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−4SU Aur at the T Tauri/Herbig
up a polar cavity transition
with µ0 = 0.8 (see Whitney & Hartmann
5.3.
Extendedcircumstellar
8
S.
V.
Jeffersetal.:Imaging
ing at the innermost edge of the disk, resulting in the artificial Additionally,
the determined dust disk mass of 1.2 · 10environmentM!
1993, for a description of these parameters). The resulting modposition. 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!! ), as shown in Figure 3, elled molecular cloud is very faint, and it was not possible toet 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 weclearly shows a nebulosity extending to the west of the im- reproduce the almost constant linear polarised intensity level
Table
3. Values
the model parameters.
boththedata same
sets, though
the data set from
1 January isonused
used exactly
dataonlyreduction
techniquesthe3.3.
obserage. Similarof
structures
SU Aurhave previously been from disk to nebulosity present in the large-scale ExPo image.
typicaland foris approximately
percent
the aroundstellarmass
Data reduction
S. V. 1Jeffers
et al.:
Imaging
SU
Aur
T and
Tauri/Herbig
transitionin this analysis as the atmospheric seeing was much lower.
et al. (2001),Woodgate
etat
al. the(2003),
reportedby Grady
Consequently, the modelled observations derived from the disk
vations andthe modelled data, the dark lane is also present in (Scholz et al. 2006; Williams & Cieza2011).
Nakajima& Golimowski (1995). Since SU Aur is a classical models with and without a remnant molecular cloud do not difsoftware has specifically been developedSymbol
For ExPo, the FLC modulates the orientation of the lin- The data-reduction
Parameter
Value
T
Tauri
star
that
is
still
heavilyaccreting,◦the model ear
imagesthatstate
wereconvolved
withsimulated
ExPo with the aimof grainminimisingatmosphericandinstru- layers of SU Aur’s disk arethe most probable fer significantly.
polarisationof theincoming lightby 90theevery
0.028s forExPo
The
sizes
in
the
surfaceinterpretation
is
that
SU
Aur
is
still
surroundedby
therem10000(2006) concluded
the 10µm
silicate
is uncorrelated
Input
The parameters
aim thatis to separate
the linearlypolarisedfeatureswitching
2). between successive A and B frames, resulting in a to- mental effects.
PSF (Figurenants
of its prenatal
molecular
cloud. To incorporate the efmostly
in the formof smallgrainsthat
followa sizedistribution
signalfromorthatof thestarlight.
The first
step is M
tal of 2000 A + 2000 B frames = 4000 frames. For each frame
infalling
remnant
cloud in our model we 5.4. Disk Model
Mass(star)
1.88
Mmolecular
with
age
the
Hunpolarised
turbulence!fects of anthat!
α equivalent width, indicatingof individual
frames, thenthe four subframes(A or B) the light is separated using a beamsplitter into two the processingwith
addedimplies
an infallingthat
envelope
using ofthe thedensity structure proa power-law
exponentof 4.8.TThis
most
Temperature(effective)
5860
K
eff
posed
by
Whitney&
Hartmann(1993),
withamass-accretion
We determined that the observed SED and ExPo images are
and(Aleftsmall
regeneration
areto important
processes
in disk, Aright , Bgrain
in sequence
form a poorthogonal beams (left and right beams), which are simultaneleft , Bright ) are−7
mass
is incombinedtheformof smallgrains, and because
there
very
(Akesonal. 2005).
We let the cloud fall best reproduced using a disk model with small dust grains
L!etareimage. This(star)
image is then combined at fourL!rate of 3·10 M$ /yr6.3
ously imaged on the EM-CCD, a concept that is the basis of larised intensity
6. Discussion
evolution. Luminosity
This
is confirmed
withrecentresults
fromonto the disk
at the exponential
turnover radius, R0 and opened that follow a very steep size-distribution power law. This imfew
large
grains,maximumcannotbeHerschelclearly
concalibrationdata
to formtheira calibrated
polarisedDsize
a dual-beam polarimeter. For every pair (A and B frames) of FLC angles withDistance
144
parsecintensity image.
The2013)
data analysis
is describedin moreobservations, this results in four different images of SU Aural.
modelsof detailaggregate
fragmentation(Oliveira1000
We have forthe first time directly imaged the circumstellar
en- etstrained.
Theand
grain
sizes
were
obtainedfromfitting
the domiSurface
turnoverpoint
R0
500 aual. (2011)density
and is briefly
summarised
here.
(with a total of 8000 for each FLC angle as described above), by Canovas et& Dominik
2005;at Birnstiel
al. Aur’s2012).
Moredevironmentfor
ofinstancethe classicalAur. Using(Dullemondinterpre- Minimum
nantsilicate∼10 µm inetSU
Spitzer
spectrum(A+B): AleftT, ATauri, BrightSU.
right , Bleftstar
grainfeature
size
amin
0.05
µm
exposure timeeach subexposure
is 4000*
tive models, The
we totaldeterminedtheof inner
radius of
the tailed
dust
disk,
4).
Supporting
simultaneously
with
SU
Aur’s
SED
(see
Figureinformation
on
the
grain
sizes,
derivedfrommodelling
Gas-to-dust
ratio
f
120
3.3.1. Individual frame processing
g/d0.028 s (total=112 s). At the end of every FLC sequence a sethe mass ofriesthe
dustframes
in thethetarget,
surface-density
power
of
small
in thelayers 2ofR!SU
Aur isInnerradiusof thegas wavelengths
disk
Rgas,inobserved
at different
isouter
not possible
given
Eachlaw,
imageimages
isthe
firstpresence
individuallycorrected
forgrainsdark, flatfieldof darkare disk,taken percomprising approxi1 et al.by(2006), who analysed−7
Cosmicareidentified
and removed
8000
the beginningand end of everythenightturbulencea and bias effects.
the power mately
law oftheframes.
dustAtgrain
size distribution,
theworkraysof
Kessler-Silacci
Spitzer
Mass-infall
rateof
the cloudṀinfall10
M! /yrtheunstable
atmosphericseeingofourobservations
in different
set of dome flat-fields were taken with a polarisation calibrator. comparing the two simultaneous left and right images of each
100
parameter, and the mass-accretion rate.
spectraof diskset al.
Polar
openingof thearoundcloud T Tauriµ0stars. Kessler-Silacci
0.8

9

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

Cartoon Models

Polarised Intensity Models

Convolved with simulated ExPo PSF

3

3

P=20%

1
0

-1
-2

-2

-1
0
1
arcseconds

2

-1

3

-3

3

-2

-1
0
1
arcseconds

2

3

3

P=20%

2
1
0

-1

P= 2%

2

arcseconds

arcseconds

0

-3

-3

-2

1
0

-1
-2

-3

-3

-3

-2

-1
0
1
arcseconds

2

3

-3

3

-2

3
P=20%

2
1
0
-1

1
0
-1

-2

-2

-3

-3
-3

-2

-1
0
1
arcseconds

2

3

-3

3

-2

3
P=20%

2

arcseconds

1
0
-1
-2

1
0
-1
-2

-3

-3
-3

-2

-1
0
1
arcseconds

2

3

2

3

-1
0
1
arcseconds

2

3

P= 2%

2

-3

-2

]

filters.
Fitted parameters
Figure 4. Fitted SED> 1000ofµmSU AUR using the model paramePrevious
modelsofthe
thedustcircumstellar
surrounding
Innerradius of
disk
Rdisk
0.17 AUSUµm5. Data are taken from Cutri & et al.
dust,inters described in100-1000
Section
1-10 µm
Aur agreeDustwellmasswiththedisk
results of our work.
ofin theMdust The1.2models· 10−4 M!
0.1-1(2012) (WISE), Høg
etµmal. (2000) (the Tycho-2 Catalogue),
SU Aur - 0.1-1 µm
Akeson et Surface
al. (2002,
2005),
based on interferometricdata
density
power-lawγ
2.4 from
10
Mermilliod(1991) (Johnson UBV), Morel & Magnenat (1978)
Maximum
grain size
66 µm
B5
A0
A5
F0
F5
G0
G5
K0
K5
M0
M5
infrared andmillimeter
wavelengths andamax
infrared photometry,(Johnson UBVRIJKLM), Ishihara
et al. (2010) (AKARI/IRC
Powerof thedisk
size distribution
4.8
Spectral Typeshowed that
SU law
AUR’s
can be fittedpby a passive flat-disk
mid-IR all-sky Survey - L18W+S9W), Ducati (2002)
Turbulence parameterα
9 · 10−3
model, which
agrees well with the model
used in3 ·thisanaly(Johnson
VJKLMN),
Ofek (2008)(Johnson
JKH),
ESA
5. Diameterof imaged
disks asa functionof spectral
Mass-accretion rateṀ
10−7 M! /yr Figure
sis for the Inferred
middleparameters
to outer regions of SU Aur’s disk. We deter- type,(1997)(Tycho-2.BT+.VT),
Droege et al.
(2006)(Cousins by
whereeach observed wavelengthrange
is indicated◦mined the Disk
inclination
of
the
disk
to
be
approximately
50
which
Monet
et
al.
(2003)(USNO-B1.0.B1+.R1),
I,
JohnsonV),
inclination (approximate)i
50 degrees
a different colour and symbol shape. SU Aur is shown as theet al. large
is in line Disk
withposition
the value62◦ axis
determined by Akeson
ISAS/JAXA
all-skyangleofof short
10-20 degrees
black star(2010)
in the(AKARI/IRC
centre of themid-IR
plot. Data
are Survey
taken from
WIDEL+WIDES+N160+N60), Kharchenko (2001) (Johnson(2005). The disk position angle of the short axis is inferred www.circumstellardisks.org.
1(2005).
mass-infall
rate isfrom Akeson
BV), Skrutskie et al. (2006) (2MASS JKsH) and Hauck et al.
from theThe
model
imagestotakenbe between
10et◦ al.and
20◦ which(1990) (Johnson UBV), Andrews & Williams (2005) (450 andmatches the position axis of the long axis of 112±24◦ deter850µm filters).
Akesonet al.
(2005).
mined byplies
that thesize
distribution is dominated by submicron- and consequently the grain mixing in the disk, is significantly
Thesized
innerparticles
radius ofwith
the almost
dust disk
0.17
au grains
agreesand
well
no of
large
dust
thewithupper higher than typically found in disks around T Tauri and Herbig
the valueof size
0.13/0.18
modelled The
by diskAkesonet al.
(2002).
grain
is poorlyauconstrained.
extends
between
0.16 stars.
6.1.
SUcontext
Aur’s disk
In theof protoplanetary disks, SU Aur is a particThe outerdisk500
radiuscomputedby Akeson
et al.
(2002) (deau andau where
the total
dust mass
is computed
to be−4
M! . The
mass240
accretion
rate ofsubmillimeterdisk material onto
interesting
because,with a massof 1.88M% it
rived to1.2
be· 10
between70 andau) usingandthe ularly
The derived
disk example
parameters
were computedby fitting
the SED−7−3
star is3 · 10 M! /yr witha a turbulence
α = 9the· 10 lies
oncomparing
the Herbig
and
T Tauri
An essential
millimeterinterferometricobservations,is parametersmaller thanandtheAe/Be
resultingimage
with transition.
the ExPo images.
This
and radiusa surface-density
lawin(asthis
a functionof radius)
is consequently
much
more
accuratefitting
SED
alone.
betweenthesetwotypesthanis that
T the
Tauri
stars
have
outer diskof 500au power
derivedpaper. An
impor-γ = difference-101
arcseconds

P= 2%

2
arcseconds

arcseconds

1

-2

-3

arcseconds

P= 2%

2

arcseconds

arcseconds

2

Disk Diameter (AU)

6

-1
0
1
arcseconds

2

3

Figure 2. Models of the disk surrounding SU Aur showing the impact of inclination on the morphology of the disk images inlinearly polarised light. The top panel shows a model inclination of i=90◦ , i=70◦ , i=50◦ , and i=30◦ in the lowermost panel. The
images in the left-hand panel show the cartoon interpretation of the disk inclination, the images in the middle panel show theraw model images, and the images in the right-hand panel are the polarised intensity models that have been convolved with asimulated ExPo PSF and processed through the ExPo data reduction pipeline. The intensity scale is logarithmic over two ordersof magnitude and is the same as that of the ExPo images.

3.Classical T Tauri stars: magnetic fields, coronae, and star-disc interactions
C.P. Johnstone1, M. Jardine2, S.G. Gregory2, J.-F. Donati3 and G. Hussain4
Classical T Tauri starの星回り(からdisk内縁)でのmagnetic field, コロナ, star-disk
interactionについて、ZDI magnetic mapとX rayのEM(emission mesure)を用いた結果。
・ 磁場の複雑さは磁場の強さと強い相関がある。弱いと単純な磁場構造で、強いと複雑である。これは星の回転速度と関係している。弱い磁場だと遅い星の回転。diskから星への磁気トルクか、あるいはPMS starの内部進化と星のspin-upに関連しているだろう。
・ZDI mapから得られた磁場構造からX-ray放射の関係がある程度説明できる、ZDI mapでは見られないsmall -scaleの構造がX-ray 放射に関わっているだろうけれど。磁場は星表面から数星半径でhot plasmaを保持することができ、そのコロナは磁場の複雑性と強く関係している。単純なdipolar fieldsはより広がったhot coronal plasmaを保持する。
・測定されたelectron densityは10^8.8~10.0 cm^-3であった。単純な構造のとき小さい。
・open unsigned magnetic flux (Fig 11の青線?)はdipole componentの強さと星の表面積で決まる。強いdipoleと広い星表面積を持つときopen magnetic fluxは大きい。
・CCTsのdiskは数星半径でtruncateする。diskのcorotation半径を決めるのは、dipole
componentの強さが主要因である。
・降着物質の軌道や予測場所は星磁場の複雑さに強く依存する。
・ものによっては、disk内縁とコロナ半径が離れている。今回のmodelではこれについて十分に説明できない。より高度なSD 磁気流体simulationが必要だろう。
Classical T Tauri stars: magnetic fields, coronae, and star-disc interactions 3
Classical T Tauri stars: magnetic fields, coronae, and star-disc interactions

17(a) AA Tau 2009(b) V2247 Oph(c) CR Cha(d) CV Cha

Figure 1. The radial components of the ZDI magnetic maps for AA Tau (Donati et al. 2010a), V2247 Oph (Donati et al. 2010b),
CR Cha, and CV Cha. We show the AA Tau and V2247 Oph maps because they are the most simple and the most complex large-scale
magnetic fields in the sample; we show the CR Cha and CV Cha maps because they have not previously been published.

12

Figure 12. Cartoon showing, for most of the stars in the sample, the coronal extents and the disc truncation radii calculated in
Sections 4 and 5. The stars are ordered from top to bottom by decreasing strength of the dipole component of the field. The horizontal
direction is measured in stellar radii from the stellar surface, marked by the dashed line, to 12 R? . The torque balance radius for AA Tau
is to the right of the plot at approximately 18 R? . The extent of the closed corona, labelled rcor , is shown in green. For each star, thelight blue bar shows the range of radii from which accretion can occur onto the star. This is between the inner edge of the disc, set to
the Alfvén surface radius, and the torque balance radius, unless the torque balance radius lies outside the corotation radius, in whichcase it is assumed to be the corotation radius. For each star, the dark blue bar shows the range of radii where the magnetic field candisrupt the disc outside the corotation radius. The vertical dotted lines show the locations of the disc truncation radii calculated using
the alternative formula of Bessolaz et al. (2008), given by Eqn. 21. As the dipole components weaken, the outer edge of the corona andthe inner edge of the disc moves inwards.

c 0000 RAS, MNRAS 000, 000–000

Johnstone et al.

use new maps produced using the same observational data
For the coronal models in Section 4, we require corobut with an updated ZDI code that fits the maps directly
nal temperatures and X-ray emission measures for each of
to the circularly polarised line profiles as a series of spherthe stars. We take from the literature multi-temperature
ical harmonic components, which is the technique used tocomponent fits to observed X-ray spectra (for references,
fit the maps for the other stars. These maps can be seen in
see Table 1). For most of these stars, these fits have lowFig. 1. Another di↵erence between the fitting method usedtemperature components that we assume are a result of acfor the CR Cha and CV Cha maps and the fitting method
cretion shocks; we therefore disregard these components. As
used for the other stars is the fact that Zeeman signatures
our model assumes an isothermal corona, we combine thefrom accretion proxies were used to fit the magnetic maps ofremaining components to get the coronal temperature andthe other stars, but were not used for CR Cha and CV Chaemission measure using(a) AA Tau 2009(b) BP Tau Feb 2006(c) V2129 Oph 2009
because no Zeeman signatures
were detected in accretion
X
EM =
EMi
(1)
related emission lines for these stars. This could potentiallyiremove magnetic energy from the axisymmetric sphericalP
harmonic components of the field, and move magnetic enEM
log
T
i
i
iP
ergy to higher l-value modes. The e↵ects of not including thelog T =.
(2)EMi
accretion related emission lines in ZDI fits were explored byiDonati et al. (2007) for the case of V2129 Oph.
For AA Tau, the components in the fit to the X-ray specThe stellar parameters for these stars are summarised intrum by Güdel et al. (2007) are all above 10 MK, andTable 1 and the published magnetic maps for AA Tau andare therefore unlikely to have been produced in accretion
BP TauDec 2006
2010(f) TWthat
Hya 2008
V2247 Oph are also (d)shownin Fig.
1. Detailed discussions(e)ofTW Hyashocks;
we therefore assume
all of these components
the adopted stellar parameters for each star can be foundcorrespond entirely to coronal emission. For BP Tau andin the papers (listed above) where the magnetic maps have
CR Cha, the three temperature component fits given bybeen published. As mass accretion can be highly variable
Robrade & Schmitt (2006) both have low-temperature comin time, in all cases we use mass accretion rates derived
ponents at approximately 2 MK, which could be a result of
from the same observations as the ZDI maps. The lowest
X-ray emission from accretion shocks; we therefore disregard
mass star in the sample is V2247 Oph, which has a massthese components. However, these components are only 18%
of 0.36 M and is therefore unlikely to develop a radiativeof the total emission measures, so this only makes a small
core at any age. Other than AA Tau and BP Tau, which aredi↵erence to the coronal models. Similarly, for V2129 Oph,both still fully convective,
the
other
stars
in
the
sampleall
we disregard the three lowest
temperature components of(g) V2129 Oph 2005(h) CR Cha(i) V2247 Ophhost radiative cores.
the nine-temperature component emission measure distribuFigure 8. Surface maps showing the predicted locations of open magnetic field (green), closed magnetic field (yellow) and accretion
footpoints (black) for most of the stars in the sample. The black contour lines show the locations of excess Ca II IRT emission inferred fromc 0000 RAS,MNRAS 000, 000–000
Doppler Imaging; this is considered a good observational constraint on the locations of accretion footpoints. These maps are arranged
from left to right and from top to bottom by decreasing strength of the dipole components of the field.

accretes onto the central star. Since some disc material will
be lost in the form of disc winds, and not accrete onto thecentral star, the measured mass accretion rates underestimate the rate at which mass moves inwards through thedisc.
For comparison, we also calculate the disc truncation radii using the alternative formulation of
Bessolaz et al. (2008), given by2

rtrunc = 2ms7

⇣

B?
140

⌘ 47 ✓

Ṁa
10 8◆ 27 ⇣

M?
0.8⌘ 17 ⇣

R?
2

⌘ 57

(21)

where ms ⇠ 1, and the quantities B? , Ṁa , M? , and R? are inunits of G, M yr 1 , M , and R respectively. The resultsfrom these calculations can be seen and compared with the

When the inner and outer radii from which accretion
can occur are estimated, we use PFSS extrapolations withthe source-surface radius set to the corotation radius, andtrace field lines from every point in the disc from which accretion can occur and assume that all such field lines support
accretion. We assume that accretion only occurs along thefield line in the direction at which the component of e↵ective
gravity along the field line in a coordinate system corotatingwith the star points inwards towards the stellar surface.
Johnstone et al. (2010) showed that the possible absence of magnetic flux on small-scales in the ZDI magnetic
maps is unlikely to have a significant e↵ect on accretion
models. This is because accretion is a↵ected primarily by
the large-scale structure of the stellar magnetic field, withthe small-scale structures having little significance.

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 Herscheling 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 A.
were
obtained.
inKarskaet al.Cooling
2013: Far-infraredmolecular lines from Low- to High-Mass Star Formother 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
Fig. 4. Relative contributions of [O i] (dark blue), CO (yellow), ∼ 104 -105 L# for our sources (van der Tak et al. 2013). The toH2 O (orange), and OH (red) cooling to the total far-IR gas cool- tal far-IR line cooling ranges from ∼ 1 to ∼ 40 L# , several oring at central position are shown from left to right horizontally ders of magnitude less than the dust cooling. Relative contrifor each source. The objects follow the evolutionary sequence butions of oxygen atoms and molecules to the gas cooling are
illustrated in Figure 4. Atomic cooling is the largest for the morewith the most evolved sources on top.
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).
van der Tak et al. 2013). The only H2 O Fig.
lines5.observedin comRotational
diagramsof CO for all objects in our sample. The base 10 logarithm of the number of emitting molecules from a
mon by the two instruments within the WISH program are the Typically, atomic cooling accounts for ∼20 % of the total line
level
u,
N
,
dividedbythedegeneracy of the level, gu , iscoolingshownisasdominateda functionbyofCO,energy of the upper level in kelvins, Eup . Detections
uground-state transitions: 212 − 101 at 179.5 µm (1670 GHz) and far-IR line cooling. Molecular line
are shown
filled circles,
whereasthree sigma
upper
areLFIRL
shown
which
is responsiblefor ∼ 15
up tolimits
85% of, withasa empty
mediancircles. Blue lines show linear fits to the data and221 −112 at 180.5 µm (1661 GHz), both dominatedby asabsorptions
corresponding
rotationaltemperatures.
line inthe G34.26+0.15
contributionof 74%.The
H2 Overticaland OHredmediancontributions
to the panel shows the dividing line between the warm
and therefore not optimal to estimate tothe
what
extent a complex
far-IRincoolingare lessthan 1%,ofbecause
manyYSOs.
of theirErrors
transitions
andresolution.
hot components
rotational
diagrams
low-mass
associated with the fit are shown in brackets.
water profile is diluted at the PACS spectral
However,as seen
HIFI observations of the lines between excited rotational states, are detected in absorption. Assuming that the absorptions arise
which dominate our detected PACS lines, are generally in emis- in the same gas, as found for the case of Orion (Cernicharo et al.
sion at the longer wavelengths probed by HIFI. In the case of 2006), they therefore do not contribute to the cooling, but heat12of the 7gas.
Still,Hthe
H2 O to the
total
FIR
lines arewith
shown
shows
diagrams foroftheTherotational
diagrams show a substantial scatter, larger
CO 10–9, the line profiles of YSOs observed
HIFIseparately.
con- ingFigure
2 O contribution
cooling
increases slightly for more evolved sources,
fromerrors∼5% of individual data points, caused by large opacsist of a broad outflow and a narrow quiescentcomponent
withall H
three sourceswherethan the2 O lines are seen in absorption.
(DR21(OH))
to30%(W3IRS5),
whereas
no
such
trend
is
seen
the relative fraction of the integrated intensity of the narrow to
broad components being typically 30-70% for low-mass sources for OH.
(Yıldız et al. 2013). For the single case of8 a high-mass YSO, W3
IRS5, this fraction is about 50% (San José-Garcı́a et al. 2013).
4.2. Molecular excitation
To summarize, PACS spectra of high-mass sources from oursample show detections of many molecular lines up to high ex- Detections of multiple rotational transitions of CO, H2 O andcitation energies. CO, H2 O, OH, and CH lines are seen towards OH allow us to determine the rotational temperatures of theall objects, whereas weaker lines of other molecules are detected emitting or absorbing gas using Boltzmann diagrams (e.g.
towards less than half of the sources. CO lines are always seen Goldsmith & Langer 1999). For H2 O and OH, the densities arein emission, CH in absorption, whereas other species show dif- likely not high enough to approach a Boltzmann distribution andferent profiles depending on the transition and the object. Table therefore the diagrams presented below are less meaningful.
Emission line fluxes are used to calculate the number of
C.1 shows the CO line fluxes for all lines in the PACS range.
emitting molecules, Nu , for each molecular transition using
Equation (1), assuming that the lines are optically thin. Here,
4. Analysis
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 light
4.1. Far-IR line coolingand h is Planck’s constant:
Emission lines observed in the PACS wavelength range are used
4πd 2 Fλ λto calculate the contribution of different species to the total line Nu =(1)hcA
cooling from high-mass protostars. Our goal is to compare thecooling of warm gas by molecules and atoms with the cooling The base 10 logarithm of Nu over degeneracy of the upper levelby dust and connect them with the evolutionary stages of the gu is shown as a function of the upper level energy, Eu , inobjects. Relative contributions to the cooling between different Boltzmann diagrams (Figures 5 and 6). The rotational tempermolecules are also determined, which can be an indicator of the ature is calculated in a standard way, from the slope b of thephysical processes in the environments of young protostars (e.g. linear fit to the data in the natural logarithm units, T rot = −1/b.
Nisini et al. 2002; Karska et al. 2013).
Because the size of the emitting region is not resolved by our
We define the total far-IR line cooling (LFIRL ) as the sum instrument, the calculation of column densities requires addiof all emission line luminosities from the fine-structure [O i] tional assumptions and therefore only the total numbers of emitlines (at 63 and 145 µm) and the detected molecules, follow- ting molecules is determined. The formula for the total number

・COのrotational diagramから、単一のwarm componentであり、300Kでlow massと同じ。いくつかの天体では、hot componentが存在を否定できない(high-J COの制限から)。
・quiescent envelopeからの放射は45~85% CO luminosity。low massでは10%程度。
highest-J COを説明するには、schockとかのcomponentが必要。

・H2Oのrotational diagramでTrot=250K、emission/absorption両方とも。この温度は100K程度で、low massより高い。光学的厚みや減光の効果でエラーは大きい。

・OHから80Kで、low massとは同じだけど、中間質量とはだと45Kほど大きい。H2Oと同様にsub-thermally excited。

・H2O 303-212とCO14-13、に対するbolometric luminosityは6-7のorderで相関が見られる。この相関は共通の物理メカニズムが存在していることを示唆する、low-massで言われているようなnon-dissociative schockとか。

6

・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。

A. Karska et al. 2013: Far-infrared molecular lines from Low- to High-Mass Star Forming Region

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" , andhigh−J CO do not exclud
hot component in several so
3. Emission from the qui∼45-85 % of the total C
PACS range. The correspo
less dense envelopes of lo
Additional physical compnecessary to explain the hig
4. Rotational diagrams of H2 O
K for all sources both fro
This temperature is aboutsources, likely due to thesources. The diagrams sho
citation and optical depth e
5. OH rotational diagrams ational temperatures of ∼80
YSOs, but higher by ∼45 K
jects. Similar to H2 O, lines
6. Fluxes of the H2 O 303 − 2strongly correlate with bolo
masses over 6 and 7 orders
correlation suggests a commsible for the line excitation
shocks based on the studies
7. Across the large luminosit
far-IR line cooling strongl
luminosity, in agreement w
The ratio of molecular andfor all those YSOs.
8. Because several H2 O lines
cooling ratio decreases fro
to high-mass YSOs.

is 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-

Fig. 13. Correlations of line emission with bolometric luminosity (left column) and envelope mass (right column) from top tobottom: CO 14-13, H2 O 303 -212 , [O i] at 145 µm and OH 163 µm
line luminosities. Low- and intermediate-mass young stellar objects emission is measured over 5×5 PACS maps. Red and blue
circles show Class 0 and Class I low-mass YSOs from Karska etal. (2013). Navy diamonds show intermediate mass YSOs from
Wampfler et al. (2013; O and OH lines) and from Fich et al.
(2010; CO and H2 O line). High-mass YSOs fluxes are measured
in the central position and shown in black. Pearson coefficient r
is given for each correlation.

tion is expected to be higher due to the carbon ionizing and CO
dissociating FUV radiation. In Section 4.1 we estimate the [C ii]
cooling in two high-mass sources to 10-25% of LFIRL .
The ratio of cooling by gas (molecular and atomic lines)
and dust (Lbol ) decreases from 1.3·10−3 to 6.2·10−5 from low tohigh-mass YSOs. It reflects the fact that H2 O, and to a smaller
extent OH, contributes less to the line luminosity in the highmass sources, because many of its lines are detected in absorption. The detection of H2 O and OH lines in absorption provesthat IR pumping is at least partly responsible for the excitation of these molecules and the resulting emission lines (e.g.
Goicoechea et al. 2006; Wampfler et al. 2013). In the case wherecollisions play a marginal role, even the detected emission linesof those species do not necessarily cool the gas.
The above numbers do not include cooling from molecules
outside the PACS wavelength range. This contribution can be
15

Fig. 14. From top to bottom: (1) Total far-IR line cooling; (2)Ratio of molecular to atomic cooling; (3) Gas to dust coolingratio, LFIRL /Lbol , where LFIRL = Lmol + Latom ; as a function ofbolometric luminosity. Low-mass YSOs are shown in red (Class0) and blue (Class I), whereas high-mass ones are in black.

significant for CO, where for low-mass sources the low-J lines
are found to increase the total CO line luminosity by about 30%(Karska et al. 2013). Thus, the CO contribution to the total gas
cooling is likely to be even larger than suggested by Table 2.

6. Conclusions
We have characterized the central position Herschel/PACS spectra of 10 high-mass protostars and compared them with the results for low- and intermediate-mass protostars analyzed in asimilar manner. The conclusions are as follows:
1. Far-IR gas cooling of high-mass YSOs is dominated by CO(from ∼ 15 to 85% of total far-IR line cooling, with median
contribution of 74%) and to smaller extent by [O i] (with median value ∼20 %). H2 O and OH median contributions to
the far-IR cooling are less than 1%. In contrast for low-mass
YSOs, the H2 O, CO, and [O i] contributions are comparable. The effective cooling by H2 O is reduced because many
far-IR lines are in absorption. The [O i] cooling increases formore evolved sources in both mass regimes.
2. Rotational diagrams of CO in the PACS range show a single,
warm component, corresponding to rotational temperature of∼300 K, consistent with low-mass YSOs. Upper limits on
16

Acknowledgements. Herschel is an E
struments provided by European-led
important participation from NASA
Christiane Nüsslein-Volhard-Foundati
German Commission for UNESCO via
and JC thank the Spanish MINECO fo
07304, AYA2012-32032 and CSD2009
Ramoón y Cajal research contract. Ast
Netherlands Research School for Astr
Academy of Arts and Sciences (KNAW
grant 614.001.008 from the Netherlan(NWO), and by the European Comm
FP7/2007-2013 under grant agreement

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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
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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)$$Kóspál et al.: RV variations in EX Lup$
and HARPS(High Accuracy Radial velocity Planet Searcher)
3. Results and analysis
1.0
望遠鏡MPG/ESO 2.2m, 3.6m telescope,
La silla observatory,Chile$

波長域：3500-9200Å, 3780-6910Å

GLS Power$

3.1. Period analysis of the RV data and orbital s

0.8

To search for periodicities in the RV values m
the absorption lines, we used the generalized(GLS) method (Zechmeister & Kürster 2009). S
0.2
ber of available HARPS points is not sufficien
an independent RV period determination, we fi
0.0
the FEROS data only, and afterwards we combin
0.5
and HARPS data points, and repeated our anal
0.4
points were weighted with the inverse square of t
0.3
ties. Figure 1 shows the GLS periodogram calcul
FEROS points. Two significant periods with false
0.2
bility (FAP) of < 10−4 are visible, one at around 1
0.1
an alias caused by the sampling (see the windowbottom panel of Fig. 1), and another one at 7.417
0.0
curve phase-folded with the 7.417-day period is pl
1
10
100
1000
Period / dFigure 2 exhibits a clear periodic signal withamplitude of ≈2.2 km s−1 . The shape of the curve
Fig. 1. Top: GLS periodogram of the FEROS RV values, show- with the increasing part approximately 1.7 times l
ing the best period of 7.417 days. The horizontal dashed line in- decreasing part. In the figure, we color-coded th
dicates the FAP level of 10−4 . Bottom: window function.
serving seasons during our five-year monitoring
Kóspál et al.: RV variations
in EX Lupsystematic differences can be seen between thesons, which suggests that the period, phase, and3.
Results
andanalysis6
the RV variations were stable for at least four y1.0
2007-2(absorption lines)
2009-1
2009 and 2012. The three data points from 2007 a
2009-20.8
3.1. 4Period
analysis of the RV data and orbital solution
2010-1tent with the later observations, implying that the
2010-22 2011-1
were largely unaffected by the largest outburst o
2012-2 for periodicities in the RV values measured from
10−4
0.6
To searchobserved which occurred in 2008.
0
the absorption lines, we used the generalized Lomb-Scargle
The shape of the phase-folded RV curve sug
0.4-2(GLS) method (Zechmeister & Kürster 2009).
Sincesibile
the existencenum- of a companion on an eccentric
FEROS
Lup. We0.2
ber -4of available HARPS points is not sufficient
toEX
carry
outfitted a Keplerian solution to the FER
HARPS
both GLS and the idl code RVlin4 (Wright & How
an independent
RV period determination, we first considered
1
0.0
resulting parameters are listed in Table 1. The R
0
the FEROS
data only, and afterwards we combined the
FEROS
0.5
handledata from different instruments by fitting-1
the relative offset between them, thus, w
and
HARPS data points, and repeated our analysis.correct
The fordata
0.4
combined FEROS+HARPS data sets, and detepointswere weighted with the inverse(emission
squarelines)of theirthe
uncertain2
0.291 km s−1 between the two instruments.
0.3
ties. Figure 1 shows the GLS periodogram calculatedset
forofthe45
calculates uncertainties for the fitted parameters
0
FEROS points. Two significant periods with false alarmprobastrapping
algorithm (Xuesong Wang et al. 2012).
0.2-2bility(FAP) of < 10−4 are visible, one at around 1 day,whichdeterminedisrameters
from the fit for the combine
also listed
in Table 1. The values derived for the
0.1
an aliasby the sampling(see 1.5
the window2.0functionin the0.0 caused 0.5
1.0
setsTheusingRV
different methods are all consistent wi
Phaseanother one at 7.417 days.
bottom panel of Fig. 1), and
0.0tainties.
curvephase-foldedwiththe
7.417-dayperiodisplottedinFig.
2. m sin i and the semi-major axis a,1
10100
1000 Fig. 2. Top: Best Keplerian fit to the absorption line RVs for the
To calculate
Period / dFigure2 exhibitsa clearperiodicsignalwith
RV semicombined
FEROS
and HARPS
RV dataset (with
RVlin).
The anstellar
mass of M∗ = 0.6 M$ for EX Lup (Gras-Vel
different
colorsofindicate
different
observingstands2005). Figure 2 shows the best fit for the comb
amplitude≈2.2 kms−1 . The
shapeseasonsof the(–1curve
is asymmetric,
the first
halfof the year,–2 approximately
for the second half).
Circles
repwhich
hasthea FAP of 6.7×10−27. After subtracting
Fig. 1. Top: GLS periodogram of the FEROS RV values, show- forwith
theincreasing
part
1.7
times
longerthan
resentFEROS
andsquares
HARPS
data
points.
Middle:residua
periodKeplerian fit from the RV data, we searche
ing the best period of 7.417 days. The horizontal dashed line in- withdecreasing
part.
In
the
figure,wecolor-codedthedifferent
obthe same color and shape coding. Bottom: Emission line odicities in the radial velocity residuals, but found−4
dicatestheKóspál et al.: RV variations
in EX
LupFAP level of 10 . Bottom: window function.
servingduringourcoding.
five-year monitoring program.
NoFAP < 10−4.
RVs
with theseasonssame colorand shape
periods withsystematic differences can be seen between the different
sea- in Sect. 2.1, we calculated RVs f
As describedrow emission
sons, which suggests that the period, phase, and amplitudeof lines that are isolated and not sup
80
broaderabsorption features. Similar but weaker6
the RV variations were stable for at least four years,
between
2007-2(absorption lines)
may distort the absorption lines and might intro
2009-1
2009 and 2012. The three data points from 2007 are also
consis-despite our best efforts to discard th
Prot = 17 d
RV signals,
4 2009-2
60
2010-1
tent
with
the
later
observations,
implying
that
the
RV
variations
sorptionlines from our analysis. To make sure tha
2010-22MASS
magnitudes
by combining
dither outburst
measurements
2 2011-1
performed a period analysis on the RV
were largely
unaffectedby thethefivelargestof EXcase,
Lupweever
2012-2of both comparison stars using an outlier-resistant algorithm. No sion lines. The algorithm found no significant per
observedwhichoccurred
in
2008.
40
0
color term was needed in this transformation. The uncertainty of a small peak is present at about 7 days with FA
The
shape ofis thephase-folded
curve uncersuggests
posthe final
photometry
the quadratic
sum of RV
the formal
Fig.the2, weplotted the RVs of the emission lines
Prot = 7.4 d-2sibileexistenceof a companion
onofantaintyof theaperture photometry,the scatterthe eccentric
photometry orbit
with aroundthe 7.417 d period. The data points have
20
FEROS
ofEX
the Lup.
individual
ditherframes,andphotometriccalibraion.
The
We fitted a Keplerian solution to the FEROS RVs
using
-4
HARPS
4
resultingphotometry
listedin Table
3. 4 (Wright & Howard 2009).
http://exoplanets.org/code/
both GLS
and theis idl
code
RVlin
The

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、$

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

5n the light curves ofs made to reproduces, see Sect. 3.3).

tail whether the obor hot spots on thers like the BVS, BC,
iodicity. Nor do wen wavelength withine constructed a simRV semi-amplitude.
have to cover a comd atypical solution.

RVem [kms-1]

RV [kms-1]

Inclination [deg]Window power

O-C [kms-1]

RV [kms-1]

GLS Power

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

O-C [kms-1]

4

0.4

$

0
0

1
0
-1

2

4
6
v sin i [km/s]

810Fig. 9. Inclination vs. v sin i. The solid lines indicate the period of
7.4 d (obtained from the RV analysis) and 17 d (a lower limit ob(emission lines)
tained from
2 the inclination constraint). The hatched area marks
the v sin i constraint from the FEROS and HARPS spectra; theorange and0 red areas mark the constraints for disk inclinationfrom Sipos et al. (2009) and Goto et al. (2011), respectively.
RVem [kms-1]

3

10−4

Window power

$

0.6

-2

inconsistent with the spot scenario. We conclude that the expla0.0
0.5
1.0
1.5
2.0
nation of the measured RV curve with photosphericspots is anPhaseunlikely hypothesis.
We also checked in Sect. 3.2 whether flows and inhomoFig.
2. on
Top:
Best surface
Keplerian
fitchromosphereto the absorption
geneities
the stellar
or in thecould pro-line RVs for the
combined
FEROS
and HARPS
RV datasetperiod(with RVlin). The
duce the measured
RV variations.
In this scenario
the RV
again
would beidentical
to the stellar
rotation observing
period. EX Lup,asdifferent
colorsindicate
different
seasons(–1 stands
anstar, isprobably
active.
However,
forM-typethe first
halfof thechromospherically
year, –2 for thesecondhalf). Circles repthe phenomena responsible for the distortions of the line profiles
resent
FEROS
squares
HARPS
points.
andthus forthe RV andvariations
are not
expecteddata
to survive
moreMiddle: residua

resulting parameters are listed in Table 1. The RVlin code can
4 handle data from different instruments by fitting parameters tocorrect for the relative offset between them, thus, we used it to fit
the combined FEROS+HARPS data sets, and determined an offset of 0.291 km s−1 between the two instruments. The code also
calculates uncertainties for the fitted parameters using a bootstrapping algorithm (Xuesong Wang et al. 2012). The orbital parameters determined from the fit for the combined data sets arealso listed in Table 1. The values derived for the different data
sets using different methods are all consistent within the uncertainties.
To calculate m sin i and the semi-major axis a, we assumed astellar mass of M∗ = 0.6 M$ for EX Lup (Gras-Velázquez & Ray
2005). Figure 2 shows the best fit for the combined data set,
which has a FAP of 6.7×10−27. After subtracting the 7.417-day
period Keplerian fit from the RV data, we searched for sine peri-

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

Fig. 4. Steady-state abundances of NH, NH2 , and NH3 as a function of the C/O gas-phase elemental abundance ratio for three different valuesof gas-phase elemental abundance of sulphur (from left to right: [S]tot = 8.0 ×10−8 , 8.0 ×10−7 , 8.0 ×10−6 ). C/O values include 0.3, 0.4, 0.5, 0.6,
0.67, 0.7, 0.75, 0.8, 0.9, 1.0, 1.1, 1.2, 1.3, 1.4, 1.5. In each panel, the solid line refers to the model predictions and the vertical dashed lines locate
C/O=0.67 (Flower et al. 2006 value) and 0.8 (our best model value). For NH2 and NH3 , the ortho and para abundances are shown separately(dashed and dotted lines respectively). The observed abundances towards IRAS 16293-2422 are represented by the hatched boxes, considering an
o/p ratio of NH2 ∼ 2, as predicted by our model.

R. Le Gal et al.: Interstellar chemistry of nitrogen hydrides in dark clouds

structed (e.g. Crimier et al. 2010, and references therein). The re- et al. 1995). In the most external part of the source where nigion can be divided into four components: the ambient molecular trogen hydrides are seen in absorption, the column
density isa.
Table
3. theAbundances
and abundanceof nitrogen
hydridesat steady-statepredicted
chemical
networks
cloud,
circumbinary
envelope, ratios
the threeprotostellar
sourcesnecessarily
muchusinglower.several
Radiative
transfercomputations
where
A1, A2 and B, and the outflow components (see Loinard et al. we have coupled non local thermodynamic equilibrium (non2013, and references therein). The total column
density ofb H2 LTE) calculations with a simple
physical model of the source
Observations
Modelstowards IRAS 16293-2422 is ∼ 2 ×1023 cm−2 (van Dishoeck (see Appendix A) have shown that the absorbing region mustN ( cm−2 )N/NHDislaire et al. (2012) c
This workd
Flower et al. (2006a)e
NH
NH2 (o)NH2 (p)NH2
NH3 (o)NH3 (p)NH3
NH2 :NH
NH3 :NH
NH:NH2 :NH3

2.2±0.8(14)5.0±0.9(13)–7.5±3.0(13)––1.4±1.0(15)7.3±4.5(-9)
1.7±0.9(-9)–
2.5±1.6(-9)––
4.7±4.1(-8)

4.3(-10)––8.2(-11)––7.5(-10)1.2(-8)2.5(-9)1.1(-9)3.6(-9)5.8(-9)8.6(-9)1.4(-8)1.3(-9)2.0(-9)2.0(-9)4.0(-9)2.5(-9)7.3(-9)9.8(-9)0.34±0.18
6.4±5.2
3:1:19

0.2
1.7
5:1:9

0.3
1.2
3:1:3

3.1
7.5
0.3:1:2.59aConstant physical conditions were adopted: T kin =10 K, nH = 104 cm−3 , ζ = 1.3 ×10−17 s−1 . Numbers in parentheses are powers of ten.
Based on Hily-Blant et al. (2010a) and revised as described in the Appendix B. For each species, N is the total column density. The adopted value of NH = 3.0±1.5 ×1022 cm−2 .
Updated version of the osu.09.2008 network used in Dislaire et al. (2012), who adopted C/O = 0.67 and [S]tot = 1.5 ×10−5 .
Our network with the parameters of our best model: C/O=0.8 (or [O]tot =1.04 ×10−4 ) and [S]tot =8 ×10−8 .
eNetwork from Flower et al. (2006a) adopting the parameters of our best model.
bcdhave a H2 density ! 104 cm−3 (i.e. nH < 2 ×104 cm−3 ), which
does not correspond to the circumbinary envelope but to theambient cloud (although the distinction is artificial since thedensity profile is continuous). For this component, the H2 column density is rather uncertain but we adopted the value derived from C18 O observations by van Dishoeck et al. (1995),
N(H2 ) = 1.5 ×1022 cm−2 . We note that Hily-Blant et al. (2010a)
employed a different value, N(H2 ) = 5.5±2.5 ×1022 cm−2 , basedon an extrapolation of the Crimier et al. profile. The NH3 column density derived by Hily-Blant et al. (2010a) was highly
uncertain because, in contrast to NH and NH , the hyperfine

plotted in Fig.5. At low and intermediate sulphur abundance,
the NH2 :NH ratio delineates a narrow range of C/O values of
0.75 − 0.80 consistent with the above constraints. At intermediate sulphur abundance the NH3 :NH ratio further constrains
the C/O ratio to 0.75. For a high sulphur abundance, however,there is no C/O which allows to reproduce simultaneously bothabundance ratios. From Fig. 2, which shows the evolution withtime of the abundances, we note that the observed abundance ratios NH:NH2 :NH3 are also well reproduced at very early times≈ 4 ×104 yr. However, at such early times, the fractional abun-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 Vは質量降着率、Vはbackgroundのgasとの相対速度。
distant
ect to distant background gas. Our force law holds even when the object is porous (e.g., a galaxy)elling mass in aこの結果はporous(多孔,銀河等)でも成り立ち、惑星形成から銀河の力学で適用できる。
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 fromcase, which is often the most interesting one astrophysically.
That
is,
they
took
V
>
c
,
where
V
is
the
speed
of
the
s an overdense wakebehind the object obtained
gravitating mass relative to the distant background, andcthe force. Do
the sound speed in that gas. (In the hypersonic calculations
ravitational pull, retarding
thesuccessors,
originalof Bondi, Hoyle, and theirc was implicitly
settigatedthis hypersoni
to zero.) While their answers differed in detail, all agreed
ical friction arisesthat
whetherthe friction forcethe
in this medium
regime varies as V , witha
Cantó
et al. (2011). Do
coefficient that includes a Coulomb logarithm. This latter
racting point particles,e.g.,a derivedstellarterm also appears
in the force
by Chandrasekhargas. He(1943), and arises from an integration in radius awaytemperaturefromthe mass.
In all derivations, at least
one of the integration
uum fluid, e.g., aninterstellar
cloud.
the total power emitt
limits is rather ill-defined.
3) provided the essential
theoryThe previous studies
made two whenimportant, simplifying(see also Rephaeli & S
assumptions. First, they neglected any accretion of backollisionless, and his
been the(1971)
groundsolution
gas by the moving has
object. Quantitatively,
asused an impu
sumption was that R " r , where R is the object’s physical radius
and the accretion
radiusgalaxr≡ 2 G M/V is thestudies of both starclusters
andframe of the backgroun
distance from the mass M within which its gravity qualigravitating object with gas also Article
occursnumber, page 1latedof 21
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
s

s

s

−2

acc

acc

2

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。highmassだと予想されるより高いaccretion ratioの傾向があった。
12 B. Lim et al.
-0.53

aperture photometry for the standard stars with an aperture size of 14.00 0 and present the coefficients in Table 2. We
carried out point spread function (PSF) photometry for target images using IRAF/DAOPHOT with the small fitting
radii of 1 FWHM (6 1.00 0), and then aperture correction wasmade by applying the di↵erential magnitude between the fitting radii and the standard aperture radius (7.00 0). The sizeof the aperture correction was determined from the aperture
photometry of relatively bright, isolated stars with a photometric error smaller than 0.01 mag for each target image.
Finally, the instrumental magnitudes were transformed tothe standard magnitude and color indices using the transformation equations below (Sung et al. 2008),
M =m

(k1

k2 C0 ) · X + ⌘ · C0 + ↵ · UˆT + ⇣(1)

where M , m , k1 , k2 , ⌘ , C0 , X, ↵ , UˆT , and ⇣ are
the standard magnitude, instrumental magnitude, primary
extinction coefficient, secondary extinction coefficient, transFigure 9. The Hertzsprung-Russell diagram of IC 1848. Several
formation coefficient, relevant color, air mass, time-variation
isochrones (0, 0.5, 2.0, and 5.0 Myr) are superimposed on thecoefficient, time di↵erence relative to midnight, and photodiagram with several evolutionary tracks (Ekström et al. 2012;
metric zero point, respectively. A negligible time-variation
Siess et al. 2000). The other symbols are the same as Figure 2.
Figure 11. The IMF of IC 1848. Circles and the dotted line repFigure 1. Finder chart for IC 1848. The size of the circles areof ⇠ 1 mmag per hour in the photometric zero points
resent was
the IMF derived from members only and members plus
proportional to the brightness of individual stars. The position
found in both the I and V bands. We adopted the candidates
recentlywithin the pre-main sequence locus, respectively. The
IMF of NGC 2264 (dashed line – Sung & Bessell 2010) is overof stars is relative to the brightest O type star HD 17505 (↵ =modifiedtransformation
coefficients(⌘ )medianfor SNUCam
fromplotted.
Dot-dashed and long dashed lines correspond to the comthe
temperature
derived fromthespectral type – e↵ective
Figure10.
Age
distributionofPMS
members.
The
age
h
m
s
0
00
02 51 08 .0 = +60 25 04 , J2000). Squares outline thepleteness limits for membership and photometry, respectively. See
Limet al.
theThisassumption
the luminostemperature relation and the color – e↵ective temperature
is about
3 Myr
with(2009).
a spreadFromof 6 Myr.
age is ratherthat
youngerthe main text for details.
observed regions.
than that
estimated forthe massive
ity functionof allstars instars.
the observed fields has a linear
relations with an appropriate weight. Finally the bolometriccorrections were obtained from the relations in Paper 0. We
slope across the entire magnitude range we found that our
4 THE INITIAL MASS FUNCTION
present the HRD of IC 1848 in Figure 9.
photometry is 88% complete down to V = 19 mag,whichphotometricand the tracks
adoptedspectraltypes
Several data
evolutionary(dashedlines)
withofanthe
ini-stars
As 15010
the IMF is a very important tool to understand starcorresponds to ⇠ 1.2M . The photometric data foraretialpresented
Table 1. (solid lines) were over plotted information processes, a number of studies have attempted
a medianof about
Myr. Given are
the agespread of– 6 electronicmass andinisochrones
starsagefrom
this 3observation
available
in 3thedescribed
in Limet al.models(2008),of all
pre-processing
Myr found in other young open clusters, such as NGC 2244 to derive the IMF in di↵erent star forming environments.
theAs
figure.
The stellar
evolution
Ekström
et al.
table (Table 3) or from the authors (B. L. or H. S.).
take into accountthewase↵ectof stellarto(2012),
removewhich
instrumental
artifacts
doneusingrotation
the IRAF (Park & Sung 2002), NGC 2264 (Park et al. 2000), NGC Bastian et al. (2010) reviewed IMF studies and the relationwith environmental conditions and mass range. Since
We checked our photometry against previousshipstudies
1 on the evolution of stars for a given mass, were used for post/CCRED packages. In order to obtain daily parameters 6530 (Sung et al. 2000), and Trumpler 14 and 16 (Hur et al. IC 1848 is one of the large-scale star forming regions in the
inthe
order
confirm
its consistency.
a few photoelec2012),
age tospreadestimatedin this work Only
is reasonable,
and MS stars, towhile
used thesystem,
models of
et al.
forMS
transformation
thewestandard
theSiess
atmospheric
Perseus spiral arm, this object must be an ideal laboratory
tric thephotometricdataforsets
for
ICstrongly
1848 were
foundto inthealthoughage
determined
PMS
stars
depends(2000)
for
PMS
stars.
The
isochrones
(0.5,
2.0,
and
5.0
Myr)
study
the IMF in a low metallicity and low spatial density
extinction coefficients and photometric zero points were ob2
open
data
base used.
WEBDA. In10,addition,
photo- Given the fact that the cluster contains a few
on thePMScluster
evolutionmodelsIn Figurethe long noenvironment.
have been constructed from the two models. The stars withtained on the same night as the cluster observations throughmassivestars it is possible to derive the IMF over a wide
tail towardagewithmayaarisefrom CCD
overestimating
the agemasses larger than 18M seem to be moving away from themetricolder
study
modernwas reported
for IC1848.
theobservationofmanyequatorial
standardstars(Menzies
range.
of stars
the Kelvin-Helmholtz
contraction
phasein theMS stage, while intermediate-mass stars (⇠ 3M ) seem to
Theinphotoelectric
photometricdata of
Hoag
et al.mass(1961)For all post-MS or MS members, the masses were esetbe
al.approaching1991) at di↵erent
airThemasses.
Wewithperformed
PMSweremodels(Sung et al.
1997,2004),
fromfor
stars
with
an
the ZAMS.
isochronean age ofsimple
5
compared
with
our
CCDanddata
9-11
commonstars.
timatedby comparing their position in the HRD to postedge-on
lying below
locus (Sungthe
et al.
Myr well predicts the position of the most massive stars intracks. The mass of PMS stars, or stars in
Thediskdi↵erencesin thethe PMSV magnitude,
B 2008).
V , andMSUevolutionaryB
It is worth noting that some of stars in the long tail may be the PMS locus without any membership criterion, were estithe HRD. However, according to Hillwig et al. (2006), HD
color indices are, respectively, V = 0.021 ± 0.020 mag,
1 17505,
17520, and BD+60
are binaries
or multiple
Image HD
Reduction
Analysis586
Facility
is developed
and dis- real, old PMS stars, since the star formation history of the mated by comparing their position in the HRD to PMS evo(B could
V ) =be0.001±0.054
mag,(UunderstandB) = 0.008±0.067
lution model tracks. In order to derive the IMF we counted
W5 regionquite complex.
In andorder tosystems, and therefore it is probable that the turn-o↵ age

The young open cluster IC 1848

The young open cluster IC 184813

of members in the low-mass regime is very similar to that ofprevious studies for other young open clusters, such as NGC
2244 (Park & Sung 2002), NGC 2264 (Park et al. 2000),
NGC 6530 (Sung et al. 2000). With complete membershipcriteria, Sung & Bessell (2010) showed a smooth increasein the IMF of NGC 2264 down to 0.25 M compared tothat of Park et al. (2000). Sung et al. (2000) and Sung et al.
(2009) have discussed the efficiencies of detecting PMS stars
using H↵ photometry (47 % from a single observation) andMIR excess (38 %), respectively. Hence we attribute the dip
in the IMF of IC 1848 to the incomplete membership selection for the stars in the Kelvin-Helmholtz contraction phase,
while the incompleteness of our photometry may influence
the IMF in the lower mass regime (< 1.2M ). The IMF ofthe members (> 1.2M ) should be a lower limit, whereasthat of all stars within the PMS locus may be an upper limit
in the same mass range.
We first derived the IMF of IC 1848 complete down to
3M through this study. The slope ( ) of the IMF is about
1.6 ± 0.2 for log m > 0.5. It appears to be steeper than the
single IMF (Salpeter 1955) and the Kroupa IMF (Kroupa2002), but close to that of Scalo (1986) for massive stars.
The slope of the IMF for IC 1848 is also steeper than thatof other young open clusters ( 1.3 ± 0.1 for NGC 6530 –Sung et al. 2000, 0.7 ± 0.1 for NGC 2244 – Park & Sung
2002, 1.2 for Orion Nebula Cluster (ONC) – Muench et al.
2002, 1.3 for W5 East – Chauhan et al. 2011, 1.3 ± 0.1
for Trumpler 14 and 16 – Hur et al. 2012, 1.1 ± 0.1 forNGC 6231 – Sung et al. 2013a). On the other hand, theslope ( = 1.7 ± 0.1) and shape of the IMF for NGC 2264
(Sung & Bessell 2010) are similar to those of IC 1848 in thesame mass range. In order to arrive at a firm conclusion on
the shape of the IMF down to subsolar mass regime, further
deep optical and/or NIR observation with extensive X-ray
observations are needed. The complete IMF of the cluster
will give an insight into the star formation processes in thelarge-scale star forming region.

15

5

ACCRETION RATES OF PMS STARS WITHA UV EXCESS

The mass accretion rate gives important information on thedisk evolution around PMS stars. Since the emergence of themagnetospheric accretion model for PMS stars, the physicalquantities associated with the accretion process have beenstudied by observation and modeling UV excess emission
or/and spectral lines such as H↵, Pa , Br , [O I] 6300 Å,etc, within such a paradigm. In this section we estimate themass accretion rate of those PMS stars with a UV excess
found from our photometric data using the relation between
the accretion luminosity and U -band brightness (Gullbringet al. 1998) and compare the results with that of other studies which used independent ways to estimate accretion luminosities for di↵erent star forming regions.
The first step is to find UV excess stars among PMS
members which exhibit either H↵ emission or a MIR excess.
Several PMS stars exhibit a UV excess in the [(U V )0 , (VI)0 ] diagram as shown in Figure 12. The majority of stars
with a UV excess are late-type members (V I > 0.7 mag).
These late-type members appear to be bluer in U V thannormal MS stars without excess emission at a given (V I)0 .

tributed by the National Optical Astronomy Observatories, which
be overestimated.
We haveconducted afor
deconvolution
is mayoperatedby the Associationof Universities
Research in As- the star formation history in W5 West, including IC 1848, stars within a given mass bin ( log m = 0.2). For the most
massive bin, the bin size ( log m) was 0.64 because of theof the brightnessof HD 17505
by assuming
the multitronomy
under operative
agreementwith thethat
National
Science the age distribution based on more complete, membership small number of massive stars. The star counts were then2 http://www.univie.ac.at/webda/
selectionis needed.
ple system consists of an O6.5 III, two O7.5V stars, and an
Foundation.
normalized by the logarithmic mass bins and the observed
O8.5 V (Hillwig et al. 2006). As a result, the HD 17505 sysKarr & Martin (2003) estimated a lower limit of 2.4 area. The IMF of IC 1848 is shown in Figure 11. To avoid theO6.5III
component
deconvolved from the others)
Myr for the age of IC 1848, assuming a sound speed for the binning e↵ect we shifted the mass bin by 0.1 and re-derivedc tem
2013(theRAS,
MNRAS
000, 1–??
should be dimmed by ⇠ 0.58 mag from its current position,
ionized gas, and an upper limit of 5 Myr, from the lifetime of the IMF (open circles), using the same procedure as above.
In the figure, the circle and dot represent the IMF ofand is then well consistent with the 5 Myr isochrone. AlO stars. Koenig & Allen (2011) inferred that the age of thethough we adopted a luminosity class III for the brightest
W5 region is younger than 5 Myr using the evolution model members with H↵ emission or YSOs from Koenig et al.
(2008b), while the dotted line corresponds to the IMF incomponent of HD 17505, it is worth noting that the genuineof Siess et al. (2000). Our estimates are in good agreement cluding all stars within the PMS locus. The shape of the IMF
luminosity may be very close to Ib or Iab supergiant becausewith those of the previous studies. On the other hand, Mo↵atthe deconvolved absolute magnitude is -6.29 mag (see Table(1972) claimed that the age of IC 1848 could not exceed 1 c 2013 RAS, MNRAS 000, 1–??
4. in Paper 0).
Myr given the existence of the binary or multiple system, HD
We have also estimated the age of IC 1848 from low17505, HD 17520, and BD+60 586. However, no supporting
Figure13.
luminosity
vs. stellar
luminosity(left) andaccretion
mass PMS
stars using
theAccretion
PMS evolution
models of Siessdiscussionof the luminosity
class massor the lifetimeof suchratestars vs. stellar mass (right). Open circles (red) denote
et al. (2000).
10 showsSquares,
the age distributionof PMS
was provided
in thethestudy,
and webelieve thatan agederivedof 6
ourFigureestimates.
triangles,
and crosses
represent
mass
accretion
rates
in di↵erent star forming regions by Hartmann
1 Myr
is unlikely.
stars. Theetmajorityof PMS stars
cover a spanetof 6al.
Myr(2011),withal. (1998),
Mendigutı́a
and
Natta
et al. (2006), respectively. Dotted lines represent the relations Lacc / Lstellar and2Ṁ / Mstellar
with an arbitrary constant.
c 2013 RAS, MNRAS 000, 1–??

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,
Σasat
0.39,
gindicate
cmthe,images
respectively,
while bands
for thinicemantleopacities
it is
0.58,
0.75,
(LD01). Dotted line shows contribution to thisfrom dust;stars dominateat=short0.44,
wavelengths.
Large 1.05
diamonds
convolution
We
follow
similarmethod
here
for
in
the
IRAC
at
24
and
70µm,
but
with
the
followinof these SEDs with the filter response function of corresponding instruments (panel (b)). The red/blue/green symbols show foreground6
Lim and Tan
intensities, relative to the LD01 IRAC 8µm value,measuredas
to thediscussed
three saturated coresby
C1/C4/C11,
respectively.
(c) Rightvalues
middle: Filterof Σ than
Note,
BT12,
higherwillofbeview
present
the map,
but cores
these previously
will be ten(1)We
searchfor70,“local
smaller
1! by 1Σ! sat
fieldsthatincontaindense
response of IRAC bands 1-4 (solid lines), MIPS
24µm
(dottedline) and PACS
100, & 160µmsaturation”(dashed lines). (d) Rightinbottom:Dust
opacities relative to effective opacity measured
in the IRAC
8µm band.
Black/blue bylines show
thin/thick ice mantle
models of OH94.
and
more
affectedsaturation,
leading
tounderestimationof thetrue Σ (2)values.
BT12.
This
helps
tolinesminimizetheeffects
foreground spatial
variation.
In addition to a standard
Solid lines show the initial, uncoagulated MRN
grain population.
Dottedand dashedindicate these models
after 10and 10 yr ofof
coagulation at n = 10 cm , respectively. The red/blue/green symbols (for C1/C4/C11 regions; slight wavelength offsets applied for
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
clarity) show κ /κ µm for the indicated Σ ranges.
We also analyze archival 70µm Herschel-PACS
images. uncertain.
The
first typeand(proposal
IDwill
KPGT-okrause-1)
were
and 24µm
maps,
that
such
hightheΣobserved
structuresofwill
be in regions
thathave been
identified
as being
satu
somewhatWe
gauge
liklihoodsaturation
also
byconsideringthemorphology(e.g.
cowith medium scanning speed and also have !6 resolution, but do not cover a very large area around the IRDC (∼ 16
wavelengths.
core/clump
we indeed
see this morphology:
thewavelengths.
70µm map peaks
at arouin extent). We estimate a 1σ noise levelof ∼ shorter
15
MJy/sr, i.e.
about
a factor of
two Inbetterthanachievedwith C1,
fast
oftwo
the
saturated
pixels,togetherwith
theiroverlapwith saturated
pixels at otherIn generalwscanning observations (Traficante et al. 2011). The second type (proposal ID KPOT-smolinar-1) were observed in theata2010)
positionclose
towiththe
C1atcore
centeridentified
by BT12
BTK.(i.e.
Thisgenerally
highest column
densityregionw
saturatedregiontobelargerwavelengths
with larger
dust and
opacities
increasing
at shorter
Galactic plane HiGAL survey (Molinari
et al.
in fastscanning
mode∼9 resolution.
These
data are usefulfor assessing intensity of the Galactic background emission, which needs to+
be interpolated from regions surroundingwiththesize
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 detec
the IRDC.
−2
Both the Herschel data sets are obtained already processed to level 2.5 in HIPE, so zero-level offsets need to be applied
structures
in
the
range
Σ
∼
0.5
−
1
g
cm
,
which
are
not
so
reliablyprobedbytheshorter
wavelength
m
can
also
be
compromisedbythepresenceofanearbysourcethatenhances
the
local
level
of
foregroundeto have measurements of absolute values of specific intensities. We adopted a model spectral energy distribution (SED)of the diffuse Galactic plane emissionGiven(Li &In
Draine2001C4[LD01])from
NIRwe
to FIR.
We fitthis24model
to theobserved
theregion
see
two
main
high
Σ
peaks
that
exhibitlocal
saturation,withthewestern
structuthelarger
beam
sizes
at
and
70µm,
we
adopt
a
less
stringest
spatialextentcriterion,
requirings
median intensities in a 2 ×2 field centered on the cloud, considering data at 8µm (Spitzer-IRAC), 24µm (Spitzer!!
MIPS), 60µm and 100µm (both IRAS)
and C4then predicted
the expect
intensities
inthe
Herschel-PACS
70µmbyband.
saturated
pixels
by
about
two
times
the
beam
FWHM,
i.e.
12
.
as
by
BT12
and
the
eastern
as
C13
BTK.
In
betweenaretwosourcesseen
mostclearlyat
24µm.
T
A single offset value was then applied to each Herschel dataset (908 MJy/sr and 241 MJy/sr for medium and fastscan data, respectively). We tested thisundoubtedly
methodoffset corrections
reported
Bernardet al.background(2010) at
l = 30
Ouragainst
methodfor estimatingtheintensity
same
adoptedandthatBT12,
i.e.istheaffectthe byMIREX
and
FIREX
methods isinthe
thisregion,so itbyis BT09
possiblethere
a hs
based on Planck data, finding agreement at the ! 10% level.
methodappliedoutsidethe WedefinedIRDC
ellipse from S06 (filter size set to 1/3 of the major axis, i.e.
thetwocores.
We find an astrometric difference offilter
abetween
few arcseconds
betweenthe Herscheland Spitzer maps.
corrected thisbythe average value of the mean positional offset of four point sources seen at 8, 24 and 70 µm, amounting to a 3.4
bythree-color
interpolation
insidethis
We inspectedbackgroundthree filamentary
control fieldsstructu
equal
In C11,
there
isis showna relatively
widespreadhigh theΣ region,
whichfluctuations
is part of ainhighly
translation at P.A. 62.7 . The resulting
image ofthe IRDC
in Fig. ellipse.
1.
While our focus is on FIR data, we
also comparetoSpitzer-IRAC
3.5, 4.5,
5.9, 8µm (GLIMPSE)
images, which
size
just
outsidetheellipse:after
foregroundsubtraction
theaverage
1σfluctuations(fromafittedtowards
C12
to
the
NE
(BTK).
The
70µm
map
also
revealsarelativelyhighΣregionextendingtoGathhave 2 resolution at 8µm. For these, we adopt 1σ noise levels of 0.3 MJy/sr, 0.3 MJy/sr, 0.7 MJy/sr, 0.6 MJy/sr,
respectively (Reach et al. 2006).
17.3%,
25.3%
and
42.3%
at
8,
24,
70µm,
respectively.
Intheoptically
thinlimit,
thesevaluesprovideanextension
is
seen
in
the
8
and
24µm
maps,
but
at
lower
Σ.
It
is
possiblethat
localized
foregroundvari
To compare multiwavelength extinction mapping pixel by pixel, we regrid the IRAC images to the 24µm MIPS framewith its 1.25 × 1.25 pixels, similar to
the 1.2
IRACwavelengths
pixels. However,for Herschel-PACS
data,
with
its thus
3.2
theuncertainty
in thederivedτ andΣ inany maps
given (seepixelalso
dueBTK,to backgroundshorterareaffectingthe
8 pixels,and
24µm
where thisfluctuations.
extension is more pro
we first carry out extinction mapping on the original pixel grid (especially searching for saturated pixels, see below),
8µm
map derived
with a finer decomposition of foreground variations).
before finally regridding to the MIPS frame
for multiwavelengthcomparison.
2.3. MIR to FIR Opacities
2.2. Radiative Transfer and Foreground Estimation
3.3. ofTheMIR
FIR Extinction Law and Evidence for Grain Growth
We adopt the 1D radiative transfer model of BT12, which requires knowing the intensity
radiation
directedto
towardBT09,
wetargetadopt
the diffuse Galactic background from the model of LD01. We w
the observer at the location just behind, Following
I , and just in front,
I , of the
IRDC. aThespectrum
infrared emissionof
fromthe IRDC is assumed to be negligible (which
imposes
an upper limiton the dustof ∼ 15in
K; non-saturated(Tan et al.
relativevaluesof temperatureΣ
derivedregions at different
wavelengths
yield
information
abo
thatThea by-productof ourextinction
mapping
is a measurementof this spectrumof theforegroundemissi2013b)), so thatofWethethenextinction
law.different
This lawis (1)
expected(Table
to varygrains undergo
via icecoagulationan
Idust= econsiderI
severaldust
models1), 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

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

10.0

H2O

1.0

1

H 2O

κλ/κ8µm

CH3OH

H2O

H2O
CH3OH

21

25

H2O

1.0

0 yr
105 yr107 yr

0.1

1

0.1

70

10

λ (µm)

78

100

◦

5

6

H

λ

7

−3

8

!!

!

!!

◦

◦

◦

!!

◦

!!

!!

!!

!!

!!

ν,0

ν,1

ν,1

−τν

ν,0

Fig. 2.— Top row: Saturated regions in IRDC C1 viewed in Galactic coordinates at 8, 24 & 70µm (left, middle, right panels; intensity
scales in MJy/sr). We have searched inside the displayed 1! × 1! square for saturated intensities in each pixel, defined to be within 2/4/8σ(white/blue/black, respectively) of the minimum pixel value. Small/medium/large circles show 8/24/70µm saturation, respectively, which
are displayed together on the left panel and separately for 24 and 70µm on the middle and right panels. The circles in the lower-rightindicate the beam. Middle row: Σ maps derived from the 8, 24, 70µm images (left, middle, right panels; intensity scale in g cm−2 ) assuming
the moderately coagulated thick ice mantle dust model of OH94. Bottom row: Ratios of Σ maps (or equivalently deviation of relative
opacity from the OH94 thick ice mantle dust model) for Σ24µm /Σ8µm (left), Σ70µm /Σ24µm (middle) and Σ70µm /Σ8µm (right).

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

36 560 40.000

570 00.000

Dec (J2000)

R CrA
IRS5N

SMM 1C
IRS7A

20.000

IRS5A

IRS7B
40.000

SMM 1A
T CrA

580 00.000

5000 AU57.00s51.00s54.00s19h 01m 48.00s

RA (J2000)

A&A proofs: manuscript no. herschel_rcra_arxiv

Fig. 1. Schematic overview of the three PACS footprints, centred at SMM 1C (red), IRS7B (green), and IRS5A (blue), overlaid on a Spitzer
4.5 µm image. The spaxel positions as recorded by PACS on the sky are at the center of each spaxel, and each spaxel measures 900. 4 ⇥ 900. 4 in thisrepresentation.
Table 1. Observational parameters: Positions give the centre coordinates of the central spaxel. All footprints are 5 ⇥ 5 spaxel grids with approximately 900. 25 separation between the spaxel centres.

RegionRA(J2000.0)

Dec(J2000.0)

P.A.a
[]

Observing IDs

Date of observation

IRS7A
IRS7B
IRS5A

19:01:55.3
19:01:56.4
19:01:48.1

36:57:17.0
36:57:28.3
36:57:22.7

0.22
2.47
2.40

1342206990, 1342206989
1342207807, 1342207808
1342207806, 1342207805

2010-10-23
2010-11-02
2010-11-01

Notes. (a) Position angle with respect to an equatorial grid coordinate system.

Fluxes at these wavelengths are thus less reliable than at otherwavelengths (Herczeg et al. 2012).
In addition to the Herschel data, we also used the ISIS spectrograph on the William Herschel Telescope (WHT) on 6 August2012 to obtain a low-resolution (R ⇠ 1 000) optical spectrumof R CrA covering 3 000–10 000 Å. The spectrum was flux calibrated against the spectrophotometric standard LTT 7987. The
analysis of this spectrum will be discussed in Sect. 4.1.1.
2.2. First look

In this section, we provide a first look at the PACS continuum
and spectral maps before they are treated by our deconvolution
algorithm.
2.2.1. Continuum maps

The total emission continuum maps of IRS7 and IRS5 at three
representative wavelengths can be found in Figs. 2–3, and continuum maps at 20 µm steps can be found in Figs. A.1–A.2.
These continuum maps show extended emission in the order of
30–6000 in size. The larger sizes are found in the longer wave-

lengths (which could be suggested to be attributed only to thelarger beam size at these wavelengths, but the shapes of the emission suggest di↵erently; see also Sect. 3.2). The highest continuum flux densities in the strongest illuminated spaxels are foundaround 85 µm in IRS7 and 105 µm in IRS5, but the spatiallyintegrated flux density is strongest around 120 µm in IRS7 and155 µm in IRS5 (since the longer wavelength data have more extended PSFs, which spreads the signal over a larger solid angle).
It is, however, difficult to draw conclusions about the stages of
evolution of the individual sources from the total emission, since
it is a combination of several compact sources and any extended
emission. For such a study, deconvolution of the emission is necessary.
The continuum emission in the IRS7 region is very extended
in comparison to the line emission (see Sect. 2.2.2), and has atleast some signal across most of the two PACS fields that cover
the region (see Figs. 2–3). Without any deconvolution, it is impossible to distinguish the point sources detected in mid-IR and
mm data (IRS7A, IRS7B, SMM 1C, and R CrA).
IRS5 consists of two separate sources, of which the northern source IRS5N is detected in SMA continuum, whereas thesouthern IRS5A is detected in Spitzer continuum and line emisArticle number, page 3 of 37

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