星形成ゼミ
2020年1⽉31⽇
古屋 玲#323 16-25

C. Gieser et al.: Chemical complexity in high-mass star formation:

3. Average spectrum (spatially averaged over 1.00 6⇥1.00 6, indicated by the white rectangle in Fig. 2) of the AFGL 2591 hot core labeled withecular emission lines with S /N > 5 (indicated by the red horizontal line, = 0.34 mJy beam 1 in the average spectrum). Unidentified lines areeled UL. A detailed table of the properties of the molecular transitions is given in Appendix A.

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Rotational evolution of solar-type protostars during the star-disk
interaction phase
F. Gallet1 , C. Zanni2 , and L. Amard3
1

Univ. Grenoble Alpes, CNRS, IPAG, 38000 Grenoble, France
INAF - Osservatorio Astrofisico di Torino, Strada Osservatorio 20, 10025, Pino Torinese, Italy
3
University of Exeter, Department of Physics & Astronomy, Stoker Road, Devon, Exeter, EX4 4QL, UK
2

Received -; accepted ABSTRACT
Context. The early pre-main sequence phase during which solar-mass stars are still likely surrounded by an accretion disk represents

a puzzling stage of their rotational evolution. While solar-mass stars are accreting and contracting, they do not seem to spin up
substantially.
Aims. It is usually assumed that the magnetospheric star-disk interaction tends to maintain the stellar rotation period constant (‘disklocking’), but this hypothesis has never been thoroughly verified. Our aim is to investigate the impact of the star-disk interaction
mechanism on the stellar spin evolution during the accreting pre-main sequence phases.
Methods. We devised a model for the torques acting on the stellar envelope based on studies of stellar winds, and we developed a
new prescription for the star-disk coupling founded on numerical simulations of star-disk interaction and magnetospheric ejections.
We then used this torque model to follow the long-term evolution of the stellar rotation.
Results. Strong dipolar magnetic field components up to a few kG are required to extract enough angular momentum so as to keep
the surface rotation rate of solar-type stars approximately constant for a few Myr. Furthermore an efficient enough spin-down torquecan be provided by either one of the following: a stellar wind with a mass outflow rate corresponding to ⇡ 10% of the accretion rate,
or a lighter stellar wind combined with a disk that is truncated around the corotation radius entering a propeller regime.
Conclusions. Magnetospheric ejections and accretion powered stellar winds play an important role in the spin evolution of solar-typestars. However, kG dipolar magnetic fields are neither uncommon or ubiquitous. Besides, it is unclear how massive stellar winds canbe powered while numerical models of the propeller regime display a strong variability that has no observational confirmation. Better
observational statistics and more realistic models could contribute to help lessen our calculations’ requirements.
Key words. Stars: solar-type – Stars: magnetic field – Stars: evolution – Stars: rotation – Stars: winds, outflows. Introduction

Classical T Tauri stars (CTTS) are magnetically active pre-mainequence stars surrounded by an accretion disk (Collier Cameron& Campbell 1993; Edwards et al. 1993; Collier Cameron et al.
995; Edwards et al. 1994; Hartmann et al. 1998). Their magetic field (up to a few kG, see Johns-Krull et al. 2009; Gregoryt al. 2012) has a strong impact on the dynamics and the strucure of the stellar environment and can lead to the truncation of
he disk and accretion of material along funnel flows down to thetellar surface, the launch of stellar winds along the opened magetic field lines (Matt & Pudritz 2005a, 2008; Matt et al. 2012;
Cranmer & Saar 2011; Matt et al. 2015; Réville et al. 2015; Johntone et al. 2015; See et al. 2018; Finley & Matt 2018), andhe ejection of material due to the magnetic star-disk interacon (Shu et al. 1994; Ferreira et al. 2000; Romanova et al. 2009;
Zanni & Ferreira 2013).
Additionally, observations suggest (see Edwards et al. 1993;
Bouvier et al. 1993; Rebull et al. 2004; Irwin & Bouvier 2009)highlight the apparent steady evolution of the three percentiles(90th, median, and 25th) of the rotational period distributionsof stars from 1 Myr to 10 Myr. The fa0866-5.30mBT8.9

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BC V9E 2E7,relaxation
Canadathermalandaccretion.
The
first
process,
appearinghttps://arxiv.org/pdf/1910.04776.pdf
School of Physics and Astronomy,Letter
Monash
VIC 3800, Australia
to theUniversity,
Editormathematicallyin the200240,
energyTsung-Dao Lee Institute, Shanghai
China equation, relies on the variationofentropyṡ. The supermassive
second one,stars:
appearingin the limit
boundary condi- needed to restor
Maximally
accretinga fundamental
Received; acceptedhydrostatic
core cannot incor
tions, reliesimposedon thebyvariation
ofequilibrium
mass
Ṁ.
The
efficiency
ofthese
ABSTRACTL. Haemmerlé , G. Meynet , L. Mayer , R. S. Klessen , T. E. Woods , A. Heger
Figure 1 show
two processes can be approximately estimated from their corContext. Major mergers of gas-rich galaxies provide promising conditions for the formation of supermassive black holes (SMBHs;
is obtained
timescales,time scales.
&responding
10 M ) by direct global
collapse because
they can triggerthe
massKelvin-Helmholtz
inflows as high as 10 10 M (KH)
yr on sub-parsec
However, the usingal.: case,
Maximally
accretingsupermassive
stars: awithout
fundamental
limit imposedhydrostatic equilibrium
channel of SMBHHaemmerlé
formation inetthis
either darkcollapse(direct collapse
prior stellar
phase) orby
supermassive
star (SMS;
andtheaccretiontime:& 10 M ), remains unknown.
M
Table
1. Masses
at the theendlimitof therun (set rateby numerical
Aims.
Here,
we investigate
in accretion
up to whichinstabilstars can maintain hydrostatic equilibrium.
5compute
R =equilibrium
260 Ra ⇥Methods.
of SMSs accretionaccreting at
1 – 1000ity) in We
10GM
M 2for hydrostaticthe modelsmodels
with various
rates
and M yr , and estimate the departures fromM
posteriori
by takingintoaccount
the finite
speedet ofsound.
metallicities.
The
newmodelsaremarkedinboldface.
⌧=,(1)
Haemmerléal.:Maximally
accretingsupermassive
stars:afundamental
limit
imposedbyhydrostatic
equilibrium
KH We find that stars accreting above the atomic cooling limit (& 10 M yr ) can only maintain hydrostatic equilibrium once
Results.
RL
1 M yr InTable
10 case,M
yrtheyatevolve100
M adiabatically
Mnumericalyrthey are supermassive.
this
with
a hylotropic
1. Masses
the endofyrthe run 1000
(set byinstabil-structure, that is, entropy is locally conserved and
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
rs: a fundamental Ṁ
limit imposed byZhydrostatic
equilibrium5.43
=0
2.29
5.25
Table
2. Stellar
mass
at
the
beginningof H-burning.4.25
Keywords.
Keywords
should
be
given
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
1respectively,
M yr
10 M yr100
M 000
yrGMis
1000theM yr gravitational conZ = and
Z
2000
M
40
dius
luminosity,andmostof &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 /
6
7

1

1

2

3,4

5

6,7

1

Département d’Astronomie, Université de Genève, chemin des Maillettes 51, CH-1290 Versoix, Switzerland
5 for Theoretical Astrophysics and Cosmology, Institute for Computational Science, University of Zurich, Winterthurerstrasse
4
5Center190, CH-8057 Zurich, Switzerland
3
Universität Heidelberg, Zentrum für Astronomie, Institut für Theoretische Astrophysik, Albert-Ueberle-Str. 2, D-69120 Heidelberg,
4
Germany
4
Universität Heidelberg, Interdisziplinäres Zentrum für Wissenschaftliches Rechnen, Im Neuenheimer Feld 205, D-69120
Heidelberg, Germany
5
1
National Research Council of Canada, Herzberg Astronomy & Astrophysics Research Centre, 5071 West Saanich Road, Victoria,
BC V9E 2E7, Canada
6
School of Physics and Astronomy, Monash University, VIC 3800, Australia
71Tsung-Dao Lee Institute, Shanghai 200240, China
1111
Received ; accepted251ABSTRACT

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

ars: a fundamental limit imposed by hydrostatic equilibrium

Haemmerlé et al.: Maximally accreting supermassive stars:

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.
Z=0
Z=Z

1 M yr 1
2.29
3.02

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

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
M 1yr to(without rotation). Sound waves can maintain full3
M yr (without
rotation).
Sound
waves
can
maintainhydrostatic
equilibrium
only below
the solid
black
line.
Above fullthe dashedblack line,only
the wholehydrodynamidrostatic
equilibrium
belowmass
the evolves
solid blackline. Above
cally. Stars
becomethe right-handof thee dashed
black
line,GR
theunstablewhole atmass
evolves sidehydrodynamidashedline. The
maximum
rateatof theatomicallycooledsidehaloesof thely.redStars
becomeGRunstableright-handis shown by a grey dashed line. The location of the hylotropic,
d dashed
line. The
maximum
rate of atomically cooled haloes
pure-accretion
regimeis indicated.
shown by a grey dashed line. The location of the hylotropic,
re-accretion regime is indicated.

structures at first order, we can derive the acceleration term with
Eq. (5). We note that rP is re-derived from the new profiles: adjacent layers are delayed to di↵erent hydrostatic models, whichuctures
atafirst
order,towederiveacceleration
term withinducescorrectionrPcan
at theorigintheof thenon-zero acceler. (5).
We
note
that(5).
rPThe
is re-derived
the new
profiles: adation
termin Eq.
condition forfrom
hydrostatic
equilibrium

is not included, since accreting SMSs are found to be slow rotators due to the ⌦ -limit, which implies negligible impact ofrotation on their structure (Haemmerlé et al. 2018a).
For zero metallicity, models with Ṁ  10 M yr 1 are already described in Haemmerlé et al. (2018b). Models of rapidly
accreting SMSs at solar metallicity have not yet been published. For the present study, we compute four new models. We
first extend the set of Pop III models by considering the rates
100 M yr 1 and 1000 M yr 1 . In order to understand the impact of metals in various regimes, we compute two models at solar metallicity with rates 1 and 1000 M yr 1 . We highlight thefact that we do not include any mass-loss, even at solar metallicity, due to the short evolutionary timescales. We use initial mod-

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,
seeGeneral100
Distance
140 pc. Sect.
Introduction
emission, it provides the line-of-sight magnetic field strength
3.2.6.
Inclination
0 90(Crutcher et al. 1993). The total magnetic field can be derived

10 2
3

Pl [%]FI [Jy/as2]

y [au]

200

10 1

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 polarization of mid-infrared radiation due to scattering of the stellar radiation and dust thermal
re-emission
radiation(self-scattering).
Results.
We findthat 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
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 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
10080
Pl [%]y [au]y [au]

1060
8 µm 13thatµm and16 µm 22µm play an importantis Wavelengths
generally assumed
magnetic
fieldsby
combining
Zeemanobservations
withcomplimentary
obserDisk00
ole in the formation and evolution of protoplanetary disks. 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
diskpolarization(Galli tions cannot
3.2.6.
Inclinationof thepatternh100the15 au
100
be 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 with respect 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
Mostobservations
were200done
the submillimepolarization pattern as 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
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,magnetic
in
M?
2.5 M
servationsuggestsacomplexfield
structure.
For
HL
Tau,
analal.
2006).
A
magneticfieldoriented
verticallywithrespect
rotoplanetarydisksinclination.
are still2.48
poorlyconstrained
observationallydegreeThis value
is zero forysis
a face-on
disk 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
?
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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