The objects follow the evolutionary sequencewith the most evolved sources on top.

van der Tak et al. 2013). The only H2 O lines observed in common by the two instruments within the WISH program are theground-state transitions: 212 − 101 at 179.5 µm (1670 GHz) and
221 −112 at 180.5 µm (1661 GHz), both dominated by absorptions
and therefore not optimal to estimate to what extent a complex
water profile is diluted at the PACS spectral resolution. However,
HIFI observations of the lines between excited rotational states,
which dominate our detected PACS lines, are generally in emission at the longer wavelengths probed by HIFI. In the case of
12
CO 10–9, the line profiles of YSOs observed with HIFI consist of a broad outflow and a narrow quiescent component withthe relative fraction of the integrated intensity of the narrow to
broad components being typically 30-70% for low-mass sources(Yıldız et al. 2013). For the single case of a high-mass YSO, W3
IRS5, this fraction is about 50% (San José-Garcı́a et al. 2013).
To summarize, PACS spectra of high-mass sources from oursample show detections of many molecular lines up to high excitation energies. CO, H2 O, OH, and CH lines are seen towards
all objects, whereas weaker lines of other molecules are detected
towards less than half of the sources. CO lines are always seenin emission, CH in absorption, whereas other species show different profiles depending on the transition and the object. Table
C.1 shows the CO line fluxes for all lines in the PACS range.

4. Analysis
4.1. Far-IR line cooling

Emission lines observed in the PACS wavelength range are used
to calculate the contribution of different species to the total line
cooling from high-mass protostars. Our goal is to compare thecooling of warm gas by molecules and atoms with the coolingby dust and connect them with the evolutionary stages of theobjects. Relative contributions to the cooling between different
molecules are also determined, which can be an indicator of thephysical processes in the environments of young protostars (e.g.
Nisini et al. 2002; Karska et al. 2013).
We define the total far-IR line cooling (LFIRL ) as the sumof all emission line luminosities from the fine-structure [O i]lines (at 63 and 145 µm) and the detected molecules, follow6

ing Nisini et al. (2002) and Karska et al. (2013). [C ii], the mostimportant line coolant of diffuse interstellar gas, is also expected
to be a significant cooling agent in high-mass star forming regions and extragalactic sources. It is not explicitly included in
our analysis, however, because the calculated fluxes are strongly
affected by off-source emission and often saturated (see also
Section 2). The emitted [C ii] luminosity is shown below for only
two sources, where reliable fluxes were obtained. Cooling inother ionic lines such as [O iii], [N ii], and [N iii] is also excluded,
due to the off-position contamination and the fact that thoselines trace a different physical component than the moleculesand [O i]. Since the only molecules with emission lines are CO,
H2 O, and OH, the equation for the total far-IR line cooling canbe written as: LFIRL = LOI + LCO + LH2 O + LOH .
Table 2 summarizes our measurements. The amount of cooling by dust is described by the bolometric luminosity and equals∼ 104 -105 L# for our sources (van der Tak et al. 2013). The total far-IR line cooling ranges from ∼ 1 to ∼ 40 L# , several orders of magnitude less than the dust cooling. Relative contributions of oxygen atoms and molecules to the gas cooling are
illustrated in Figure 4. Atomic cooling is the largest for the more
evolved sources in our sample (see Table 1), in particular for
NGC7538 IRS1 and AFGL2591, where it is the dominant line
cooling channel. For those two sources, additional cooling by[C ii] is determined and amounts to ∼10-25% of LFIRL (Table 2).
Typically, atomic cooling accounts for ∼20 % of the total line
far-IR line cooling. Molecular line cooling is dominated by CO,
which is responsible for ∼ 15 up to 85% of LFIRL , with a median
contribution of 74%. H2 O and OH median contributions to thefar-IR cooling are less than 1%, because many of their transitions
are detected in absorption. Assuming that the absorptions arisein the same gas, as found for the case of Orion (Cernicharo et al.
2006), they therefore do not contribute to the cooling, but heating of the gas. Still, the contribution of H2 O to the total FIR
cooling increases slightly for more evolved sources, from ∼5%(DR21(OH)) to 30% (W3IRS5), whereas no such trend is seen
for OH.
4.2. Molecular excitation

Detections of multiple rotational transitions of CO, H2 O andOH allow us to determine the rotational temperatures of theemitting or absorbing gas using Boltzmann diagrams (e.g.
Goldsmith & Langer 1999). For H2 O and OH, the densities arelikely not high enough to approach a Boltzmann distribution andtherefore the diagrams presented below are less meaningful.
Emission line fluxes are used to calculate the number of
emitting molecules, Nu , for each molecular transition using
Equation (1), assuming that the lines are optically thin. Here,Fλ denotes the flux of the line at wavelength λ, d is the distanceto the source, A is the Einstein coefficient, c is the speed of lightand h is Planck’s constant:
4πd 2 Fλ λ
(1)
hcA
The base 10 logarithm of Nu over degeneracy of the upper levelgu is shown as a function of the upper level energy, Eu , inBoltzmann diagrams (Figures 5 and 6). The rotational temperature is calculated in a standard way, from the slope b of thelinear fit to the data in the natural logarithm units, T rot = −1/b.
Because the size of the emitting region is not resolved by our
instrument, the calculation of column densities requires additional assumptions and therefore only the total numbers of emitting molecules is determined. The formula for the total numberNu =・1-10^6 Lsunの広いluminosity rangeに渡って、far-IRのcoolingはborometric luminosityに強い相関がある、low massでの傾向と一致し。molecular/ atomicのcooling ratioは 4程度
でlow massと同じ。。
・low massと比較しhigh massでは、H2Oが吸収線となっていることから、gasとdustのcooling ratioは1.3 x 10^-3から6.2x10^-5ほどに減ずる。

lecular lines from Low- to High-Mass Star Forming Regions observed with Herscheldied with PACS in12); Wampfler et al.
-mass objects from
This allows us to∼ 1 to 106 L" , andis 200 pc, whereasrison, the full PACS
50## ), which correor of 3 smaller thanh-mass objects. On
w-mass sources arefactor that is come of low- and highhould compare justy the central spaxel
sults, however (seensitions is based onh samples and theirelations are quantinumber of sources
to r ≈ 0.6 and 5 σ

ns between the seinosities as well as
rce, the larger is its
arly, the more masprotostar, the larger
ies. The strength ofanges and envelope
responsible for thebjects, Karska et al.
en with PACS withw walls, most likely
ssion, on the otherative shocks at thense envelope. In the
d the strength of rahe emission can bein lines, its molecing by dust for the6
L" .
strongly (5σ) withnt with studies on
et al. 2013). Under
he shock energy, theas been interpreted
wer/ velocity being
protostellar surface
ction rate.
oling, Lmol /Latom , isthough a large scat4 times higher than
was included in thecrease for the highemission accounts(Goicoechea et al.
rces, this contribu-

5.Radial velocity variations in the young eruptive star EX Lup
A ́ . K ́o s p ́a l 1 , M . M o h l e r - F i s c h e r 2 , A . S i c i l i a - A g u i l a r 3 , P . A ́ b r a h
́a m 4 , M . C u r ́e 5 , T h . H e n n i n g 2 , C s . K i s s 4 , R. Launhardt2, A. Mo ́or4, andA. Mu ̈ller6
EX lupのoutburstによる変動のを5年のRVのsurvey data(radial velocity)で調べた。
測光(optical/infrared)観測では14-dayの活動周期が知られている。
$

受信機FEROS(fiber-fed Extended range Optical spectrograph)$$$and HARPS(High Accuracy Radial velocity Planet Searcher)
望遠鏡MPG/ESO 2.2m, 3.6m telescope, La silla observatory,Chile$

波長域：3500-9200Å, 3780-6910Å$

line；Ca 8662Å, Ca II K

結果として；
RVは周期7.417d、振幅A=2.2km/s。
RVの挙動の解釈として、$

1.２つの幾何学的なacrretion columnsが存在し、片方は7dayである。

$

2.companion(brown dwarf)の存在。Keplerian fitしてみると、中心星0.6 Msunの周$Kóspál et al.: RV variations in EX Lup
りにM2=(14.7±0.7)Mjup、e=0.24がえられる。

GLS Power

1.0
0.8
0.6
0.4
0.2
0.0

6.Interstellar chemistry of nitrogen hydrides in dark clouds
R. Le Gal1, P. Hily-Blant1,2, A. Faure1, G. Pineau des Forˆets3,4, C. Rist1 and S. Maret1
窒素N系の気相での化学反応計算とIRAS 16293-2422の観測結果との比較。
先行研究に対して新しい点は、$

1.実験室系の結果,特にN->N2変換のradical-radical反応を含んでいる。

$

2.H+,H3+でorth/para変換の理論的結果を取り入れた$

3.核spin ratioの関係式を取り入れた

彼らの計算結果ではC/Oの比と硫黄Sの気相の存在量が重要らしい。
観測結果(NH:NH2:NH3=3:1:19)を再現するbest resultとしては硫黄の存在量として[S]total=8.0x10^-8でC/O=0.8。ただしNH3はこの計算でも観測より低い(範囲ギリギリ?)。
R. Le Gal et al.: Interstellar chemistry of nitrogen hydrides in dark cloudsmical Friction in a Gas: The Supersonic Casemass traversing
a Steven
relatively
Aaron T. Lee1 andW. Stahler1 sparse gas experiences a retarding force
dium.
In a previous contribution (Lee & Stahler 2011), we determi
partment, University of California, Berkeley, CA 94720
7.Dynamical
Frictionin a Gas:We
The SupersonicCase’s
now extend
our analysis to the supersonober velocity
2013 / accepted: 8 was
Novembersubsonic.
2013
Aaron T. Lee1 and Steven W. Stahler1
ions created in ABSTRACT
the gas far from the gravitating object, and thereb
ガスの力学摩擦をsupersonicの場合に解析解を求めた。
a traversing
large,
bounding
surface.
mass
a relatively
sparse gas experiences
a retarding Various
force created by itsterms
disturbance ofin
the the perturbation series
um. In a previous
contribution (Lee & Stahler 2011), we determined this dynamical friction force
結果として、
atmentoftheflow
streamlines.
Nevertheless,
we are able to derive
velocity was subsonic.
We
now extend
our analysis
to the supersonic regime.
As before, we consider
ns created in the gas far from the gravitating object, and thereby obtain the net influx of linear
large,
bounding
surface. F
Variousin theseries formally
ement
find
that
=termsṀ
Vperturbation, whereṀ isdiverge,thenecessitating
rate anof mass accretion onto t$ streamlines.
of the flow
Nevertheless, we are able to derive exactly the force itself. As in thefind that F =, where Ṁ is the backgroundrate of mass accretion onto gas.
the object Our
and V its force
instantaneouslaw holds even when the
spect
toṀ Vdistant
ect to distant background gas. Our force law holds even when the object is porous (e.g., a galaxy)elling mass in a wind.
Quantitatively,
force in the supersonicregime is less than that derived
xpelling
mass
in athewind.
Quantitatively,
the force in the supersonicevious researchers, and is also less than was found in numerical simulations through the mid 1990s.
s to revisit the problem using modern numerical techniques. Assuming our result to be correct, it isprevious
researchers, and is also less than was found in numerical sim
y fields of astrophysics, ranging from exoplanet studies to galactic dynamics.
dynamicsto— waves— ISM: generaldisksusing— galaxies: kinematics
and dynamics
ors
revisitthe— protoplanetary
problemmodernnumerical techniques. Assumany fields of astrophysics, ranging from exoplanet studies to galactic
fluid elements follow hyperbolic orbits in the star’s reference frame and land in an infinitely thin, dense spindle
behind the object. Bondi & Hoyle analyzed the transfer oflinear momentum from the spindle to the star, and hence
obtained the force. Dodd & McCrae (1952) similarly investigated this hypersonic limit, as did, much more recently,
Cantó et al. (2011). Dokuchaev (1964) first treated a finitetemperature gas. He determined the force by integrating
the total power emitted by the object in acoustic waves(see also Rephaeli & Salpeter 1980). Ruderman & Spiegel(1971) used an impulse approximation in the reference
frame of the background gas. Finally, Ostriker (1999) calculated the force by integrating directly over the wake, whose
density she obtained through a linear perturbation analysis.

object passes through—a rarefied
rodynamics
waves — ISM: general — protoplanetary disks — gal
ounding matter toward it. As a result,an overdense wake behind the object
avitational pull, retarding the originalal friction arises whether the medium
cting point particles, e.g., a stellar
um fluid, e.g., an interstellar cloud.
) provided the essential theory when
lisionless, and his solution has beenudies of both star clusters and galaxavitating object with gas also occurs
tuations. A partial list of topics andhe interaction of planets and gaseouserquem, & Papaloizou 2013); the orn-envelope binaries (Ricker & Taamf massive stars in dense molecularal. 2010); the coalescence of masth isolated galactic nuclei (Narayan
laxies (Armitage & Natarajan 2005);
ntracluster gas by infalling galaxies
mionkowski 2004).
pread occurrence of gaseous dynamitill no generally accepted derivationr 70 years of effort. The flow in theating mass is complex both tempomany simulations have shown (see,& Sawada 1987). Theorists seeking a
on for the time-averaged force havetegems to circumvent a detailed den. In the course of their classic stud, Bondi & Hoyle (1944) took the starro-temperature gas. In their model,

fluid elements follow hence frame and landve object passes through a rarefied behind the object. Bon
These researchers
on the supersonicrrounding matter toward
it. focused
As aprincipally
result,linear momentum froms an overdense wake behind the object obtained the force. Do
ravitational pull, retarding the original tigated this hypersoni
ical friction arises whether the medium Cantó et al. (2011). Do
racting point particles, e.g., a stellar temperature gas. He
uum fluid, e.g., an interstellar cloud. the total power emitt3) provided the essential theory when (see also Rephaeli & S
ollisionless, and his solution has been (1971) used an impu
studies of both star clusters and galax- frame of the backgroun
gravitating object with gas also occurs lated the force by integ
situations. A partial list of topics and density she obtained th
the interaction of planets and gaseous
These researchers f
Terquem, & Papaloizou 2013); the ormon-envelope binaries (Ricker & Taam case, which is often theof massive stars in dense molecular That is, they took Vt al. 2010); the coalescence of mas- gravitating mass relati
both isolated galactic nuclei (Narayan the sound speed in tha

8.Sejong Open Cluster Survey (SOS). II. IC 1848 Cluster in the H II Region W5 West
Beomdu Lim1, Hwankyung Sung1, Jinyoung S. Kim2, Michael S. Bessell3 and Rivkat
Karimov4
IC1848：giant star formation のCas OB6 associationの若いopen cluster
SOS: Sejong Open cluster Surveyの一部
ウズベキスタンのMaidanak astronomical observatoryのAZT-22 1.5m望遠鏡を使用
1” resolution、 UBVIとHαフィルター＝＞(U-B,B-V)から、redding <E(B-V)>=0.660±0.054
公開されている近中間赤外のアーカイブデータを用いて、Rv=3.1を確認$

Spiter IRAC data；$ 397 young stars$

Hα測光；$$$$

total；462 PMSs$

257 PMSs

clusterの距離を決める為にZAMS fittingをして、距離の補正として11.7±0.2mag。
median ageは3Myr。
IMFのベキはΓ=-1.6±0.2 (log mass>0.5)で、Scalo IMFに近い(Salpeter/Kroupa IMFより)。
同様の質量帯のopen clusterのNGC 2264と同じ傾向。
The young open cluster IC 18483

我々の測光から、UV超過を伴う51 PMSsを見つけた。
Table 1. Photometric data and spectral types for 12 bright stars
ID

BD +59 556
BD +60 584
BD +60 586
HD 17505
HD 17520
HD 17688
HD 237007
HD 237011
HD 237012
HD 237015
HD 237018
HD 237019

V

10.06
9.87
8.48
7.06
8.27
10.64
9.44
10.05
9.67
9.44
9.01
9.73

B

V

0.36
0.28
0.30
0.40
0.32
0.56
0.33
0.30
0.47
0.24
0.14
0.47

U

B

-0.12

reference

Spectral typereference

Hoag et al. (1961)
Høg et al. (2000)
Johnson & Hiltner (1956)
Johnson & Hiltner (1956)
Johnson & Hiltner (1956)
Bigay (1963)
Johnson & Hiltner (1956)
Høg et al. (2000)
Høg et al. (2000)
Haug (1970)
Høg et al. (2000)
Johnson & Hiltner (1956)

B9 V

Hoag & Applequist (1965)

O7.5 V
O6.5 III((f))
O9 V
A7 V
B0 V
B1.5 V

Hillwig et al. (2006)
Hillwig et al. (2006)
Conti & Leep (1974)
Fehrenbach (1966)
Morgan et al. (1955)
Reed (2003)

B4 D
B9.5 V
O8 V

Reed (2003)
Fehrenbach (1966)
Morgan et al. (1955)

U luminosityからaccretion luminosityに変換した。
-0.67
-0.64
-0.68
-0.08
-0.30

accretion massはaccretion luminosityとモデルから求めた星半径,質量を用いて計算した。
＝＞星の質量rangeは0.5 Msunから2Msunで、1.4x10^-8Msun yr^-1。この結果は他の領域-0.43での結果と一致する。high massだと予想されるより高いaccretion ratioの傾向があった。
-0.53

The young open cluster IC 1848

15

Figure 13. Accretion luminosity vs. stellar luminosity (left) and mass accretion rate vs. stellar mass (right). Open circles (red) denote
our estimates. Squares, triangles, and crosses represent the mass accretion rates derived in di↵erent star forming regions by Hartmann
et al. (1998), Mendigutı́a et al. (2011), and Natta et al. (2006), respectively. Dotted lines represent the relations Lacc / Lstellar and2Ṁ / Mstellar
with an arbitrary constant.

3.1. Saturated regions and measurement of the spectrum of the Galactic foreground

After a global investigation of the IRDC, we focus on three 1! ×1! regions around BT12 and BTK core(Fig. 2), C4 (Fig. 3) and C11 (Fig. 4) to search for local
saturation
at 8, 24, and 70µm. As seen in the
TABLE
1
a (cm
−1 )
these figures, the Telescope
8µm saturation
is morewidespread than seen
BT12, Peri.e. Gas
in C1
andC4,2 gwhich
is aBand andBackground-Weighted
Dust byOpacities
Massof searching in local regions and thus reducing the masking effect of foreground variations. There is gebcorrespondenceof saturated
regions
acrosstheIRAC4.5
different wavebands,
althoughcan bePACS70
disruptedPACS10
by th
Dust
Model
IRAC3.5
IRAC6
IRAC8 thisMIPS24
9.Far-infrared extinction
mappingof infrared
dark
clouds3.52 µm
4.49 µmto be5.91µmextended
7.80 µm
23.024µm,µm 74.0µmat 70µm
103.6 µm
discrete sources. There is a tendency
forcsaturation
more
at 8 andthan(
Wanggi Lim1 andJonathanC. Tan1,2!!
in C11).
The
region in 5.45
C1 at 70µm3.63
is not greater
is larger 0.185
if theWD01
RV =3.1 size of the 2σ saturated
8.73
5.45 than 12
3.92 but it0.392
relaxed4σ, so we consider this likely
saturated 4.86
core. In general,the2, 4, 8σ-defin
WD01
RV to
= 5.5
11.5 to be a7.59
6.13 comparing
4.10
0.420
0.193
OH94
thin wemantle,0 yr
19.3morphologies,(14.0) 11.8 (9.57)(7.77) that
6.89 (6.77)
4.88are revealing
0.762
0.389
regions,
see coherent,
contiguous
which8.56
indicates
these levels
real
clou
Spitzer MIPSとHerschel
PACSのイメージを用いて、−3
OH94
thin mantle,
105 yr,106 cm24.7reasonably(17.9) 14.6well-estimated.
(11.9) 10.4 (9.43) 8.24 (8.09)
6.86
1.14
0.603
and thatthe imagenoiselevelsare
7
6−3
OH94 thin mantle, 10 yr, 10 cm
26.5 (19.2) 16.2 (13.2) 11.7 (10.6) 9.18 (9.01)
8.33
We derive the specific intensitiesof(26.1)
the Galactic
foreground towards C1, C4, C11,
also 1.42
including0.746
mea
IRDC G028.37+00.07にある三つのmassive
sterlesscore/clumpsに対して、
OH94 thick mantle, 0 yr
36.0
15.1
(12.3) 16.7 (15.1) 9.34 (9.17)
10.6
3.13
0.950
IRAC
bands
1
to
3
at
the
locationsof
8µm
pixelsshowing2σ-saturation.
We
plot
these
intensities,
56−3
OH94 thick mantle, 10 yr, 10 cm
43.5 (31.5) 17.8 (14.5) 18.7 (17.0) 10.7 (10.5)
13.2
3.96
1.27n
−3 In45.7
面密度Σ mapを求める(独自の?)解析を行った。
the 8µm
on7 yr,Figure1b.
general,agree10.9
wellwith theLD01 model.
The1.44
24µ
OH94
thick values,
mantle, 10106 cm(33.1) these
18.2 measurements(14.8) 19.4 (17.6)(10.7)
14.8
4.55
consistently about a factor of two smaller, while the 70µm values are a few tens of percent larger. The#やっていることは、光学的厚みとopacityを考慮したみたい。
a ratios
ourchoiceof (Draine
normalizing
at 8µm. In the IRAC bands 1 and 2, the observed inCommonare
total sensitive
to dust mass toratio
of 142is adopted2011).
b References: WD01 - Weingartner & Draine (2001); OH94 - Ossenkopf & Henning (1994), opacities have been scaled from values in parentheses to insignificantly lower than the LD01 model of total emission. However, there is reasonable agreement withfrom scattering.
＝＞thick ice mantle
model(Ossenkopf
and (dotted
Henning1994)と一致する結果が得られた。
of Galactic emissionline).
This IRDC is at 5 kpc distance, and this appears to be
c component
Mean wavelengths weighted by filter response and background spectrum.
thatforeground stars are not making significant contributions to the foreground emission as measured on
この結果は、低密度から高密度に渡ってicescales opticalin thesedepth
saturartedwhereτν = cores.
κν Σ andmantleが成長していることを意味する。
κν is total opacity at frequency ν per unit total mass and Σ is th
surfacedensity.
The
value
of
I
is
to
be estimated via a suitable interpolation from the region surroundinν,02
Lim and Tan
FIR Extinction
Maps towards the cloud. However,while Iν,1 is to be derived from the observed 3.2.
intensity
to given locations
Following
the method
described
in §2,Igiven
the foregroundobserved
Galactic
backgroundemission,and the ofobserved
intensityand
justbackgroundin front of intensities
the IRDC,
ν,0,obs ,estimatesregion we derive
extinctionemission,
maps at I8,
24 ,and70µm,
g cm−2under
the assumptionocontaminatedwiththeforeground(IRDCsaredisplaying
typically Σ
at inseveralkpc
distancein the Galact
ν,foreactually
the thickobserve
ice mantle model of OH94 (middle rows of Figs. 2, 3, 4). The choice of this thick, rather than thi
Iν,1,obs= Iν,fore and+ Iν,1dust model is motivated by the observed opacity
law (below),
is different from BT12’s use of the thmodel. For 8µm-derived maps the resulting variation of opacity per unit mass and thus Σ is quite small
and
it can make more than a factor of 3 difference at 70µm.
Iν,0,obs= Iν,fore+ Iν,0 . begins to start underestimating th
The value of Σ at which our saturation-based
extinction
mappingsurfacedensityis(BT12)
Therefore, like MIREX mapping, FIREX mapping requires measurement of Iν,fore towards each region to)τν,sat forln(IFollowing BT12, we estimate Iν,fore empirically by looking
spatially
dark regions that are
ν,0 /Iν,1independent, set by the noise level of the im
sat = some=intensity range
i.e. they have the same observed intensity toΣwithinκκνν
8 µm GLIMPSE images with 2!! resolution and 1σ noise level of 0.6 MJy/sr, BT12 defined saturated pix
where
I
is
set
equal
to
the
2σ
noise
level,
i.e.
1.2,
0.5,
30
MJy/srfor the8, 24,However,
70µm images.
these twithin a ν,1
2σ range above the observed global minimum value in a given
IRDC.
for theAtIRDC
Fig. 1.— (a) Left: Three-color image of IRDC
G028.37+00.07(red: 70µm,
green: 24µm
& blue: 8µm).
Small circles
indicate massive!!
thesaturated
cores
have
averagevaluesofI=92.7,
53.7,
1302.3
MJy/sr,respectively.
Thus
for
thic
ν,0
exhibitsaturation,thesepixelsto be
distributed over a region at least 8 in extent.
cores studied by BT12 and BTK. Squares are
regions inspected
for local saturation aroundcores C1,
C4 & C11.needed
Dashed ellipse shows
the
IRDC boundary from S06. (b) Right top: Total IR to sub-mm SED (solid line) of the Galactic plane in the “MIRS” region−2
at l ! 44
opacities,= 0.44,
0.39, 1.05
cmthe
,images
respectively,
while bands
for thinmantleopacities
it is
0.75,
We followΣasat
similarmethodheregforin the IRAC
at ice24 and70µm,
but with
the0.58,
followin!
Note,
discussedby BT12,higher valuesof Σ1thanwillofbeview
present
the map,
but cores
these previously
will be ten
(1)
We as
searchfor “local
saturation”
in smaller
by 1Σ! sat
fieldsthatincontain
dense
and
more
affectedbysaturation,
leading
tounderestimationof thetrue Σ (2)values.
BT12. This helps to minimize the effects of foreground spatial
variation.
In addition to a standard
Theweabove
thatranges.
we should
to probe
to higher
of of
Σ in
70µm
maps
t
range
also analysis
consider predicts
4σ and 8σ
Thisbeis able
becausewe regard
the values
estimatethethe
noise
level
in th
and 24µmuncertain.
maps, andWe
that
such
hightheΣ liklihood
structuresofwill
be in regions
have been
identified
as being
somewhatwill
gauge
saturation
also bythat
consideringthemorphology(e.g. satu
cowavelengths.
In core/clump
we indeed
see this morphology:
thewavelengths.
70µm map peaks
at arou
oftwo
theshorter
saturated
pixels, together
with theirC1,
overlapwith saturated
pixels at otherIn generalw
at a position
close
to theC1at
core
center identified
by BT12
BTK.(i.e.
Thisgenerally
highest column
densityregionw
saturatedregionto belargerwavelengths
with larger
dust and
opacities
increasing
at shorter+
withthesize
C1-N
N2 Daffected(3 − 2)bycore
by Tanet al.
We alsoinfer thatThethe ability
70µm map
is r
but
the
is alsotheidentified
relative noiselevelsin (2013b).
the different
wavebands.
to detecstructures
the range Σby
∼ the
0.5 −
1 g cm−2
aresource
not sothat
reliablyprobedthe shortercanalso be incompromised
presenceof, awhich
nearbyenhancesthebylocal
level of wavelength
foreground me
In theregionwe sizes
see twomain
Σ peaks
thataexhibit
local saturation,
withcriterion,
the western
structuGiven
the C4largerbeam
at 24
andhigh
70µm,
we adopt
less stringest
spatial extent
requirings!!
saturatedby about
two times
theby
beam
FWHM,
i.e. 12. two sources seen most clearly at 24µm. T
as C4 by pixels
BT12 andthe eastern
as C13BTK.
In betweenare
Our method affectfor estimating
the backgroundsame
adoptedandthatBT12,
i.e.istheundoubtedly
the MIREX
and FIREX intensity
methods isinthe
thisregion,so itbyis BT09
possiblethere
a hs
filtermethodoutside the defined IRDC ellipse from S06 (filter size set to 1/3 of the major axis, i.e.
betweenthe applied
two cores.
by interpolation
this ellipse.
We inspectedbackgroundthree filamentary
control fieldsstructu
equal
In C11, thereinside
is a relatively
widespreadhigh theΣ region,
whichfluctuations
is part of ainhighly
size
just outside
after The
foregroundsubtraction
the average
1σ fluctuations(froma fittedtoGa
towards
C12 to thethe ellipse:
NE (BTK).
70µm map
also reveals
a relatively
high Σ region
extendingth
17.3%,
25.3%
and in42.3%
8, 24,
70µm,
respectively.
In theoptically
thinthat
limit,
these values
provideanextensionis seenthe 8atand
24µm
maps,
but at lower
Σ. It
is possible
localized
foregroundvaritheuncertainty
in thearederivedτ andΣ inany maps
given (seepixelalso
dueBTK,to backgroundshorter
wavelengths
affectingthethus
8 and
24µm
where thisfluctuations.
extension is more pro
8µm map derived with a finer decomposition of foreground variations).
2.3. MIR to FIR Opacities
3.3. The MIR to FIR Extinction Law and Evidence for Grain Growth
Following BT09, we adopt a spectrum of the diffuse Galactic background from the model of LD01. We w
relative valuesof Σ
derived inmapping
non-saturated
regions at different
wavelengths
yield
information
abo
thatThea by-productof ourextinction
is a measurementof this spectrumof theforegroundemissiofWethethendustconsider
extinction
law.different
This lawis models
expected(Table
to varygrains undergo
via icecoagulationanseveraldust1), asespeciallythe thin growth
and thickmantle moddepositionof volatiles,
as water,CO,to form ice mantles (see Fig. 1d). The Σ maps
6 cm−3 ) grains (OH94). For all models we adopt
erately
coagulated(i.e. suchcoagulation
formethanol
105 yr atand
nH=10
and dust)
4 are derived
moderately-coagulated
thick ice
mantleof OH94,
deviations
inplus
mass to assuming
refractorythe
dust
mass ratio of 142 (Draine
2011,
c.f. model
the valueof 156sousedin BT09
e.g. choice
Σ70µm /Σtell us about
deviationsof the actual
dust
laweffectivefrom theOH94
This
will8µm
not, from
affectunity
measurementof relative
opacities.
Finally,
weextinction
obtain theopacities
In the bottom
rows of Figs.
3 and 4 we
present filter
maps responseof Σ24µm/Σ8µm ,and
Σ70opacityand Σµm /Σ8µm
7wavebands,
i.e. convolutionof the2,background
spectrum,functionfunction
(se
Some
of
the
variation
in
these
ratio
maps
is
due
to
the
different
saturationlevels,
i.e.
the
70µm
ma
and Table 1).
λ Iλ/(8µm I8µm)

10.0

Dust + Stars
Dust

1.0

C1
C4
C11

0.1

1.0

Spitzer IRAC
Spitzer MIPS
Herschel PACS

Filter response

0.8
0.6
0.4
0.2
0.0

10.0

0.1 g cm-2
0.2 g cm-2
0.3 g cm-2

H2O

1.0

1

H 2O

κλ/κ8µm

CH3OH

H2O

H2O
CH3OH

21

25

H2O

1.0

0.1

1

0 yr
105 yr107 yr

0.1

70

10

λ (µm)

78100◦

10.Warm gas towards young stellar objects in Corona Australis - Herschel/PACS observations from the DIGIT key programme
Johan E. Lindberg1,2, Jes K. Jørgensen2,1, Joel D. Green3, Gregory J. Herczeg4,5,
Odysseas Dionatos1,2,6, Neal J. Evans II3, Agata Karska5 and Susanne F. Wampfler1,2
外部からの放射の影響がどうなっているか。中間質量星から低質量原始星コアへの影響。
R CrA dark cloudで２つの星形成領域のHerschel/PACS data。
CO,OH,H2O,[C II],[O I]と連続波を調べた。
点源と広がったソースを切り分けるdeconvolution methodを開発した。
COのrotational diagramも計算した。
deconvolutionの結果として、数1000AUに広がった連続波と、点源に付随しないline
emissionを見つけた。
line emissionから温度として点源でも広がった成分でも、CO(282±4K)、hotCO(890±84K)、OH(79±4K)、H2O(197±7K)。
近くの明るい天体からの影響でrotational temperatureが上がってはいない。けれども、
H2O,OH,Oのphotodissociationした生成物の放射はwarmガスの周りでenhanceされている。warm gasは外部からの放射がないところと比較し、原始星やその周りに付随している。
Johan E. Lindberg et al.: Warm gas towards young stellar objects in Corona Australis

Johan E. Lindberg et al.: Warm gas towards young stellar objects in Corona Australis

A&A proofs: manuscript no. herschel_rcra_arxiv

Fig. 3. As Fig. 2, but for IRS5. Contour levels are at 3 Jy for the continuum maps, and as in Fig. 2 for the line maps. The blue colour maps havethe same ranges for each separate molecule, with maximum fluxes of 1.5 ⇥ 10 16 W m 2 for CO, 2 ⇥ 10 16 W m 2 for OH, 1 ⇥ 10 16 W m 2 foro-H2 O, and 2 ⇥ 10 15 W m 2 for [O i].
Fig. 2. IRS7 contour maps of continuum emission at three wavelengths (green with black contours), and some important spectral lines (the sameas in Table 4 of Green et al. 2013, blue with red contours). The first column in each panel shows the total emission, whereas the second column
in each panel shows the residual emission after deconvolution (emission not associated with the point sources). Such maps of the whole band ofcontinuum and all spectral lines can be found in Appendices A–B. The point sources used for the deconvolution are marked with crosses, except
for R CrA, which is marked with a star symbol (the sources are identified in the top row). The dots indicate the PACS spaxel centres. The Herschel
PSF for each observation is shown in the top left corner. Contour levels are at 10 Jy for the continuum maps, and at 5 , 10 , 15 , 30 , 60 , and90 for the line maps. The blue colour maps have the same ranges for each separate molecule, with maximum fluxes of 7 ⇥ 10 16 W m 2 for CO,
1 ⇥ 10 15 W m 2 for OH, 4 ⇥ 10 16 W m 2 for o-H2 O, and 1 ⇥ 10 14 W m 2 for [O i].

sion. It is difficult to distinguish the two sources in the PACS
data.
Article number, page 4 of 37

2.2.2. Spectral line maps

The sum of the PACS spectra of the spaxels closest to IRS7A andSMM 1C is shown in Fig. 4. The detected species are indicatedby di↵erent colours. The total emission spectral line emission
maps of some important lines are found in Figs. 2–3, and all
spectral line emission maps can be found in Figs. B.1–B.10. The

CO emission in IRS7 is generally found to be more extended
than the OH and H2 O emission. The emission from the atomic
species C ii and O i detected in IRS7 is even more extended. In
particular, the [C ii] line emission peaks o↵-source in a positioneast of IRS7A. Compared to the CO, OH, and H2 O emission
in the IRS7 cloud, which is mainly centred on the three point
Article number, page 5 of 37