星形成ゼミ
2020年1⽉31⽇
古屋 玲#323 16-25

18.

power-law index q and the measured continuum intensity
distribution derived from the visibilities V(s) as a function of the
NOEMAのラージプログラムCOREの最初の論⽂
uv distance s. Following the theoretical work of Adams (1991)
and Looney et al. (2003), the density power-law index p canbe
化学的進化の解析を物理進化の議論に適⽤するテストケース
derivedassuming
spherical
symmetry:
An observational and modeling case study of the AFGL 2591 VLA 3 hot core

Chemical complexity in high-mass star formation:V(s) / s p+q13 / s↵ .

ほかに17天体を観測

arXiv:1910.05081v1 [astro-ph.SR] 11 Oct 2019

C. Gieser1, 2 , D. Semenov1, 3 , H. Beuther1 , A. Ahmadi1, 2 , J.C. Mottram , Th. Henning1 , M. Beltran4 , L.T. Maud5 , (2)F. Bosco1, 2 , S. Leurini6 , T. Peters7 , P. Klaassen8 , R. Kuiper9 , S. Feng10, 11 , J.S. Urquhart12 , L. Moscadelli4 ,T. Csengeri13 , S. Lumsden14 , J.M. Winters15 , S. Suri1 , Q. Zhang16 , R. Pudritz17 , A. Palau18 , K.M. Menten13 ,19 , Á.
19 35865corr
proofs:manuscript no.
R. Galvan-Madrid18 , F. Wyrowski13 , P. SchilkeA&A
Sánchez-Monge, H. Linz1 , K.G. Johnston14 , I. Jiménez-Serra20 , A&A proofs: manuscript no. 35865corr
S. Longmore21 , and T. Möller19
Table 1. Data products of AFGL 2591 used in this study.
Totheovercome(Affiliations can be found afterreferences)this problem, the intensity profiles can be fit in the
Fourier-plane (e.g., Beuther et al. 2007b).Beam
TheSizedensity
profilePosition
Angle
Frequency
Spectral Resolution
rms Noise
Received / Accepted⇥ ✓min temperature
PA⌫v⌫
AFGL領域内に
can be determined using the previously ✓majderived(00 ⇥00 )()(GHz)(km s 1 ) (MHz)power-law index q and the measured continuum
intensity
同様な天体が他に3つCORE data
1400AU相当＠3.3
kpc
the visibilities
V(s)as a function 66
of theABSTRACTdistribution derived fromNOEMA
continuum
0.46⇥0.36
218.873
0.59 mJy beam 1
いずれもVLAで検出
uv distance s. FollowingNOEMA
the theoretical
work
of
Adams
(1991)
line data
0.47⇥0.36
66
217.224 220.773
3.0
2.2
1.2 mJy beam 1 channel 1
Aims. In order to understand the observed molecular diversity in high-massstar-forming
regions,
wehaveto
determinetheunderlying
physical(なのにUCHIIでなく、and Looney et al. (2003),
the density
index p can
Mergedline data power-law
0.47⇥0.36
67 be
217.224 220.773
3.0
2.2
1.5 mJy beam 1 channel 1
and chemical structure of those regions at high angular resolution and over a rangeofevolutionary
stages.
IRAM
30 m LI
11.8⇥11.8
0217.224 221.274
0.27
0.2
2.2 Jy beam 1 channel 1
derived assuming sphericalsymmetry:Hot
Methods. We present a detailed observational and modeling studyof theCoreとclaimしているが)hot core VLA 3 in the high-mass star-forming
region
AFGL
2591,
which
IRAM 30 m UI
11.2⇥11.2
0229.172 233.222
0.25
0.19
1.5 Jy beam 1 channel 1
is a target region of the NOrthern Extended Millimeter Array (NOEMA)
large
program
CORE.
Using
NOEMA
observations
at
1.37
mm
with
p+q
3
↵
Archivaldata
等⾼線は1.4mm連続波
A&A proofs: manuscript no. 35865corr
V(s) / s/s .
(2)an angular resolution of ⇠0.00 42 (1 400 au at 3.33 kpc), we derived the physical and chemical structure of SMA
the source.
modeledthe observed 28
SiO lineWe
data
6.8⇥5.6
216.976 217.214
1.1
0.8
30 mJy beam 1 channel 1
温度構造が両輝線で異なるのは
molecular abundances with the chemical evolution code MUSCLE (MUlti Stage ChemicaL codE).
Results. With the kinetic temperature tracers CH3 CN and
H2 CO we observe a temperature distribution with a power-law index of q = 0.41 ± 0.08.
H2COが⾼温に温度計としての感度ないから
3.3.1. Temperature structure
Using the visibilities of the continuum emission we derive a density structure with a power-law index of p = 1.7 ± 0.1. The hot core spectra reveal(rot.
diagramの提⽰なし;
high molecular abundances and a rich diversity in complex molecules. The majority of the molecules have an asymmetric spatial distribution around
The molecules H2 CO and CH3 CN can probe the kinetic temperFig. <5.1Vector-averaged
complexwevisibilities
V explainof the 1.37
continthe forming protostar(s), which indicates a complex physical structureon scales
400 au. Using MUSCLE,
are able tothe mmobserved
ature when multiple transitions at di↵erent energy levels are obCH313CNを持ちいてtauをuumchemical
as a function
uv distance.
A In
power-law
fit the(redobservational
solid line) wasFig. 4.abundance(a) H2 CO andCH
derivedXCLASS
molecularof 10
outof temperature
14 modeled maps
species
at an withestimated
hot coreage ofof⇠21
100 years.
contrast toserved (e.g., Zhang et al. 1998; Rodón et al. 2012). The CH3 CN
3 CN
含めて評価しているかの記述なし)
in orderderivethe power-law
index ↵. The gray data
at a thresholdof 10modeling. The black
contours
showdensity
the 1.37mm
continuum
emission is strong around hot cores and can trace dense highanalysis,
our chemical
predicts
a lower
power-lawindex of performed
p < 1.4. Reasons
forto
this
discrepancy
are discussed.

https://arxiv.org/pdf/1910.05081.pdf

1temperature gas (Fuente et al. 2014), whereas H2 CO is used topoints were
excludedfit. Theerror pictureon the offit the(one
standard
emission withlevels athigh5, 10,
20, 40,and 80 observations( = 0.59 mJy). chemical
Conclusions.
Combiningspatialresolution
withbeamdetailed
modeling
allowsfrom
us tothederivea concise
physicaltrace the temperature of the colder extended envelope (Mangum
the red fordashedlines. CORE sample, and thengray dash-dottedindicates
the 2591
outer hotradiusof Thethe radial
tem- aredeviation)
and Thechemicalstructure of circle
the famous
AFGL
core.
next steps
to conductisashown
similarbyanalysis
the whole& Wootten 1993). The upper state energies Eu /kB of the obperaturefit (bottom
panel).
The
beam size(0.00 47⇥high-mass
0.00 36) is starshownin
use thisanalysis
to constrainthechemicaldiversityinformationtoamuch
greater
depth.
C. Gieser et al.: Chemical complexity in high-mass star formation:
served transitions of H2 CO and CH3 CN are listed in Table A.1.
the lower right corner of the H2 CO temperature map. (b) Radial temThe vector-averaged
complex visibilities of the 1.37 mm
Keyperaturewords. profiles
ISM: individual
objects:
AFGL
2591
–
Astrochemistry–
ISM:
molecules–
Stars:
massive
To derive the physical parameters of the molecular gas suchobtainedthein thetop panel.
The data
are continuum emission, computed using the miriad software
C. from
Gieser etal.:data
Chemicalcomplexity
in high-massstar formation:
as column density N and rotation temperature T rot , we fit thebinned in steps of half a beam size (⇠700 au) for H2 CO (green data
(Sault et al. 1995), are plotted against the uv distance in Fig. 5.
spectral lines with XCLASS6 (eXtended Casa Line Analysis Softpoints) and CH3 CN (blue data points). A power-law fit was performed
ware Suite, Möller et al. 2017). The XCLASS software models
visibilities were shifted such分⼦流のblue/redが⾒ている空間スケールで異なる;multipleのためと解釈that the position of the 1.37 mm
on spatial scales in the range 700 4 000 au to determine the temper-theThe
1. Introduction
innermost
severalthousand
au
and
rich
molecular
spectraatmolecular lines by solving the 1D radiative transfer equation
ature power-law index q. The gray dash-dotted line indicates the outer continuum peak is at the phase center at ↵(J2000) = 20:29:24.88
assuming local thermal equilibrium (LTE) conditions and an(sub)mmwavelengths
areobserved(e.g.,
Belloche
et
al.
2013;
and (J2000) = +40:11:19.47. With a power-law fit (minimumof the ofradialfit. Atstars
radiiissmaller
than field
700 auofthe
Theradius
formationthetemperaturemost massive
an active
re-tem-Sánchez-Monge
isothermal source. In dense sources, such as hot cores, the LTE
2
et
al.
2017;
Beltrán
etal.
2018).
At
later
stages,method), an index of ↵ = 0.87 ± 0.05 is obtained on a physperature
profile cannot
be determined
the finite
beam size(grayassumption is a reasonable approximation for rotational lines; insearch(for reviews,
see, e.g.,
Beuther etdueal.to2007a;
Bonnell
2007;
Hii⇠3
(UC
Hii)
regions
form wherewith
theaprotostars
ical
scale
from
000
67
000
au.
Visibilities
uv
distance
shaded area). The transparent data points were excluded from the fit-ultra-compact
Sect. 3.3.3 we estimate an average H2 density of ⇠107 cm 3 inZinnecker & Yorke 2007; Smith et al. 2009; Tan et al. 2014; ionize their surrounding envelope, which is observed in strong
ting. The best fit is shown by the red solid line. The error of the fit (one s > 200 k are excluded from the fit due to an increase in thethe central part within a radius of ⇠2 500 au. The critical density
Schilke
2015;
Motte iset shown
al. 2018).
high-massemissionat cm withwavelengths(e.g.,
Garay&
Lizanois < 107 cm 3 for most of the observed molecules (CH3 OH bestandard
deviation)by theSince
red dashedlines. star-forming free-freenoise level.
Combinedtheprevious
determinationofthe
2. AFGL 2591 continuum emission at 1.37 mm. The green conregions (HMSFRs) are less abundant than their low-mass coun- 1999; Palau et al. 2007; QinFig.
et3.3.1),
al.
Klaassen
et al.
2018).
tours
show2008;
themm continuum
emissionwithlevels at 3 (dashed), ing the only exception; see Table A.1). Therefore, the population
temperature index q (Sect.
the1.37
density
index
p can
be
cal1
terparts and are typically located at large distances (> 1 kpc), it
3,
5,
10,
20,
40,
and
80
(
=
0.59
mJy
beam
). The beam size of the molecules can be described by a Boltzmann distribution
culated according to Eq. (0.
2 00as00
46⇥0.
36)
is shown wavelengthin Vtheoflower
corner.
green star indi- with a single characteristic temperature T ex = T rot = T kin . In
is challenging
to study
Fig. 5. Vector-averaged
complexvisibilities
the left
1.37
mm ThecontinThese authors
find individuala value of collapsingq = 0.40 ±objects.
0.01 and a tempera- Numerous
millimeter andsubmillimeter
obsercates the position of the 1.37 mm continuum peak. The white rectangle XCLASS, the line profiles are assumed to be Gaussian, but optias qa +functionof uvshows
distance.
A power-law
fitaverage
(red
solid
line)
ture
of
250
± 20
at 13 CN
000 temperature
au for2591derived(theirTables1 andvations
Fig.The
4. (a)
Handmaps
withtakes
XCLASS
majorityofK CH
high-mass
starAFGL
formation(HMSF)
revealrichchemistry
nitrogen2 CO
puum= ↵towards
3.hot cores(3)waswas
thearegion
withinwhich an inspectrumobtained in the cal depth e↵ects are included. Each molecule can be describedin orderto derive
theadata
power-law↵.(Sect.
Thevangray data
at a 2).
thresholdof 10infrared. regions
The black
1.37
mm
continuum
a recentstudy
by Barrshow
etclouds.
al.the(2018),
hot onabundant
place
inInthedensestof contours
molecularBasedob- and performed
spectralline
for line identification
3.2).
oxygen-bearing
molecules(forreview,seeindex
Herbst&by multiple emission and absorption components. The transition
1
pointswerea density
excluded
the fit.
Theof
error
the±fit0.1.
(onestandard
emission
at 5,a10,
20, 40, andof80
(± 59
=can
0.59
).Dishoeck
CS is with
detectedwithtemperature
714
K fromabeam
regionWeobtainindex
p =on(Herpin
1.7
This
properties (e.g., Einstein coefficients, state degeneracies, partiservational
andlevels
theoreticalconsiderations,
HMSF
bemJy
divided2009)
as well power-law
asfromsulfur-bearing
species
et al.
tion functions) are taken from an embedded SQlite3 databasedeviation)
is
shown
by
the
red
dashedlines.
Thegray
dash-dotted
circleindicatestheouterradiusoftheradialtem130 au.
This is consistent
with theettemperature
profile from2009).
resultwith
Palau etorganical. (2014),
where(COMs)
these authors
find
into<several
evolutionary
stages (Beutheral. 2007a; Zinnecker
Inagrees
addition,complexmolecules
are very00
00containing entries from CDMS and JPL using the Virtual Atomic
peraturefit(bottom
Theofbeamsize(0. 47begins⇥ a0.radius
36)
isof
shown
our study,
where
this temperatureis reachedat
56 au.inabundant
p = 1.80
0.03 for
AFGL
Observations
ofdefinition
HMSFRsfind&Yorke2007):
the panel).
formation
massivestars
in infraredin±thoseenvironments.
theby
the 2591.
James Following
Clerk
Maxwell
Telescope(JCMT)and the IRAM 30 m and Molecular Data Centre (VAMDC, Endres et al. 2016).
the lower
right
cornerharboringof the H2 CO
temperature
map. (b)Radial tem-Herbst
similarvaluesfor the density
power-law
index
van
der
Tak
telescope(Bisschopet al. (e.g.,
2007;
der
Wiel
et mm
al. 2011); anddark
clouds(IRDCs)potentiallyshort-lived
high-massThe
vector-averaged
complexvisibilities
of
the
1.37
&
van
Dishoeck(2009),
a COM
consistsofAFGL
sixvanor
more
package ofprovides
new functions
for CASA
Fig.
7.
Outflow
structure
of
2591
at
di↵erent
spatial(Ceccascales traced byThe
IRAMXCLASS
30 m observations
CO (left panels),
SMA observationsof SiO
perature
profilesobtained
from
the
data
in
the
top
panel.
The
data
are
480
1900
GHz
with
the
Herschelspaceobservatory
et
al.
2000;
Beuther
et
al.
2002;
Mueller
et
al.
2002;
Hatchell&starless
cores
and
lowto
intermediate-mass
protostars(e.g.,
Pilatoms.
Molecules
can
form
in
the
interstellar
medium(ISM),
continuum
emission,computedusing
themiriadsoftwaretomodelinterferometricandsingle-dish
data.
By
using
the
(centralpanels),
and
NOEMA
+
IRAM
30
m
observationsofSO(rightpanels).
Upperpanels:The
background
gray
scale
shows
the integrated
3.3.2.
Density
structure
relli
et
al.
2010;
van
der
Wiel
et
al.
2013;
Kaźmierczak-Barthelbinnedin 2006;
steps ofhalf a beam
size2006;(⇠700Sanhuezaau) for H2etCOal.(green
van
der
2003;
Beuther
et intensity
al. 2007b;
et
al.
2009)
aand redoftheZhang
centralline
emission.
TheandblueFig.
contours showmyXCLASSFit
the blue- and redshifted
linetheemission.
The fitbeam
is givenin the lowerlai
et al.
Rathborne
et al.
2012;datawhere
function,user canthesizemodel
parameters
the Tak
gas
and
dust
shielded
fromstrongstellarradia(Saultet al.
1995),
areet plottedagainsttheuv
distancein5.
al. 2014).
Bisschop
et al. (2007)
andstar
Ceccarellial. (2010)
re-1.37 mm continuum peak. The orange box indicates the change of field of view.
points)andpeakdata points).
A power-law
wasperformed
leftSect.
cornerin each panel.
The greenmarks theetpositionof the3 CN
discussion
is given
5.1.
(source size ✓source , rotation temperature T rot , column density N,
Zhang
al.CH
2015)
anddensity
proceeds
formhigh-mass
protostelTheetflux(blueof tothe
1.37
mm fitcontinuum
istiondetailed(at avisibilities
visualextinction
of
Ain
⇡
200
mag
in
the
dense
parts
of
The
were
shifted
such
that
the
position
of
the
1.37
mm
vLowerported
that
AFGL
2591
has
a
“line-poor”
spectrumcompared
panels: Spectra of the outflow tracers extracted from thetopositionof the 1.37
mm velocity
continuumo↵set
peak from(greenthestar).
The systemic
velocityon 3.spatial
scales
in
the
700
au
to brightness
determinethetemper1 range
linewidthv, andsystemic
velocityvo↵ ) of00 4 000
00
00
1
Fig.
Average
spectrum(spatiallyaveraged
over
1.
6⇥1.
6,
indicated
by
the
white
rectangle
in
Fig.
2)
of
the
AFGL
2591
hot
core
labeled
with
00
00
55.47
mJy
beam
correspondingtoatemperalar
objects(HMPOs)
with
M
>
8
M
showinggas
accretionothermassivebutbywithbeam
sizes >line.
10 Thethebluecompact
HMSFRs;Tan 2591
etpeakal.hot2014).
Asummary
ofisthedetected? 1. 6⇥1. 6, indicated by the white rectangle1 in Fig.
5.5 kmcenters protostars,
indicatedthemolecules
green
dashedand red dashed lines show the integration intervals for the blue- and redshifted line
3. Average spectrum (spatially averaged over2)continuumof the AFGL
corelabeled
with
is
at
the
phase
at
↵(J2000)=
20:29:24.88
toobservational
data
by
using
di↵erent
optimization
algorithms.
ature
power-law
index
q.
The
gray
dash-dottedline
indicatestheoutermolecular emission lines with S /N > 5 (indicated by the red horizontal line, = 0.34 mJy beam in1 the average spectrum). Unidentified lines are
Theranging
integrationinterval
ofbe
theresolved
central line
emission isandset between the two ranges.
emissionthecorecannotsufficiently
ecularlines
S /N >2),
5 (indicatedby the red
horizontal
line, Motte= 0.34 mJy
average
spectrum).
Unidentifiedlines emission.
are aroundture
8.5with
Ktable(Fig.
whereas
gas
has
>
200
andemission
molecular
outflows(e.g.,
ettransitions
al. temperatures
2002;
et beamal. KininthetheISM
is given
McGuirefromdiatomiclabeled
UL.ofAofdetailedof the properties
ofBeuther
the themolecular
given700
in Appendix
and(J2000)= in+40:11:19.47.
Withahotpower-law
fit(minimum
In contrast, the myXCLASSMapFit function can be used to fit
3.3.3.
H
densityand (2018)gas
mass
2 column
theradialradii transitions
smaller isthan
au theA.temeledradius
UL. A detailed
table of temperature
the properties offit.
the Atmolecular
is givenin Appendix
A.
is heavily
beam
diluted. Our 1.37 mm spectral line data at sub-the H2 CO abundance inmodel does not include
Table 4. Model input parameters.
sity of H13
Best-fit Model Parameters
2 CO (Fig. 8) shblueshifted NE outflow
Inner radius rin
700 auParameter
Value or Range
Outer radius rout
29 000 au
The 2 and ⌧HMC dis
周囲のmassive
YSOs起因Radiation field:
Temperature power-law index q
0.41
space is shown in Fig. 11
1
UVはoutﬂow
cavityTemperature
T
at
r
253.7
K
CR ionization rate ⇣CR5( 17) s
0in
fit model lies in a paramから⼊射したとしてモデル
Density
power-law
index
p
1.0
1
⌧HMC ⇡ 3 000 6 000 y
X-ray ionization rate ⇣X
0s
3
Densitynatr
1.7(7)cmと述べているが、
0in
rin or rout as within our mExtinction at the cloud edge Av
100magMass M
487 M
ﬁg.7からの知⾒を
small 2 (Fig. 11).
C. Gieser et al.: Chemical complexity in high-mass star formation:
HMC
stage
⌧
4
642
years
UV photodesorption yield
1( 05)
HMC
どう考慮したかの明⽰なし
C. Gieser
et al.:
Chemicalcomplexity in high-mass star 10/14
formation:
Within the best-fit m
Wellfittedmolecules
Grain properties:2
variation of the modeled
0.5295Grain radius rg
0.1 µmat formation:
the time of the best-fi
The evolutionary past of the physical structure and chemical
C. Gieser et al.: Chemical complexity in high-mass starb
3
DCN, and H2 CO abund
Notes.
= a⇥10.
As a(b)initialabundances,
Dust density ⇢d
3 g cm composition of AFGL 2591 is not known.
過去
sulfur-bearing species S
we use the best-fit IRDC abundances from the study by Gerner
Gas-to-dust mass ratio ⌘150
et al.(化学進化の計算開始時)(2015). For 19 well-studied IRDCs,The
columndensitiesof derived from the observations given in the center up to ⇠4 00
columndensities
Surface di↵usivity EDi↵ /EBind
0.4
18 molecules including deuterated species
wereareobtainedusing
as input
listed inTable 2. The H2 column density is calcu- wards. The nitrogen-bea
Mantle composition
Mg x Fe xIRAM
SiO(olivine)
1
430 m observations at 1 and 3 mm.
Thefromcolumnlated
thedensities
1.37 mmofcontinuum emission at the peak position, are enhanced in the ce
HMC期にある時間は解いた結果
the molecules were averaged and modeled
to create inabundances
Pre-HMC stages
as analyzed
Sect 3.3.3. The CO column density is inferred as well. A similar profi
これはof a template IRDC. These authors derive an IRDC chemical
CH3 OCHO, but the abu
IRDC stage a
IRDC期以降現在までの時間であってage of ⌧IRDC = 16 500 years using MUSCLE(Table
A.4
and
A.8
in
9
In the current MUSCLE version, uncertainties of the observed column contrast, the abundance
その時間にはIRDC期終了後から
Inner radius rin
12 700 au
Gerner et al. 2015).
densities cannot yet be included as an additional constraint.
region > 10 000 au (>hyper
compact
HIIが形成された時間もOuter radius rout
0.5 pcFinally, the physicalstructureof the AFGL
2591 hot core has
missing the short-spacings
Temperature T 0 at rin
11.3 K to be defined. Our NOEMA data are含んでいるはず(しかし明⽰的に区別していない)
in the continuum yielding a missing flux of 84 %, so we can only
Temperature power-law index q
0.0 (isothermal)
determinea lower limit for the outer radius of 7 000 au (Fig. 2).
3
Density n0 at rin
1.4(5) cm
From the observed radial temperature profile shown in Fig. 4(b),we can adopt an upper limit for the inner radius of 700 au which
Density power-law index p
1.5
is half これはgivena beam size. The
density index p is varied from 1.0 to(⼿で⾊々与えてみたのか？)⌧IRDC
16 500 years
2.5. The temperature profile is fixed to the radial profile derived
HMC stage:
現在 presented in Sect. 3.3.1 (Eq. 1) and thusfrom the observations
Inner radius rin
50 700foraueach physical model the temperature T 0 at the inner radiusr
Inner radius step rin
50 au is calculated according to T 0 = 255 K ⇥ ( 691inau ) 0.41 . In total, Fig. 10. Comparison of the observed (red) and modeled (blue) column
densities.
Outer radius rout
7 000 we
31 run000 theau chemical code ALCHEMIC on 2912 physical models(#rin ⇥ #rout ⇥ #p = 14 ⇥ 13 ⇥ 16). For each model, it takes
H2以外は解いた結果
Outer radius step rout
2 000 au⇠75 s using 32 cores to compute the chemical evolution withfrom the observed C18 O abundance assuming an isotopic ratio
Temperature power-law index q
0.41 ALCHEMIC up to 100 000 years on 40 radial grid points.
of 16 O/18 O = 500 (Wilson & Rood 1994). For molecules whose✓◆ qlines from the vibrational ground state and from vibrationallyTemperature T 0 at rin b
T 0 = 255 K ⇥ 691rinau
excited states were observed, an average value for the column
4.2. Best-fit modelDensity power-law index p
1.0 2.5
density is given as input.
2
Using
a
analysis,theobservedcolumndensities
NarecomobsDensity power-law index step p 0.1
The best-fit parameters from the model with the lowest 2
paredto the computed model column densities N
for each are summarized in Table 5. The best-fit model has
5
総量はgiven
a hot core
⌧HMC
10 10physicalyears model and time step. Vasyunin et al. (2004)modstudied the lifetime ⌧index of[Eq.6]
HMC ⇡ 4 600 years, with a density power-law⌧HMC
100 logarithmic
steps of the abundances of chemical models due toerror propagationp = 1.0. A comparison between the modeled and observed colthe large uncertainties and approximations of rate coefficients. umn densities is shown in Fig. 10. For 10 out of 14 molecules(b)
Notes. a(b) = a⇥10b . (a) Taken from Gerner et al. (2015).
Constraint
These authors
estimated an error of 0.5 1.5 orders of mag- the modeled and observed column densities agree within the unfrom the observations described in Sect. 3.3.1.
nitude depending on the size of the molecule. Wakelam et al. certainties. For SO, DCN, and HNCO the model column den(2005) found that for older hot cores with ⌧ > 104 years, the un- sity is too high compared to the observations with Nmodel = 40,
Nobs
certainties in the chemical model can be large for
speciesFig.some
12. Radial
abundance profilesof the 18,
MUSCLE
best-fit model, relative
to thefor
radialdensity
profilein the bottom
right panel.
48, andrespectively,
while
HH22COthemodeledcolumnTable 5. Best-fit model of the AFGL 2591 VLA 3 hot core.

e
d

at
al
s.
y

al
d
al
d
y
nuC
e
m
e
les
2

e-

e
a
or

6)

i-

as:

me that during the disk-accretion
phase, the stel-outflow types take advan- 2.1. Rotational evolution& Ferreira 2013). Since these two last
cretioncurtainsandtheexerted
star directly
mass
Theslow
torque(red)directlyby theaccretes
MEs
onto
the andstar anismentum.
The
radiusR
corresponds
totheregionaroundinitially
rotating
stars.
Foreacht
https://arxiv.org/pdf/1910.03995.pdf
tage of magnetic
fieldlinesenergythat still connect
the star with the
Rotational
evolutionof solar-type
protostars
duringthe star-diskswheredriving
power
from
a
fractionoftheThe
temporale↵ect
evolutionof the surface
stellar andangu
momentumfrom the
disk.
lated
a di↵erential
rotation
betweenthe starthe dipolar magnetosphere truncatesaccretion
disk, they canthe
exchangemass, energy,
andto
angularmomentumaccretion
rates
areinvestigatedṀ=
10interactionphase
acc evolution ogoverned by the angular momentum
前主系列星(PMS)システムにおける星の対流層の進化と円盤の進化がwith
both2008):
theontostar and thethe disk.
impactof the accretion
columnsstellar
MEs.
The
torquethat
exertedby stellar
magnetic
field
lines
coup
We
assumetheisconvective
mass
accretion
rate
evolvesaccordcanbe approximatedby (Bessolaz
et al.
envelope:(lowerpanel)
M
/yr
with
Qacc
=Garatti
1 (dotted)a(Fig. 2. Evolution of Ṁ? as function
caseof Gallet
fast (blue)andカップリングして進化する様⼦を初めて考察(これまでは以下のJconvはTheinworks& Bouvier(2013,
2015)
and
Amard
F. Gallet of, C.time
Zanni
, and
L.by
Amard
to
a
region
of
size
R
of
the
MEs
can
be
expressed
as(Armita
simpledecayfunction
basedonCarattiO
etal.
0slow4 (red)1relation
ablishes
betweentheaccretionandthe
1/7 rotating
12 a
˙ + İ ⌦ =を計算+
etal. (2016),
aimed
at comparing
theoreticalmodels andobser- J̇
initiallystars.
Foreachrotators,
twoinitialmass
定数だった);
本研究では
=I ⌦
+
,
1.2
B
R
BBB Univ.
C&
Clarke
1996;
Matt&
Pudritz
2005b)
?
i.e.
Ṁ
/
t
)
C
dip
10
7
acc stars, make thevations
the rotational
CCCAlpes,
Grenoble
IPAG, 38000 Grenoble,
rates
are, CNRS,
investigatedṀbetween
10 of(upper
panel)evolution
and10of solar-mass
BB@accretionacc =France
werandconsequently
themassaccre(5)INAF - Osservatorio
AstrofisicodiTorino,StradaOsservatorio
20,10025,
Pino
Torinese,
Italywhere J
= I ⌦ is the angular momentum of
hypothesisthat during
the accreting TTauri phase the 2stellar6 ro-! 1.2
2 AM /yr with Qacc = 1 (dotted)(lowerpanel)
and10%(solid).
✓◆1.2is its yrsに渡って数⽇程度
GM
UniversityofExeter,
DepartmentofPhysics&
Astronomy,
StokerRoad,
Devon,
Exeter,
EX4
4QL,
UK
[結果]
太陽質量程度のPMSの表⾯部の回転を数M
B
R
? Ṁ
acc
tive
envelope,I
moment of inertia, and ⌦tationis
kept
constant
thankstoa‘disk-locking’mechanism,
but
?
t
t
dip
disk
disk
d outflow
rate.
We
take
into
account
twovalues=fdecay(t)q =R1, velocity. The term
1 few
, is kGのdipolar磁場がthe sum of all of the exte
they did not take into account the actual
detailsof any magneReceived -; accepted MEのタイムスケールで維持させようとすると、a4
3
Kratioisadimensionlessparameter.
Following
Long
et
al.
t the stellar surface, whiletただし、Ṁ
Rt is to acting onis the torqutopshericstar-diskinteraction
model. The
aim of thistinit
work
必要
1%
and
10%
.
wind /Ṁacc = Qacc =add self-consistentelements to a physical[課題]この予測磁場強度は計測されることもあるが、例外的であるため,model of the star-disk with the core-envelope angular momentum exchang
and Zanni & Ferreira (2009), we assume
Kt = 0.5.
6 19
where
qtdisk=represents
B/B/&[⌦the⌦ofmomentum/⌦MEs(LivioPringleangularvariationof&theconvectiveeand to the evolutionary
trackspresented
inthe
Galletk has dissipated(t > ⌧disk),interaction
the stellar
wind canzstartingMEs? ]ourwhere
isage
simulation
(i.e.
10
yr)
Gallet
et
al.:
On
th
init
The early pre-main sequence phase during
which solar-mass stars are still likely surrounded by an accretion
モデルの改良(星表⾯に多重極磁場構造など)にとりくむ
to theMatt
developmentof the radiative
core(see e.g.
A
Bouvier(2013,
2015)andandcontracting,Amard etal.do(2016)
into&
order
to deter-1996;
a puzzling stage of their rotational evolution. While solar-massstars
are accretingthey
not seemspinup
Armitage
Clarke&Pudritz
2005a)
takes
into
the
disk
lifetime,
and
t
the
age
considered.
The
mass
accr
wer
from accretion
anymorewind
mineandunder thewhich stellar
physical conditions
the stellar rotation period Gallet & Bouvier 2015; Amard et al. 2016). We note
substantially.
Magnetosphericejectionscountthedi↵erential
rotation
thefofdecay
star
andthe MEs.
It is usually assumed that the magnetospheric star-disk
interaction
tends
tomaintainthestellarrotation
periodconstant(‘diskisthus
definedṀrelated(t)between(t). velocity
isto=theṀvariation
angulardue toT
can be kept approximately constantrate
while
the
protostars
are stillas
acc
acc,initΓK=0.21,
K
は1程度の正定数;Klocking’),but
thishypothesis
has
never
been
thoroughly
verified.
Our
aim
is
to
investigate
theimpactofthestar-diskinteractionΓ
2.3.3.
Stellarwinds
acc
themechanism
onlyon thespin-down
mechanismleft.
As Toinmodel
Galof accretionthe moment
oftorqueinertia of the convective
envelope
accretingcontracting.
thestellarspin
evolution
we thefactorcanbeapproximatedas: <1はsubstellar spin evolution during the accretingpre-main andsequencephases.
Weparametrizeas
常に回転を増す働き
acc
Keplerian回転Γ
Es (ZanniWe&devised
Ferreira
2013)
resultinteraction
into Γ
account
includethespin-down
by aandstellarwind, employa model forthe torques
acting onfrom
the stellarthe
envelope
based on torquestudies ofofexerted
stellar winds,
we developedaacc, maxthe stellar contraction and the growthΓ this!
2013,2015),
duringphase
themass-lossisp
3/2
Γ tot,stellar
new prescription for the
coupling founded on numericalof star-diskproposedinteraction andejections.
ME we= magnetospheree.g.,Γ star-disk
whereashave neglected
the changeof theing
thesimulations
parametrization
Matt ettorqueal. (2012).
Never- mass,
10% ejection,The
exertedbythemagnetizedwindsonmagnetosphere
andtoanaccretion
disk
that
producesthe by magnetosphericRΓthe long-term evolutiont
We then used this torque model
followofthe stellar
rotation.
Γ
=
K
Ṁ
GM
R
,
ME
Γ
acc
acc tomomentum
acc of, solar-type? t protostars
inertia
due toRomanova+09
mass accretionfrom the surrounding dipresent
new prescription
fortheangular
Galletet al.:On K
therotrotational
evolution
q surface/
Strong dipolar magnetic
fieldthecomponents
up totheless,a fewfieldkG weare required
toa
extract
enoughangularmomentum
so as
keep
Γ
thetheand
BOREAS
routine(Cranmer
&
Saar
2011)
can
be
expressed
as
(see
Weber
& Davis1967
onreconnectionofmagneticlines
connecttot,
1%
associatedstellar outflows.
In the following
a star and ana surrounding
accretionR
disk
Γ of solar-type stars approximatelyexchange
surface rotation rate
constant for between
a few Myr. Furthermore
efficient enough spin-down
torqueR with= coronation
radius (これより外側
co based loss Γwind, 10%
Galletetand
al.:Oncorresponding
the rotationalof rate,solar-type
be provided
either one of the following:a stellarwindwitha mass outflow
rate⇡ 10%evolutionoftheis
accretion
proportional
to the torque’massaccretion‘internalas=+rate Ṁİ acc⌦ and.
on resulting
thenumerical
magnetohydrodynamic(MHD)
simulationsof protostars
starcanwithitsbysurrounding
disk.
The
inflationat to whichllar
angularvelocity,
luminosity,
radiusas
2.3.3.
Stellarwinds
では星表⾯と円盤は磁気的に直接作⽤せor a lighter stellar wind combined with a disk thatis truncatedaroundthe corotation radius entering a propeller regime.
Γexpresses
2 rot < 1
Zanni
& Ferreira(2013)
that show the
impactofMEs
onKepleriantheTocomputethe torque
,in
wethe
follow
MacGreg
withK=Ṁ0.7.
ArstelK
the fact that
thetruncatio
MEs
ro
wind,
1%momentum
rot
specific
angular=⌦and accretion
poweredstellaris
winds
play andynamic
important role in thedisk
spinwindevolutionof solar-typetudes of theMagnetospheric
dipolar ejections
magnetic
field
lines
very
ず) [カッコ]内 < 0の時、回転抑制[R
wind
?
A
The
torqueexertedbythemagnetizedstellar
winds
on
the
stellarlar angularmomentumevolution.
did
notconsiderGhoshrate.
& ner(1991). It implies
that
both
theMEscore andthe eos along
prevent
sharp
transitionatt=⌧stars. However, kGa
dipolarmagnetic fields
are neither uncommon
or ubiquitous.
Besides,
it is during
unclearWehowmassivestellarwindsacan
disk
ata
sub-Keplerian
Assuming
R
/
R
,
the
torquet
2.3.3.
Stellarwinds(truncation
radius)
>
K
R
]
逆に
R
gion.
As
it
is
betterspecified
in
the
next
section,the
propo
surfacecanbeexpressed
as
(see
Weber
&
Davis
1967)
with
outflows
can exchangemass
andangube powered
while numericalmodelsthatof the propeller
regimedisplay a strong
variability
that
has noobservational
confirmation.
BetterLamb-likestar-diskmagnetic
coupling
since
it has beentested to in solid body rotation with di↵erent angular velociti
常に回転を抑制observational regime,
statistics and more realistic
models
could
contributetoahelp
lessenour calculations’requirements.
bewrittenas:が⼩さいとき,raction
we
use
the
wind
mass
loss
rate
provideratherinefficientspin-down
torque(seeṀ
e.g.
Matt
et al.
tity
J of accountangularmomentum
thenality
constant
Kacc
takes
into
the星表⾯の回転は増す
factis that
inrAthe
where
is
the
wind
mass
loss
rate
andexchanged
is trtb
mentum
with
both
the
disk
and
the
star.
Because
of
the
=
Ṁ
r
⌦
(10)
wind
The torqueexertedbythemagnetizedstellar
winds
on
the
stellar
2
3
Stars: solar-type – Stars: magnetic field – Stars:
evolution–
Stars:
rotation–Stars:
winds,
outflows
2010), and we preferred taking into account scenarios2 that 6rely core and the envelope
a time-scale ⌧ (hereaf!3/2 overrotation
tion
region,
the
disk
is
not
in
Keplerianbecauseooc valuefield
lines
reconnection,theseoutflows
detach
from
B
R
6
7
value
of
the
Alfvénradius.
Weuse
theexpression
betweentheaccretion-poweredandthesurface can be expressed as(seeṀon
Weber
Davis
1967)
? 6
7
outflows.
We didnotconsidereither
X-wind(Shu etdip
al. 1994)
envelopeR
coupling
time-scale). The
quantity
J is tht
whereis the&wind
mass
loss
rate andr isthe average
67664Krot
775 ,thatfraction=
K(Ferreiraet al.
scenarios
that
are
still
based
angularmomentumthe core andoutflows
thathave
extracteda
relevant
ofthe
theenvelope
disk
gnetosphereandcontinue
theirasmagneticME
MEvaluepropagationof ortheReX-wind
Alfvén radius.
We usethe2000)
expression
obtainedbyMatt
etthreeal.percentiles(2012)
3
2..whichFigure2displaystheevolution
of
the
mass
loss
Introduction
highlight
theapparent
steadyevolutionoftheR
Rtmech- change instantaneouslycoon(2012)
phenomenologicalmodels and we
preferred
to include
in ordertoto thehave region
the same an=disconnectedṀwind rA ⌦? from
(10)
Matt
et al.
windarethecentralpartofthemomentum.
TheradiusR
correspondsar(90th,
median,and star25th) of the larrotational
period distributionst2
3
anismsforwhich
self-consistentMHD
numerical
solutions
are
ity.
We
also
assume,asin
Allain(1998),
that⌧
isClassical T Tauri stars (CTTS) are magnetically active pre-main
2
3
m
stars
from 1 Myr
to 10toMyr.
The fact that the two extreme
nequenceof andtime
fromthe diskaccretionphase
the666currentlyofavailable.
777up
低降着率で回転を抑制、⾼降着率で回転を促す
tem
are
in
betweenopen
magnetic
surfaces
that
where
K
=
0.21.
It
is
interesting
to
note
that
this torque
sp
the
star
where
the
dipolarmagnetospheretruncates
the accr
stars surroundedbyan accretion thedisk (Collier
Cameron6
7
agiven
model.
ME
percentiles
remainalmostconstant
in time confirms6that the evoB R
7
2
2
6
7
where
Ṁ
is
the
wind
mass
loss
rate
and
r
is
the
average
67
windA
666 be on
777
777 is& Campbell 1993; Edwards et al. 1993; Collierr Cameron
BRp R>?K2/3=figureK 6666 et al. q
R not, random(11)
lution theof the disk
rotationbut strongly
depends
the
hored
into
the
star
and
those
anchored
into
(see
me
of
our
models.
The
clearly
shows
the
7
disk.
It
can
approximatedby(Bessolaz
et aal.lower2008):
down
the
star
only
for
R
while
trunca67995; Edwards
al. 1998).We
Their4Article
mag-the
666 in a fast q t
777 Rforvalueet al.
of1994;the Hartmann
Alfvén etradius.
use
expression
obtainedrot coinitialconditions,thatis, whetherby
the starisinitially
number,page
2Rof 512
r
=
K
,
Ṁ
K
v
+
⌦
A
1
?
etic
field
(up to a &few FerreirakG, see Johns-Krullet al. 2009; Gregory or slow rotating regime. Since this constant rotational
664MEs775up.
omZanni2013).
phase
is 1theradius,thespinstellarsurfaceMattetal.
(2012)
0
1/7
en
the
APSW
mass
loss
and
the
BOREAS’s
Fig.
3.
External
torquesasfunction
ofthemassaccretionrate(loweraxis)
2
2
2t al. 2012) has a strong impact on the dynamics andrate
struc4122
p comparable to the disk lifetime, it suggests a BlinkB
betweenṀ+ ⌦simulations
R?the
CCCK2 vof
wind
opening
the
field
lines
limitsthesizeescthe? R?
ureof thestellarof
environmentand can lead radiusto thetruncation
of
where
v 3=
2GM
/R
is
the
escapevelocity.
We
assumetheBdip
2 magnetospheric
We
used
the
resultspresented in Za
andcorresponding
truncation(upper
axis)
for
given
set
of
stellarmagnetic
star-diskinteraction
and
thisobserved
behaviour.
FiB
m
B
C
he
disk
and
accretionofmaterialalong
funnel
flows
down
to
the
R
=
K
,
ttellar
=
⌧
.
B
C
6
7
valuesm=0.2177andK=
0.0506,
provided
by
Matt
et
al.
t
t @ rotational evonally,of conthe physicalof thisdisk
777 2.5
2 the
ea
over
the
star
and
magneticallyparametersat 6661 Myr(Mwinds=2 R
1.0
M disk, R? are=
R , Pregardless
Bp = &
1originFerreira(2013)2 toA calibrate the constants appearing in E
?B
? = 7 days,
surface,which
the launchof stellar
along
the
opened
mag-

16.

1

2

3

conv

1
2

conv

?

conv

3

conv

conv

?

ext

c e

rad

?

conv

?

ext

c e

ABSTRACT

Context.
Aims.

ME

acc

Methods.

rot

rotacc, maxtot, 10%ME

Results.

tot, 1%

wind, 10%CO

wind, 1%

int

c e

rad

c e

Conclusions.

2wind A

wind

?

A

2
p

1

wind

esc

t

co

c e

wind

A

?

t

2/3
rot

Key words.

conv

2
?

2 2
2 esc
?

c e

m

2 2
? ?

?

?

2

GM Ṁpwhich typically does not extend beyond the corotation 4 where
and 9. vThese
small fractionthe= results
2GMexplore/R isonly
the aescapevelocity.ofWe
where
Kspaceis =aofdimensionlessparameter.
Followingrameterthe star-disk
interaction
problemandLongevenvaluesm0.2177andK=
0.0506,
providedby!
(2005)
andZanni&
Ferreira(2009),
weassumeK=
0.5.
orque summary
limitedsamplesuggests K
thatisthe
K in, KTable, andK
values
(2012).
The
constant
given
2.
GM
Fig. 3.
External
torques
as parametersfunction ofmass.
(6) not constant
and
are
a
function
of
the
of
the
mo
Inand
agreement
with
the
APSW
scenario(Matt⌦corresponding
radiusout some important properties of the total extersuch as the massaccretion rate.truncation
However, thesimple(upper
formulaa
2005a), we assume
that during
the disk-accretionpha? (2012). The
K1theseis givenin Table constraints
2.
lution,observationalsuggest that during
666& Pudritz 2005a,p2008;
777 constant? thisacc
etic
fieldrThe
lines
(Matt
Matt
et
al.
2012;
=
K
R
,
(11)
kG).
accretion
torque(blackdottedline),
themaximum
spinq
7
A
16
?
acc
early-PMS
phase,
alarge
fractionoftheangularmomentumofIn
agreement
with
the
APSW
scenario(Matt&
Pudritz
775
esc
?
?
6
Cranmer & Saar 2011;6
Matt et al. 2015; 2Réville
et
al.
2015;
John4
2
2
2 =+
thethatstarneedsthe
tothe
beMEs
removed.
Fig. 2. Evolution of Ṁ? as function of time in case of fast (blue) andup
torqueallowed(K1)(grey
dottedline),
torque2005a),
weassumeduringdisk-accretion
phase,
thestelṀ
K
v
⌦
R
acc
acc
wind
tone et al. 2015; See et al.
2018; Finley& Matt ?
2018),
t
? and
2 esc
slow (red) initially rotating2stars. For each rotators, two initial mass
lar
wind
takes
its
driving
power
from
a
fractionoftheenergy(dot-dashedline),
thestellarwind
torque
for
Q
=
1%
(grey
Moreover, acc
the magnetospheric star-disk interaction scenario
he ejectionof material due to the magnetic star-disk interacME
wind
accretion rates are investigated Ṁacc = 10 10 (upper panel)t and 10 7
dissipatedbytheimpactoftheaccretioncolumnsonto
thestellar
p
is
still
the
main
paradigm
tointerprettheangularmomentumondashed(Shu et 1/3
al.
1994;
Ferreira
et
al.
2000;
Romanova
et
al.
2009;
acc
rot 10% (solid).
MEline)vand =Qacc 2GM= 10%/R(blackdashedline),
andaofrelation
the
total
external(lower panel) M /yr1with Qacc = 1 (dotted)
and
where
is theescapevelocity.
We
assumeThisestablishes
betweenthetheevolution
CTTS(see
Bouvier
etaccretional. 2014, andfor athereview). The
Zanni &?Ferreira2013).
esc?
?surface.
torque= acc + suggest
foretQ= power
1%
solid
line)
and the angularext observations
ME +
wind
acc1993;
first(orange
attempts
to modelthe star-disk
momentum ex2 values
winddriving
andconsequently
betweenmass accreAdditionally,(see
Edwardsal.
m=0.2177andK=
0.0506,
provided
by
Matt
et
al.
2
? et=al.10%
in CTTS(see takee.g. Königl
1991) wereon the GhoshBouvier
1993;(redRebullsolid
et al. line)
2004; Irwin&tion
Bouvier2009)
Qacc
are plotted.
and thewindchange
outflow
rate. We
into accounttwobasedvalues
2.3.2.
Magnetosphericejections(2012). The constant K is given in Table
2. (1979) model, which was& Lamboriginally
3 developed for X-

16.

A&A proofs: manuscript no. 35432corr

A&A proofs: manuscript no. 35432corr

⻘は90%パーセンタイル
NGC6530

NGC2362

the temporal evolution of the mean magnetic field intensity as
a function of stellar parameters. In Cranmer & Saar (2011) themean magnetic field BCR is given bys8⇡⇢⇤ kB Te↵
0.55
BCR = 1.13
(2)
⇥
⇤ ,
µmH
1 + (x/0.16)2.3 1.22

where x = Ro/Ro with Ro = Prot /⌧conv the Rossby number withProt the rotation period, ⌧conv the convective turnover timescale(see Wright et al. 2011; Oglethorpe & Garaud 2013, for moredetails about the convective turnover timescale), Ro = 1.96, ⇢⇤the photospheric mass density, kB the Boltzmann’s constant, Te↵
the e↵ective temperature, µ the mean atomic weight, and mH themass of a hydrogen atom.
The BOREAS routine reproduces reasonably well the meanmagnetic field and the mass loss rate of the present day Sun(hB i ⇡ 2-7 G and Ṁ ⇡ 2 10 14 M .yr 1 ). We note that forless massive stars, this formalism produces magnetic fields withsmaller intensity than what is measured using spectropolarimetry (Morin et al. 2008, 2010). When extrapolated to convective
TTauri stars, this model does not provide magnetic fields consistent with the total mean stellar field intensity, but the strength
produced by the BOREAS model is compatible with the dipolar components of CTTS (a few hundred Gauss2 ). We assume
that our torque models only depend on the dipolar field intensity
since, due to its slower radial decrease, the dipolar component ismainly responsible for the large-scale magnetic interactions of
the star with its surroundings.
1

http://lasp.colorado.edu/~cranmer/Data/Mdot2011/
It is important to note that in our models, the Bdip value corresponds
to the dipolar field intensity at the magnetic equator, which is the value
entering the definition of the magnetic dipole moment. Spectropolarimetric observations that perform spherical harmonics decomposition of
the surface stellar field, usually provide the values of the components attheir magnetic pole (see e.g. Donati et al. 2011), which in the case of adipole is twice its equatorial strength.
2

⾚は25%パーセンタイル

ONC

Cep OB3bh Perseus

Fig. 4. Angular velocity evolution of stellar convective envelope as function of time for fast (up) and slow (down) rotator models in case where

Angularvelocityof stellar
convective
envelope
as thefunctionof APSW
time forfast (up)models in case whereṀacc,init= 10 9 1M evolution.yr 1 . The blackand thegrey solid lines
represent
MEs andprocesses
withand
Qacc slow= 1%(down)
and Qacc rotator= 10%, respectively,
9= and10 theM.yr . field
The isblack
the greyroutine
solid lines
represent
and
APSW
withlines
Qaccrepresent= 1% andQacc =including
10%, respectively,magnetic
from andthe BOREAS
from Cranmer& the
SaarMEs(2011).
The
red andprocesses
the blue dotted
the models
magnetic
is fromthe BOREAS
from
Cranmer
Saar (2011).
and the blue
dottedmagnetic
lines represent
modelsthe MEs fieldand APSW
processeswith Qaccroutine= 1% andQacc= 10%,&
respectively,
but Thewith red
the numerically
imposedfield Bmodthe. Thefree including
contraction(i.e.
withoutanyexternal
interaction)
is
shown
by
the
magentalong-dashedline.
It
represents
theincreaseoftherotationrate
duringa . The frees and APSW processes with Qacc = 1% and Qacc = 10%, respectively, but with the numerically imposed magnetic field Bmod
freecontraction
The angularvelocity is isscaledto thevelocitylong-dashedof the present line.
Sun. The
blue and thetilted square
androtation
associated
ion
(i.e.
withoutphase.
any externalinteraction)
shown
byangular
the magenta
It represents
theredincreaseof therate during aerror bars represent, respectively, the 90th percentile and the 25th percentile of the rotational distributions of solar-type stars in star-forming regions
traction
phase. The angular velocity is scaled to the angular velocity of the present Sun. The blue and the red tilted square and associatedand young open clusters obtained with a rejection sampling method described in Gallet & Bouvier (2013).
rs represent, respectively, the 90th percentile and the 25th percentile of the rotational distributions of solar-type stars in star-forming regions

Dark Collapse or Supermassive Star: a fundamental limit imposed by hydrostatic equilibrium
L. Haemmerlé1, G. Meynet1 , L. Mayer2 , R. S. Klessen3 , T. E. Woods4 and A. Heger51Département d’Astronomie, Université de Genève, chemin des Maillettes 51, CH-1290 Versoix, Switzerland; 2 Center
for Theoretical Astrophysics and Cosmology, Institute for Computational Science, University of Zurich, Winterthurerstrasse 190, CH-8057 Zurich, Switzerland; 3 Universität Heidelberg, Zentrum für Astronomie, Institut für Theoretische Astrophysik, Albert-Ueberle-Str. 2, D-69120 Heidelberg, Germany; 4 National Research Council, Herzberg
Astronomy & Astrophysics, 5071 West Saanich Road, Victoria, BC V9E 2E7, Canada; 5 School of Physics andAstronomy, Monash University, VIC 3800, Australia
E-mail contact: lionel.haemmerle at unige.ch
Major mergers of gas-rich galaxies provide promising conditions for the formation of supermassive black holes (SMBHs;
5
4
5
−1
>
∼10 M⊙ ) by direct collapse, since they can trigger mass inflows as high as 10 − 10 M⊙ yr on sub-parsec scales.
However, the channel of SMBH formation in this case, either dark collapse (direct collapse without prior stellar phase)4
or supermassive star (SMS; >10M⊙ ), remains unknown. Here, we derive a criterion on mass and accretion rate for∼
SMS formation by testing the consistency of hydrostatic equilibrium in case of rapid accretion. We compute hydrostatic
models of SMSs accreting at 1 – 1000 M⊙ yr−1 , and estimate the departures from equilibrium a posteriori by taking−1
into account the finite speed of sound. We find that stars accreting above the atomic cooling limit (>∼10 M⊙ yr )can maintain hydrostatic equilibrium only once they are supermassive. In this case, they evolve adiabatically witha hylotropic structure, i.e. entropy is locally conserved and scales with the square-root of the mass coordinate. Our
results imply that stars can become supermassive by accretion only at the rates of atomically cooled haloes (∼ 0.1 − 10M⊙ yr−1 ). Once they are supermassive, larger rates are possible. Dark collapse occurs if the accretion rate exceeds
the atomic cooling limit before a SMS already formed.
Accepted by A&A
https://arxiv.org/pdf/1910.04776.pdf

Characterization of mid-infrared polarization due to scattering in protoplanetary disks
Stefan Heese1 and Sebastian Wolf1
1

論⽂のタイトルが変更されているが、次ページは上の論⽂
University of Kiel,
Institute of Theoretical Physics and Astrophysics, Leibnizstrasse 15, 24118 Kiel, Germany

E-mail contact: sheese at astrophysik.uni-kiel.dethermal relaxation and accretion. The first process, appearing
mathematically in the energy equation, relies on the variationof entropy ṡ. The second one, appearing in the boundary conditions, relies on the variation of mass Ṁ. The efficiency of thesetwo processes can be approximately estimated from their corresponding global timescales, the Kelvin-Helmholtz (KH) time
and the accretion time:

hift zBC⇠ 7V9E(Mortlockal. 2011; Wu merger (Bañados et al. 2019) and has an inferred mass similar to
2E7, etCanada2018;
Wang et al. 2018). Assuming that of the massive disc obtained in the simulations.
6
https://arxiv.org/pdf/1910.04776.pdf
Schoolof forPhysics
andthroughAstronomy,Letter
Monash
VIC 3800, Australia
to theUniversity,
Editor
rmed
at
z
⇠
30,
example7
Tsung-Dao
Lee Institute,200240,China
mpliesthe accumulationof mass ShanghaiatThe
propertiesof SMSs accreting at the rates of atomicallyyr 1 in compact objects for a bil- cooled haloes (0.1 – 10 M yr 1 ) have been studied both anaMaximallyaccretingstars: a fundamental limit
Received; acceptedundin atomicallycooledprimordial supermassive
lytically and numerically in the last decade (Begelman 2010;s have been destroyed imposedby an exter- byhydrostatic
equilibrium
Hosokawa
et al. 2012; Schleicher
et al. 2013; Hosokawa et al.
t al. 1997; Omukai 2001; Bromm & 2013; Sakurai et al. 2015; Umeda
ABSTRACTet al. 2016; Woods et al.
1 , G. Meynet1 , L. Mayer2 , R. S. Klessen3,4 , T. E. Woods5 , A. Heger6,7
2008; Latif et L.
al.Haemmerlé
2013; Regan
et al. 2017; Haemmerlé et al. 2018b,a; Haemmerlé & Meynet 2019).
Theygalaxies
are foundprovide
to evolvepromising
as red supergiant
protostars
along
the
Context.
Major
mergersof gas-richfor theformationof supermassive black holes (SMBHs;
1
Départementd’Astronomie,
Universitéde Genève, chemindes Maillettes 51, CH-1290
Versoix,conditions
Switzerland
native
route
has
been
found
in
the
Hayashilimit(Hosokawa
et
al.
2012,
2013),
beforecollapsing
for Theoretical
Astrophysics
and Cosmology,
Computational
Science,inflows
University&2 Center
105 M) by direct
collapsebecauseInstitutetheyforcan
trigger mass
asZurich,
highWinterthurerstrasse
as 104 105 M yr 1 on sub-parsec scales. However, the5 ofs (Mayer
et al. Zurich,
2010,Switzerland
2015; Mayer et al.:due tothe GR instability
at several
10 M (Umeda
etfundamental
al. 2016;
190, CH-8057
HaemmerléMaximally
accretingsupermassive
stars: awithout
limit imposedhydrostatic equilibrium
channelof
SMBH
formation
in
this
case,
eitherdark
collapse(direct
collapseprior stellar
phase) orby
supermassive
star (SMS;
3
Universität
Heidelberg,
ZentrumfürAstronomie,
Institutfür
Theoretische
Astrophysik,
Albert-Ueberle-Str.
2,
D-69120
Heidelberg,
rio has 4two advantages compared to Woods et al. 2017; Haemmerlé et al. 2018b). The case of galaxy
Germany
10formation
M ), remains
st,&4star
does notunknown.
need to be mergers provides conditions for which rates higher than those
Universität Heidelberg, Interdisziplinäres Zentrum für Wissenschaftliches Rechnen, Im Neuenheimer Feld 205, D-69120
Table
1.
Massesat astheendlimittherun (set
byhaloesnumerical
instabilHere,weobtained,
investigate
thein accretion
rate
up toarewhichstarsHere
canwemaintainequilibrium.
h Aims.
largerinflows
arehighofofatomicallycooledpossible.
address hydrostatic
theHeidelberg,
Germany5512forNational
Research
Councilof Canada,
Herzberg
Astronomy
Astrophysics
Research
Centre,5071
Victoria,
questionof&whether
or accretion
not
SMSs
can
form
bySaanichaccretion-parsec
scales.
Thecomputegas accumulates
Methods.
Wehydrostatic
modelsofSMSs
accretingat
1 –West1000
MRoad,
yr at,ratesand estimate the departures from equilibrium a
ity)
in
10
M
the
modelswithvarious
rates
and
9
1
BC
V9E
2E7,
Canadadiscthat reaches
⇠ 10
Mnew
;account
what
&the
100
M markedyr. We
investigate
the conditions in which hydro6 metallicities.
posteriori
by takingintofinitespeed
ofinsound.
The
modelsareboldface.

19.

needed to restor
core cannot incor
Figure 1 show
is obtained using

M
R = 260 R ⇥M(1)by hydrostatic equilibrium
Haemmerlé et al.: Maximally accreting supermassive stars: a fundamental limit imposed

GM⌧KH =We find that ,stars accreting above the atomic
本論⽂で論じる点
Results.
cooling limit (& 10 M yr ) can only maintain hydrostatic equilibrium once
RL
1 M yr InTable
10 case,M
yrtheyatevolve100
M adiabatically
Mnumericalyrthey are supermassive.
this
with
a hylotropic
that is, entropy is locally conserved and
1
1. Masses
the endofyrthe run 1000
(set byinstabil-structure,
For the sound-cr
scalesZwith
of inthe5.43= 0 the square
2.29 rootity)
4.25 with various
5.25
10mass
M coordinate.
for the models
accretion rates andmarkedsupermassiveinboldface. by accretion at the rates of atomicallycases
Conclusions.
Our3.02
results metallicities.
imply that The
starsnewcanmodels
only arebecomecooled haloes
(⇠
Z=Z M
8.57
of
ideal
ga
0.1
larger rates
are possible.
⌧accr10 M= yr ).,Once they are supermassive,(2) mean molecular
1 M yr
10 M yr
100 M yr
1000 M yr
1e5
Msunで逆転
rs: a fundamental Ṁ
limit imposed byZhydrostatic
equilibrium5.43
=0
2.29
5.25
Table
2. Stellar
mass
at
the
beginningof冷却優勢H-burning.4.25
Keywords.
Keywords
should
be
given
加熱優勢
Z=Z
3.02
8.57
evant for the env
1 M yr Table
10 M
yr mass100
M beginning
yr
1000
M yr
2. Stellar
at theof H-burning.
where
M
is
the
stellar
mass,
R
and
L
are the photospheric ra- consider only th
Z=05000 M
20 000 M
50 000 M
100 000 M
ntroduction
happens
next however remains unknown. The disc might either
図2の星の位置
1respectively,
M yr
10 M yr
100
M 000
yr
1000theM yr gravitational conZ = and
Z
2000
M
40
Mis
dius
luminosity,
and
G
most
of &theradiu
directly to a black hole (dark collapse,
Mayer
BonoliZ=05000 M
20 000 M
50 000 M collapse
100 000 M
formationofthe first⌧supermassive
black
holesevolution(SMBHs) 2019),
Z=Z
2000
M , the40 000M first form a supermassiveorstar (SMS),
which
then colstant.
When
⌧
is
essentiallygoverned
longer
than
a ye
accr
KH
ne ofTable
the mainproblems
understandingfor- mass
3. Mass
fractioninof ourthe convective
coreofat thea stellaroflapses to form a SMBH once accretion stops or general relativisby
thermalcontraction
and
the
star
is
expected
Table
3.
Mass
fractionoftheconvective
core
at
a
stellarmassof to beon ofcosmic(e.g. Rees 1978, 1984; Volonteri & tic
We
note
the
e↵ec
200
000 Mstructures.
(GR)
instability
setsthermallyin (e.g. Woods et al.
2019).
Interestingly,
000 MThe. mosthappens
Introduction
next however
remains unknown. The disc might either
elman
2010; Woods
et aal.2002019).
extreme
SMBHs
distant
observed
quasarseems to mass
be hostedrange,
by a galaxy
relaxed
after
KH
time,
as
nuclearreactionsareignited.
Whencollapse
directly
toablack
holethe(darkmostcollapse,
Mayer &
Bonoliins
9
Mredshiftyr
M7 (Mortlock
yr
100
M
yr
1000
M
yr
eM
formationof theM
first1at
supermassiveblack(SMBHs)z10⇠holes
et
al.
2011;
Wu
2019),
or
first
form
a
supermassive
star(SMS),
whichthencol• ⇠ 10
1 M yr
10 M yr
100 M yr merger
1000 M (Bañadosyr
et al. 2019) and has an inferred mass similar to
one of the
main
problems20%
in our understandingof the for- lapses
to
form
a
SMBH
once
accretion
stops
orgeneralrelativisZ
=
0
5%
0.5%
0.4%
Z=
0Wang
20%
5% Assuming
0.5%⌧ofaccr⌧⌧(e.g.
, 2018;accretion
the evolution,
leading
to the. 2015;Bañados
et al.
et dominates
al. tic
2018).
of magnitude. A
KH
ation
cosmicstructures
Rees
1978,
1984;
Volonteri&that 0.4%
ofmassivedisc obtainedin the simulations.
(GR) instability sets in (e.g. Woodsetal.the2019).
Interestingly,
Z
=
Z
35%
1%
Z
=
Z
35%
1%
gelmanof
2010;
Woods
et al. 2019).
The mostat
extremeseed
theseobjects
formedz ⇠ SMBHs
30, forthe
examplemost distantthroughobserved quasar seems to be hosted by a galaxyofenvelope,theentropycannotve swellingM ⇠ 10 M at redshiftz the⇠ 7 (Mortlock
et al. 2011; Wu because
regimeof pure amerger(Bañadosetal. 2019)
and has an inferred
mass similar tobe radiatedctcollapse,
this
value
impliestheaccumulation
of
mass
at
The
properties
ofSMSs
accreting
at
the
rates
of
atomicallyal. 2015; Bañados et al. 2018; Wang et al. 2018). Assuming that of the massive disc obtained in thesimulations. Fig. 2. Internal profiles of the model at Ṁ = 1000 M yr ,1 of the model at Ṁ = 1000 M3
eage
seed
of
these objectsz ⇠ 30,
for1
examplethrough
rates
ofefficiently
1 formed– 10atMyrnot
in compact
objectsforbilaway
enoughtorestorethermalequilibrium
againstisincluded,sinceaccreting
SMSsa arefoundcooled
to be slowro- 2. (0.1
Fig.
Internal
profileshaloes–10Myr)10havestudiedboth anarates
around
10
Z
=
Z
and
M
=
8
⇥
M . beenUpperpanel:
Luminosity
proectFig.
collapse,
this
value
impliessince
the ofaccumulationof massat
The
properties
ofSMSs
accreting
at
the
rates
of
atomically1.
GlobaltimescalesEq.
(1–2–3),computed
with
Eq.
(4).
5
is
not
included,accreting
SMSs
are
found
to
be
slow
rotators
due
to
the
⌦
-limit,which
implies
negligible
impactofyears.
areinfoundatomicallyprimordial
file(solidblue)
andluminosity(dotted
Lower
= Z
=in
8 the
⇥Eddington
10lastMdecade. Upperpanel:white).
Luminosity
numerically(Begelman
2010;
erage
ratesSuchof 1 – rates
10 M yrcompactin
objects
for a bil- cooled
cooled haloes(0.1 – 10 M yr lytically) have beenZandstudiedbothandana- Mtheperturbativee↵ectof
accretion(Hosokawa&
Omukai2009).
The
sound-crossing
time
is
shown
for
ideal
gas
and
radiation
rotationon
their
structure(Haemmerlé
etal.
2018a).
The
use
glo
tatorsduetothe⌦ -limit,
which
implies
negligible
impactn years.
Such rates
found
in atomicallycooledprimordial
lyticallyby
and annumericallyin thelastofdecadefile(Begelman
2010;
panel:
Entropy
profile(solidblue)
and hylotropic
fit swhite).
/of
M
es,
where
H2aremolecules
have
been
destroyed
exter(solid
blue)
and
Eddington
luminosity(dottedL
Hosokawa
et
al.
2012;
Schleicher
et
al.
2013;
Hosokawa
et
al.
For
zero
metallicity,models
with
Ṁ

10
M
yr
are
alloes,where H The
molecules
have been
destroyedby anforexterpressure.
accretion
time
is shown
the indicatedrates.
Hosokawa
etal. 2012; Schleicher et al. 2013; Hosokawa
et al.white).
(dottedrotationon
their
structure(Haemmerlé
etal.
2018a).
UV
field(e.g.
(e.g.
Haimanal. if
1997;
Omukai
2001;
Bromm
& is
l UV field
Haimanet al. 1997;etOmukai
2001;
Bromm
In
principle,
theaccretion
time
short
enoughcompared2013;
et etal.
Umeda
al.conclusive
2016;
Woods ethowe
2013;
Sakuraiet (2018b).
al. 2015;
Umeda
etofal.Sakurai
2016;
Woods
al. 2015;
readydescribed
in&Haemmerléetal.
Modelsrapidly
panel:
Entropy
profile(solidetblue)and hylotropic
fital.s /School of Physics and Astronomy, Monash University, VIC 3800, Australia
Tsung-Dao Lee Institute, Shanghai 200240, China
1
1
Received ; accepted
7

5

1

1

1

ABSTRACT

Context. Major mergers1 of gas-rich galaxies provide promising conditions for the formation of supermassive black holes (SMBHs;

1
1& 105 M ) by direct collapse because they can trigger mass inflows
as high as 1041 105 M yr 1 on1sub-parsec scales. However,thechannel of SMBH formation in this case, either dark collapse (direct collapse without prior stellar phase) or supermassive star (SMS;& 104 M ), remains unknown.
Aims. Here, we investigate the limit in accretion rate up to which stars can maintain hydrostatic equilibrium.
Methods. We compute hydrostatic models of SMSs accreting at 1 – 1000 M yr 1 , and estimate the departures from equilibrium a11
posteriori by taking into account the1finite speed of sound.1
1
Results. We find that stars accreting above the atomic cooling limit (& 10 M yr ) can only maintain hydrostatic equilibrium once
they are supermassive. In this case, they evolve adiabatically with a hylotropic structure, that is, entropy is locally conserved and1111
scales with the square root of the mass coordinate.
Conclusions. Our results imply that stars can only become supermassive by accretion at the rates of atomically cooled haloes (⇠0.1 10 M yr 1 ). Once they are supermassive, larger rates are possible.

Key words. Keywords should be given

1

•

1

1

1

1

1

1

1

9

1

5

1

2

1

1

1/2
r

undamental limit imposed by hydrostatic equilibrium

Haemmerlé et al.: Maximally accreting supermassive stars:

太字が今回の新しい計算

Table 1. Masses at the end of the run (set by numerical instability) in 105 M for the models with various accretion rates andmetallicities. The new models are marked in boldface.

ars: a fundamental limit imposed by hydrostatic equilibrium

この線より下は
静⽔圧平衡を維持し進化

本論⽂の主題の
質量降着率

Z=0
Z=Z

計算終了
5.63

10 M yr 1
5.43

100 M yr 1
4.251000 M yr 1
5.25
8.57

Table 2. Stellar mass at the beginning of H-burning.
Z=0
Z=Z

⾚波線より右下は
⼀般相対論的
不安定性に陥る

1 M yr 1
2.29
3.02

1 M yr 1
5000 M
2000 M

10 M yr 1
20 000 M

100 M yr 1
50 000 M

1000 M yr 1
100 000 M
40 000 M

Table 3. Mass fraction of the convective core at a stellar mass of
200 000 M .
Z=0Z=Z

1 M yr 1
20%
35%

10 M yr 1
5%100 M yr 1
0.5%1000 M yr 1
0.4%
1%Fig. 3. Limits to hydrostatic equilibrium for accretion above

g. 3.10 Limits
3
1 hydrostatic equilibrium for accretion above
本論⽂の主題のM 1yr to(without rotation). Sound waves can maintain full3
M yr (without
rotation).
Sound
waves
can
maintainare found to be slow rohydrostatic
equilibrium
only below
the solid
black
line.
Above full is not included, since accreting SMSs質量降着率
the dashedblack line,only
the wholehydrodynamidrostatic
equilibrium
belowmass
the evolves
solid blackline. Above tators due to the ⌦ -limit, which implies negligible impact of
cally. Stars
becomethe right-handof therotation on their structure (Haemmerlé et al. 2018a).
e dashed
black
line,GR
theunstablewhole atmass
evolves sidehydrodynamidashedline. The
maximum
rateatof theatomicallycooledsidehaloesof the
For zero metallicity, models with Ṁ  10 M yr 1 are ally.red
StarsbecomeGRunstableright-hand[結論]is shown by a grey dashed line. The location of the hylotropic,
ready described in Haemmerlé
etal.
(2018b).
Models of rapidly
d dashedline.
Themaximum
rateofatomicallycooledhaloesゼロメタルで100 Msu/yrなら
pure-accretion regime is indicated.
have not yet been pub⽔素核融合開始
SMS形成直後にSMBH形成可能
shown by a grey dashed line. The locationof the hylotropic, accreting SMSs at solar metallicity(ZAMSと呼んでいいのかな？） lished. For the present study, we compute four new models. We
re-accretion regime is indicated.
4.69
著者は明⽰的に述べていないが,first extend the set of Pop III models by considering the rates
この条件ではdark
100 M yr 1 and 1000 M
yr 1 . In ordercollapseは起こらない.
to understand the imstructures at first order, we can derive the acceleration term withdark collapseが起るかどうかは,対流核が”前主系列星”質量の0.5%となる
Eq. (5). We note that rP is re-derived fromthe new profiles: adpact of metals in various regimes,
we compute two models at so図３の⿊実線と⾚破線の計算範囲(ZAMSと呼んでいいのかな？）
jacent layers are delayed to di↵erent hydrostatic
models, which
lar metallicity with rates 1 andをひろげるとわかる(はず)1000 M yr 1 . We highlight the
4.30
uctures
atafirst
order,towederiveacceleration
term with fact that we do not include any mass-loss, even at solar metallicinduces
correctionrPcan
at theorigintheofthenon-zero acceler. (5).
We
note
that(5).
rPThe
is re-derived
the new
profiles: ad- ity, due to the short evolutionary timescales. We use initial modation
termin Eq.
condition forfrom
hydrostatic
equilibrium

20.

December 19, 2019

Characterization of mid-infrared polarization due to scattering in
protoplanetary disks

ABSTRACTS. Heeseof mid-infrared polarization due to scattering
1 et al.: Characterization
1
2
S. Heese , S. Wolf , and R. Brauer

Context.
It is generally assumed that magnetic fields play an important role in the formation and evolution of protoplanetary disks.
1

Key
words.
Polarization,
Radiative
transfer,
Scattering,
Protoplanetary
disks
Fig.
19:
Polarization
map forthe referencedisk.
For details,
seeGeneral[背景]100
Distance
140 pc. Sect.
Introduction
emission, it provides the line-of-sight magnetic field strength
中間⾚外線偏光観測の例は3.2.6.
Inclination
0 90(Crutcher et al. 1993). The total magnetic field can be derived

Pl [%]y [au]FI [Jy/as2]

y [au]

Institute of Theoretical Physics and Astrophysics, University of Kiel, Leibnizstraße 15, 24118 Kiel, GermanyOneobservationallyfieldsis to measurepolarized emission and absorption produced by magnetically2
CEAwaySaclayof– Serviced’Astrophysique,constraining
Orme des Merisiers,magnetic
Bât 709, 91191
Gif-sur-Yvette,
Francealignedelongated dust grains. The fact that radiation also becomes linearly polarizedhttps://arxiv.org/pdf/1910.06183.pdf
by light scattering at optical to millimeter
December 19, 2019wavelengths complicates magnetic field studies.
Aims. We characterize the linear polarizationof mid-infrared radiation due to scattering of the stellar radiation and dust thermal
ABSTRACT
re-emission
radiation (self-scattering).
Context. It is generally assumed that magnetic fields play an important role in the formation and evolution of protoplanetary disks.
Methods.
We computed
themagnetic
radialfields
polarization
profiles
at wavelengths
acrossby the
N and Q bands for a broad range of circumstellar
One way of observationallyconstraining
is to measure polarized
emissionand absorption produced
magneticallyalignedconfigurations.
elongated dust grains.These
The factsimulations
that radiation also
becomesas
linearlypolarized
by light scattering
at optical to millimeter
disk
serveda basisto analyze
the correlations
between selected disk parameters and the resulting
wavelengths complicates magnetic field studies.
linearpolarization.
Aims. We
characterize the linear図.2
polarization of mid-infrared radiation due to scattering of the stellar radiation and dust thermal
re-emission
radiation(self-scattering).
Results.
We find that the thermal re-emission radiation is stronger than the scattered stellar radiation for disks with inner holes smaller
Methods. We computed the radial polarization profiles at wavelengths across the N and Q bands for a broad range of circumstellar
200
than⇠ 10 au within
the considered
range. The
mid-infrared
polarization
due to scatteringshows
disk configurations.
These simulations
served as a basisparameterto analyze the correlations
betweenselected disk parameters
and the resulting
200 several clear trends: For
10 1
linear polarization.
scatteredstellar radiation only, the linear polarization degree decreases slightly with increasing radial
distance, while it increases with
Results. We find that the thermal re-emission radiation is stronger than the scattered stellar radiation for disks with inner holes smaller
80
図.6
10 2
100
100 and thermal re-emission
radialthermalre-emission
radiation
only anda combination
oftrends:
scatteredstellar radiation
than ⇠ 10distance
au within thefor
consideredparameter
range. The mid-infrared
polarization
due tofor
scatteringshows several clear
For
scattered stellar radiation only, the linear polarization degree decreases slightly with increasing radial distance, while it increases with
10 3
60
radiation.
radial distance for図.4
thermal re-emission radiation only and for a combination of scattered stellar radiation andthermalre-emission
00radiation.
Thelinear polarization degree decreases with increasing disk flaring and luminosity of the central
10 4 star. An increasing inner radius
40
The linear polarization degree decreases with increasing disk flaring and luminosity of the central star. An increasing inner radiusshiftstheincreaseofthelinearpolarization
degreefurther
outside,whilealargerscale
heightincreasesthelinearpolarization
degreeshifts the increase of the linear polarization degree further outside, while a larger scale height increases the linear
5
100 polarization degree
100
20
for smallradialdistancesand decreasesand
this degree
further outside.
For longerfurther
wavelengths,i.e., towardFor
the Qlonger
band in ourstudy, the
for
small
radialdistances
decreases
this degree
outside.
wavelengths,i.e.,10toward the Q band in our study, thelinear polarization degree converges more slowly.
99 %
10 6
linearpolarization
degreeconverges
moredueslowly.
Conclusions.
We found several
clear trends
for polarization
to scattering. These trends are the basis to200
distinguish polarization
200
0
200
100
0
100
200
200
100
0
100
200
due to scattering fromWe
polarization
to dichroicclear
emissionand absorption.
Conclusions.
founddueseveral
trendsfor polarization due to scattering.
These
trends
are
the
basis
to
distinguishpolarization
x [au]x [au]
Key words.
Polarization,from
Radiative
transfer, Scattering,
Protoplanetary
disks
due
to
scatteringpolarizationdue
todichroic
emissionandabsorption.
200
no. paper_tc
200

10 1
10 2

100

80
Pl [%]

3y [au]FI [Jy/as2]

y [au]

きわめて少ないが,
1060
8 µm 13thatµm and16 µm 22µm play an importantis Wavelengths
generally assumed
magnetic
fieldsby
combining
Zeemanobservations
withcomplimentary
obserDisk00
観測結果の解釈を提⽰しておきた。
ole in the formation and evolution of protoplanetarydisks. For vations of linear polarization (Heiles & Haverkorn 2012). How↵
1.0 2.5
10 4
40
xample,magnetic
fieldscancausestrongmagneticbraking,
Introduction 1.05 1.4
emission,
it provides
ever, the requirements for spatiallyresolved Zeeman
observa- the line-of-sight magnetic field strengthindering
formation of5 au
adependence
centrifugally
supported
disk(Galli3.2.6.
Inclinationof thepolarization
patternh100the
15 au
MIR
pol.tionsが検出するものは,
100
cannotbe met with current instruments(Brauer et al. 2017).
100
10 5 total magnetic(Crutcheretal.
1993).
The
field can be 20derived
rin
au oriented
50 au
t al. 2006).
A magnetic0.1fieldvertically withrespect
吸収偏光
+
放射偏光
+
散乱偏光
Only
spatially
unresolved
observations
arepredicted
tobefeasirout midplane
100 au 1000
au momentum
generally
assumedthatmagnetic
fieldsanimportant
98 %
osIn
the
disk
increases
angulartransport
in play
ble under
certaincircumstances.by combining Zeeman observations
withcomplimentary
obser
10 6
the
disk,
causing
MDustpreceding
10 6chapters,
M
10 3 M we considered a face-on
200
200
0
ccretion
disks
by
means
of
magnetorotational
instability
even
0
100
200
200
100
0
100
200
en aideal
incentro-symmetrictheformation
disks.
T Tauri starand evolution of protoplanetary
table
1参照
Most
observations
were200done
the submillimepolarization pattern as[設定]shown
inpolarization
Fig.
19. ForBevationsofin100linear2012).
How
magnetohydrodynamics
x [au] polarization (Heiles & Haverkorn
x [au]
M?
0.7 M (MHD). In nonideal MHD simter wavelengthbraking,regime. Rao et al. (2014) obtained spatially remple,
magnetic
can causestrongchanges
magnetic
lations,
thepresencehow
of a fields
verticalmagnetic
field ispatternimportant2theR polarization?
fore RTshowing
withinclinarequirements
solvedpolarization
observationsever,of IRASthe
16293- 2422 using the for spatially resolved Zeeman observa[結果]
4000areK consistent with observations
oderingprovide? accretion
rates
that
thebriefly
formationofthe
a centrifugally
diskArray.
(Gallition,etwediscusschangeoffields
the supported
net Submillimeter
linear polarization
The magnetic
derivedfrommetthisobHerbigAestartionsfieldcannotbewith(left)current
et al.for
2017)
Fig.
20:
Intensitymaps
and instruments
polarization (Brauer
maps (right)
Simon
al. 2013).
Despitethis importance,magneticin
偏光率の変化要因として考察した物理量：M?
2.5 M
servationsuggestsacomplexfield
structure.
For
HL
Tau,
analal.
2006).
A
magneticfieldoriented
verticallywithrespect
rotoplanetarydisksinclination.
are still2.48
poorlyconstrained
observationallydegreeThis value
is波⻑依存性,
zero forysis
a ダスト質量(分布),face-ondisk ow- an
R? withR
Only
spatiallyof
are 60predicted
to be
feasi
inclination
20 mm(top), observations
40 (center) and(bottom).
Forof spatially resolved polarization
observations
atunresolved
1.25
Stephensetal.
2017).
T
9500
K
?
円盤諸元
he
disk
midplaneincreasesofangularmomentumtransport
in details,
ing
to thecentro-symmetry
the polarization
pattern.
With inthe underorientationof the magnetic
seecertainSect.
3.2.6.
ble
circumstances.
One way of observationally constraining magnetic fields is to (Stephens et al. 2014) shows thatfield
iscoincident
with
the
major
axis
of
the
disk.
Table
1:
Model
parameters
creasing
inclination,
thenet
linearpolarization
degreeincreases
retionpolarized
disksemission
by means
instability evenmeasureproducedof
by magnetorotational
dichroic emission and

IRAS 09002-4732: A Laboratory for the Formation of Rich Stellar Clusters
Konstantin V. Getman1 , Eric D. Feigelson1, Michael A. Kuhn2 , Patrick S. Broos1 and Gordon P.
Garmire1
1

Department of Astronomy & Astrophysics, Pennsylvania State University, 525 Davey Lab, University Park, PA
16802; 2 Department of Astronomy, California Institute of Technology, Pasadena CA 91125
E-mail contact: gkosta at astro.psu.edu
IRAS 09002-4732 is a poorly studied embedded cluster of stars in the Vela Molecular Ridge at a distance of 1.7 kpc.
Deep observations with the Chandra X-ray Observatory, combined with existing optical and infrared surveys, produce acatalog of 441 probable pre-main sequence members of the region. The stellar spatial distribution has two components:
most stars reside in a rich, compact, elliptical cluster, but a minority reside within a molecular filament several
parsecs long that straddles the cluster. The filament has active distributed star formation with dozens of unclusteredprotostars. The cluster pre-main sequence population is ≤ 0.8 Myr old and deeply embedded; its most massive
member is extremely young producing an ultracompact H II region. The cluster total population deduced from the
X-ray luminosity function is surprisingly rich, twice that of the Orion Nebula Cluster. The cluster core is remarkably
dense where strong N-body interactions should be occurring; its Initial Mass Function may be deficient in massive
stars. We infer that IRAS 09002-4732 is a rare case where a rich cluster is forming today in a molecular filament,
consistent with astrophysical models of cluster formation in clouds that involve the hierarchical formation and merging
of groups in molecular filaments.
Accepted by The Astronomical Journal
https://arxiv.org/pdf/1910.04873

Chemical complexity in high-mass star formation: An observational and modeling casestudy of the AFGL 2591 VLA 3 hot core
C. Gieser1,2 , D. Semenov1,3 , H. Beuther1 , A. Ahmadi1,2 , J.C. Mottram1 , Th. Henning1 , M. Beltran4 ,L.T. Maud5 , F. Bosco1,2 , S. Leurini6 , T. Peters7 , P. Klaassen8 , R. Kuiper9 , S. Feng10,11 , J.S. Urquhart12 ,L. Moscadelli4, T. Csengeri13 , S. Lumsden14 , J.M. Winters15 , S. Suri1 , Q. Zhang16 , R. Pudritz17 , A.
Palau18 , K.M. Menten13 , R. Galvan-Madrid18 , F. Wyrowski13 , P. Schilke19 , Á. Sánchez-Monge19 , H.
Linz1 , K.G. Johnston14 , I. Jiménez-Serra20 , S. Longmore21 and T. Möller19
1

Max Planck Institute for Astronomy, Königstuhl 17,69117 Heidelberg, Germany; 2 Fellow of the International
Max Planck Research School for Astronomy and Cosmic Physics at the University of Heidelberg (IMPRS-HD); 3First detection of submillimeter-wave [13 C I] 3 P1 –3 P0 emission in a gaseous debris diskof 49 Ceti with ALMA
Aya E. Higuchi1 , Yoko Oya2 and Satoshi Yamamoto2
1

National Astronomical Observatory of Japan, Osawa, Mitaka, Tokyo 181-8588, Japan; 2 Department of Physics, The
University of Tokyo, Hongo, Bunkyo-ku, Tokyo 113-0033, Japan
E-mail contact: aya.higuchi at nao.ac.jp
We have detected the submillimeter-wave fine-structure transition (3 P1 –3 P0 ) of 13 C, [13 C I], in the gaseous debris diskof 49 Ceti with the Atacama Large Millimeter/submillimeter Array (ALMA). Recently, the [C I] 3 P1 –3 P0 emission
has been spatially resolved in this source with ALMA. In this dataset, the F =3/2–1/2 hyperfine component of [13 C
I], which is blue-shifted by 2.2 km s−1 from the normal species line, [C I], has been identified in the outer part of
the 49 Ceti disk, thanks to the narrow velocity widths of the gas components. The [C I]/[13 C I] line intensity ratio
is found to be 12±3, which is significantly lower than the 12 C/13 C abundance ratio of 77 in the interstellar medium.
This result clearly reveals that the [C I] emission is optically thick in 49 Ceti at least partly, as speculated by theprevious work. As far as we know, this is the first detection of [13 C I] 3 P1 –3 P0 emission at 492 GHz not only in debris
disks but also in the interstellar medium.
Accepted by ApJL
https://arxiv.org/pdf/1909.09748



Nonideal MHD Simulation of HL Tau Disk: Formation of Rings
Xiao Hu1 , Zhaohuan Zhu1 , Satoshi Okuzumi2 , Xue-Ning Bai3,4 , Lile Wang5 , Kengo Tomida6 and andJames M. Stone7
1

Department of Physics and Astronomy, University of Nevada, Las Vegas, 4505 South Maryland Parkway, Las Vegas,
NV 89154, USA; 2 Department of Earth and Planetary Sciences, Tokyo Institute of Technology, Meguro, Tokyo, 1528551, Japan; 3 Institute for Advanced Study, Tsinghua University, Beijing 100084, China; 4 Department of Astronomy,
Tsinghua University, Beijing 100084, China; 5 Center for Computational Astrophysics, Flatiron Institute, New York,NY 10010, USA; 6 Department of Earth and Space Science, Osaka University, Toyonaka, Osaka, 560-0043, Japan; 7Department of Astrophysical Sciences, 4 Ivy Lane, Peyton Hall, Princeton University, Princeton, NJ 08544, USA
E-mail contact: xiao.hu.astro at gmail.com



Nonideal MHD Simulation of HL Tau Disk: Formation of Rings
Xiao Hu1 , Zhaohuan Zhu1 , Satoshi Okuzumi2 , Xue-Ning Bai3,4 , Lile Wang5 , Kengo Tomida6 and andJames M. Stone7
1

Department of Physics and Astronomy, University of Nevada, Las Vegas, 4505 South Maryland Parkway, Las Vegas,
NV 89154, USA; 2 Department of Earth and Planetary Sciences, Tokyo Institute of Technology, Meguro, Tokyo, 1528551, Japan; 3 Institute for Advanced Study, Tsinghua University, Beijing 100084, China; 4 Department of Astronomy,
Tsinghua University, Beijing 100084, China; 5 Center for Computational Astrophysics, Flatiron Institute, New York,NY 10010, USA; 6 Department of Earth and Space Science, Osaka University, Toyonaka, Osaka, 560-0043, Japan; 7Department of Astrophysical Sciences, 4 Ivy Lane, Peyton Hall, Princeton University, Princeton, NJ 08544, USA
E-mail contact: xiao.hu.astro at gmail.com
Recent high-resolution observations unveil ring structures in circumstellar disks. The origin of these rings has been
widely investigated under various theoretical scenarios. In this work we perform global 3D nonideal MHD simulations
including eﬀects from both ohmic resistivity and ambipolar diﬀusion (AD) to model the HL Tau disk. The nonideal
MHD diﬀusion profiles are calculated based on the global dust evolution calculation including sintering eﬀects. Disk
16
ionization structure changes dramatically across the snow line due to the change of dust size distribution close to thesnow line of major volatiles. We find that accretion is mainly driven by disk wind. Gaps and rings can be quickly
produced from diﬀerent accretion rates across the snow line. Furthermore, AD leads to highly preferential accretionat the midplane, followed by magnetic reconnection. This results in a local zone of decretion that drains the massin the field reconnection area, which leaves a gap and an adjacent ring just outside of it. Overall, under favorable
conditions, both snow lines and nonideal MHD eﬀects can lead to gaseous gaps and rings in protoplanetary disks.

Accepted by ApJ
https://arxiv.org/pdf/1904.08899

An unusually large gaseous transit in a debris disc
Daniela P. Iglesias1,2, Johan Olofsson1,2, Amelia Bayo1,2 , Sebastian Zieba3 , Matas Montesinos1,2,4,
Jonathan Smoker5 , Grant M. Kennedy6 , Nicolas Godoy1,2 , Blake Pantoja7 , Geert Jan Talens8 , Zahed Wahhaj5 and Catalina Zamora1,2
1

Instituto de Fisica y Astronomia, Facultad de Ciencias, Universidad de Valparaiso, Av. Gran Bretana 1111, 5030

An unusually large gaseous transit in a debris disc
Daniela P. Iglesias1,2, Johan Olofsson1,2, Amelia Bayo1,2 , Sebastian Zieba3 , Matas Montesinos1,2,4,
Jonathan Smoker5 , Grant M. Kennedy6 , Nicolas Godoy1,2 , Blake Pantoja7 , Geert Jan Talens8 , Zahed Wahhaj5 and Catalina Zamora1,2
1

Instituto de Fisica y Astronomia, Facultad de Ciencias, Universidad de Valparaiso, Av. Gran Bretana 1111, 5030
Casilla, Valparaiso, Chile; 2 Nucleo Milenio de Formacion Planetaria - NPF, Universidad de Valparaiso, Av. Gran
Bretana 1111, Valparaiso, Chile; 3 Universitat Innsbruck, Institut fur Astro- und Teilchenphysik, Technikerstrasse 25,
6020 Innsbruck, Austria; 4 Chinese Academy of Sciences South America Center for Astronomy, National Astronomical
Observatories, CAS, Beijing 100012, China; 5 European Southern Observatory, Alonso de Cordova 3107, Vitacura,Santiago, Chile; 6 Department of Physics, University of Warwick, Gibbet Hill Road, Coventry CV4 7AL, UK; 7
Departamento de Astronomia, Universidad de Chile, Camino el Observatorio 1515, Las Condes, Santiago, Chile,
Casilla 36-D; 8 Institut de Recherche sur les Exoplanetes, Departement de Physique, Universite de Montreal, Montral,
QC H3C 3J7, Canada
E-mail contact: dpiglesi at uc.cl
We present the detection of an unusually large transient gas absorption in several ionized species in the debris discstar HD 37306 using high-resolution optical spectra. We have been analysing a large sample of debris discs searchingfor circumstellar gas absorptions aiming to determine the frequency of gas in debris discs. HD 37306 stood out
showing remarkably broad absorptions superimposed onto several photospheric Ca II, Fe II and Ti II lines. The
observed absorptions, unlike typical exocometary transits, lasted for at least eight days. Here we analyse simultaneous
spectroscopic and photometric data of the event and evaluate diﬀerent scenarios that might explain the observed
features. We conclude that the most likely scenario might be an exocometary break-up releasing a significant amountof gas close to the star, producing an occulting ring/torus shape.
Accepted by MNRAS
https://arxiv.org/pdf/1910.04747.pdf

Fingerprints of the protosolar cloud collapse in the Solar System II: Nucleosynthetic
anomalies in meteorites
Emmanuel Jacquet1 , Francesco C. Pignatale1,2 , Marc Chaussidon2 and Sébastien Charnoz2
1

Muséum national d’Histoire naturelle, UMR 7590, CP52, 57 rue Cuvier, 75005, Paris, France; 2 Institut de Physiquedu Globe de Paris (IPGP), 1 rue Jussieu, 75005, Paris, France

Fingerprints of the protosolar cloud collapse in the Solar System II: Nucleosynthetic
anomalies in meteorites
Emmanuel Jacquet1 , Francesco C. Pignatale1,2 , Marc Chaussidon2 and Sébastien Charnoz2
1

Muséum national d’Histoire naturelle, UMR 7590, CP52, 57 rue Cuvier, 75005, Paris, France; 2 Institut de Physiquedu Globe de Paris (IPGP), 1 rue Jussieu, 75005, Paris, France
E-mail contact: emmanuel.jacquet at mnhn.fr
The isotopic heterogeneity of the Solar System shown by meteorite analyses is more pronounced for its earliest objects,
the Calcium-Aluminum-rich Inclusions (CAIs). This suggests that it was inherited from spatial variations in diﬀerent
stardust populations in the protosolar cloud. We model the formation of the solar protoplanetary disk following its
collapse and find that the solid-weighted standard deviation of diﬀerent nucleosynthetic contributions in the disk is
reduced by one order of magnitude compared to the protosolar cloud, whose successive isotopic signatures are fossilized
by CAIs. The enrichment of carbonaceous chondrites in r-process components, whose proportions are inferred to havediminished near the end of infall, is consistent with their formation at large heliocentric distances, where the earlysignatures would have been preferentially preserved after outward advection. We also argue that thermal processing
had little eﬀect on the (mass-independent) isotopic composition of bulk meteorites for refractory elements.
17Distortion of Magnetic Fields in Barnard 68
Ryo Kandori1 , Motohide Tamura1,2,3 , Masao Saito2 , Kohji Tomisaka2 , Tomoaki Matsumoto4 , Nobuhiko
Kusakabe1 , Jungmi Kwon3 , Takahiro Nagayama5 , Tetsuya Nagata6 , Ryo Tazaki7 and Ken’ichi Tatematsu2
1

Astrobiology Center of NINS, 2-21-1, Osawa, Mitaka, Tokyo 181-8588, Japan; 2 National Astronomical Observatoryof Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan; 3 Department of Astronomy, The University of Tokyo, 7-3-1,
Hongo, Bunkyo-ku, Tokyo, 113-0033, Japan; 4 Faculty of Sustainability Studies, Hosei University, Fujimi, Chiyodaku, Tokyo 102-8160; 5 Kagoshima University, 1-21-35 Korimoto, Kagoshima 890-0065, Japan; 6 Kyoto University,
Kitashirakawa-Oiwake-cho, Sakyo-ku, Kyoto 606-8502, Japan; 7 Astronomical Institute, Graduate School of Science
Tohoku University, 6-3 Aramaki, Aoba-ku, Sendai 980-8578, Japan
E-mail contact: r.kandori at nao.ac.jp
The magnetic field structure, kinematical stability, and evolutionary status of the starless dense core Barnard 68 (B68)
are revealed based on the near-infrared polarimetric observations of background stars, measuring the dichroically polarized light produced by aligned dust grains in the core. After subtracting unrelated ambient polarization components,
the magnetic fields pervading B68 are mapped using 38 stars and axisymmetrically distorted hourglass-like magnetic
fields are obtained, although the evidence for the hourglass field is not very strong. On the basis of simple 2D and 3D
magnetic field modeling, the magnetic inclination angles on the plane-of-sky and in the line-of-sight direction are determined to be 47◦ ±5◦ and 20◦ ±10◦, respectively. The total magnetic field strength of B68 is obtained to be 26.1±8.7µG. The critical mass of B68, evaluated using both magnetic and thermal/turbulent support, is Mcr = 2.30 ± 0.20
M⊙ , which is consistent with the observed core mass of Mcore = 2.1 M⊙ , suggesting nearly critical state. We found a
relatively linear relationship between polarization and extinction up to AV ∼ 30 mag toward the stars with deepest
obscuration. Further theoretical and observational studies are required to explain the dust alignment in cold and dense
regions in the core.
Accepted by Accepted by PASJ
https://arxiv.org/pdf/1910.09121

Eﬀect of wind-driven accretion on planetary migration
C.N. Kimmig1, C.P. Dullemond1, W. Kley2
1

Zentrum für Astronomie, Heidelberg University, Albert-Ueberle-Str. 2, 69120 Heidelberg, Germany; 2 Institut für
Astronomie und Astrophysik, Universität Tübingen, Auf der Morgenstelle 10, 72076 Tübingen, Germany
E-mail contact: mail at carolin.kim