星形成ゼミ2019/11/15 SFN#322 26-35 (辰馬)26. Gaia 19ajj: A Young Star Brightening Due to Enhanced Accretion + Reduced Extinction
27. The New EXor Outburst of ESO-Hα 99 observed by Gaia ATLAS and TESS
28. Chronology of Episodic Accretion in Protostars - an ALMA survey of the CO and H2O snowlines
29. Understanding the origin of the magnetic ﬁeld morphology in the wide-binary proto- stellar system BHR 71
30. On the origin of wide-orbit ALMA planets: giant protoplanets disrupted by their cores
31. New four-color optical photometry of the young stars V521 Cyg and FHO 27
32. ALMA observations of water deuteration: A physical diagnostic of the formation of protostars
33. A stellar relic ﬁlament in the Orion star forming region
34. Exploring the conditions for forming cold gas giants through planetesimal accretion
35. Origin of hydrogen ﬂuoride emission in the Orion Bar: An excellent tracer for CO-dark H2 gas clouds

1http://arxiv.org/abs/1910.07290

30

On the origin of wide-orbit ALMA planets: giant protoplanets disrupted by their cores
J. Humphries1 & S. Nayakshin1
1

Department of Physics and Astronomy, University of Leicester, Leicester LE1 7RH, UK

E-mail contact: rjh73 at le.ac.uk
Recent ALMA observations may indicate a surprising abundance of sub-Jovian planets on very wide orbits in protoplanetary discs that are only a few million years old. These planets are too young and distant to have been formed via
26 than the gas clumps born in the classical Gravitational
the Core Accretion (CA) scenario, and are much less massive
Instability (GI) theory. It was recently suggested that such planets may form by the partial destruction of GI protoplanets: energy output due to the growth of a massive core may unbind all or most of the surrounding pre-collapse
protoplanet. Here we present the first 3D global disc simulations that simultaneously resolve grain dynamics in the discand within the protoplanet. We confirm that massive GI protoplanets may self-destruct at arbitrarily large separations
from the host star provided that solid cores of mass around 10–20 Earth masses are able to grow inside them during
their pre-collapse phase. In addition, we find that the heating force recently analysed by Masset and Velasco Romero(2017) perturbs these cores away from the centre of their gaseous protoplanets. This leads to very complicated dust
dynamics in the protoplanet centre, potentially resulting in the formation of multiple cores, planetary satellites, andother debris such as planetesimals within the same protoplanet. A unique prediction of this planet formation scenario
is the presence of sub-Jovian planets at wide orbits in Class 0/I protoplanetary discs.
Accepted by MNRAS
http://arxiv.org/pdf/1909.04395
2

2011). Both analytical work (Nayakshin 2015a) and 3D simulations
After its birth, this protopla(Humphries & Nayakshin
2018) on the&evolutionof GI2018)
protoplanets
ourprotoplanets
paper in Section 7: if rocky cores rapidly fo(Humphries
Nayakshindemonstrate ofthat
the disc. If accreted dust is ain their pebble-rich parent
discs.theNewly
bornofGI
protoplanets
the resultant release of energy may
accretemajoritymm
and above are
sized dustprotoplanets,
grains that enter
settle to the centre and form
1
1
initially very extended withof ∼1 au.
After a characteristic
objects
andleave super-Earth and potentially Satur
J. Humphries & S. Nayakshin
their radii
Hill spheres,considerablyenhancing their total
metal
content.
Nayakshin 2010, 2011). T
cooling time, hydrogen moleculesin thethese
protoplanet
dissociate
andto growstranded
at rapidly
tens to hundreds of au.
Once
accreted,
pebblesareexpected
andsettle
全束縛エネルギーに対する、解放される
potential energy and heat t
1of Physics and Astronomy,
oftheLeicester,
Leicester
LE1
UK coreit collapses toUniversity
a tightly boundanalogue (Bodenheimer
1974;
数Myrの円盤で中心星から遠くにある土星質量くらいの惑星の•Department
insideJupiter
protoplanet,likelybecoming
locked7RH,
into a massive
Assuming that the accretion
重力ポテンシャルエネルギーの割合
Helled et al. 2014). During
this1951;
cooling
phase
protoplanetsare(Kuiper
Boss,
Wetherill& Haghighipour
2002; Helled &surroundinggas
instantaneo
存在が、ALMA観測による円盤ギャップによって示唆されている
2
A
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E-mailcontact: rjh73 at le.ac.uk vulnerable to tidal disruption
via 2008;
interactions
centralstar 2010, 2011).
Schubert
Boley withet al.the2010;
Nayakshin
Stevenson 1982; Pollack et
or with other protoplanets, which
may destroyof these
nascent
we estimate
massneeded to unbind
Nayakshin(2016) many
showedvia 1D
populationHere
synthesis
models the core
total
corresponding
feedbac• コア集積では若くて遠すぎるので形成するのが難しい(ペブルplanets if they migrate tothat
separations
closer inside
than 20youngau (Humphries
GI protoplanetwith a radius of !a few au, which is ty
core formation
GI protoplanets may
in fact releaset
et al. 2019).
GM
collapseplanets.
We
assumethattheprotoplanetc Ṁ cenough heat to remove some or all of their gaseous envelopes.
′ can bを使っても若すぎる/十分な量のペブルがない)
5188
J. Humphries
and S.pebble
Nayakshin
Efb =
dt ,
26
↑原始惑星が崩壊する
Additionally,
accretion
plays
a
key
role
in
the
evolutionofa polytropicmaking
sphere with mass Mp , 0radiusrrcp , and adiabaThis process could therefore downsize GI protoplanets,
GI protoplanets. Observations(Tychoniec
et al. 2018)suggestthat for 7/5.
The Virial
theorem tells us that the total energy o
GI
a
more
physicallyplausible
scenario
hatching
ALMA
where M
rc are(Humphries & Nayakshin
2018)
on 0theevolution of GI
protoplanets
ofasourpaper in Section 7: if rocky cores rapidly
form
insideGIthe m• 古典的な重力不安定だと低質量な惑星は形成できない
c and
young
Class
protoplanetary
discs
containasmuchhundreds
its gravitational binding energy and is th(Etot ) is halfplanets. This idea is also attractive because of parallels
withofotherrespectively.
If wetheseassume t
in their pebble-rich parent
discs.
Newlyinborn
GI protoplanetsareprotoplanets,
the resultant
releaseofenergymay
disruptof Earth
massespebble-sizedgrains.
These
will
be
focusedby of their initialastrophysical systems, e.g. galaxies losing much
coreSaturnis constantinitially very extendedinsidewith protoplanets
radii of ∼1 as
au.theyAftera borncharacteristic
objects
andleave super-Earth and potentiallymass andcoreswe tak• 重力不安定で形成された原始惑星の中のコアの成長によって、
are(Boley&Durisen2010)
andgaseous mass due to energetic feedback from supermassive
black
2
3GM
rewrite rc in terms of Mc ancooling time, hydrogenaccreted
moleculesin theprotoplanet
dissociateandstrandedat tens to hundreds of au.p
duringanysubsequentmigration(Johansen&
Lacerda.
holes (Di Matteo, Springel & Hernquist 2005). Etot = −重力エネルギーが周りのガスに渡されることで、原始惑星が崩"# 1/3
it collapses to a tightly2010;
boundOrmel
Jupiter &analogue
1974;
5rp
Klahr (Bodenheimer
2010;
Baruteau,
Meru&
Paardekooper
3 4πρcIn Humphries & Nayakshin (2018), we used a sink particle apHelled et al. 2014). During
this analyticalcooling phase
protoplanets
are and 3D simulations
Efb =
G(Mc
2011). Bothwork
(Nayakshin
2015a)
壊する → コアが残されて惑星が形成できるのでは
proximation to study the accretionofApebbles
ontoGIEprotoplanets.
After
its
birth,
thisEprotoplanet
is5able toaccrete gas a
3
2
A
N
LY
T
I
C
A
L
S
T
I
M
AT
S
vulnerable to tidal disruption
via interactions
with2018)
the central
star that protoplanets(Humphries& Nayakshin
demonstrate
the disc.
accreted
The sink particle approach provides a reliable estimate
forIfthetotal dust is able to grow large enough,where
weahaveassumedthator with other protoplanets,
which
may
destroymanyandof these
nascent
accretethemajority
of
mm
above
sized
dust
grains
that
enter
Here
we
estimate
the
core
mass
neededtounbindnewlyborn
円盤内と原始惑星内の粒子の力学を分解して、3次元で円盤全体
settletothecentre
and
form
a
core
(Helled,Podolak&mass of pebbles captured by a protoplanet, but it does not allow us•~
10地球質量is liberated
instantaneously.
planets if they migratetheir
to separations
closerthan 20 au(Humphries
Hill spheres,considerablyenhancing
their
total
metal
content.
GI
protoplanet
witharadius
of
a
few
au,
which
is
typical
Nayakshin2010,
willprereleaseto study their subsequent evolution. In this paper,
we improveon 2011). This process forのシミュレーションを行った
as a canfraction
ofregionsthe total
et al. 2019).
Once accreted, these pebbles
are expected
growandcollapse
settleplanets. Wehydrodynamically,
assume thatthe protoplanet
be modelledasofenerg
potential
energyand heat thecentralthe
コア質量
our previous
worktobymodelling
therapidly
protoplanet"index# 1/3
Additionally, pebbleinsideaccretion
plays a key roleintheevolutionofthe protoplanet,likelybecoming
lockedintoamassivecoreapolytropic
sphere
with
mass
M
,
radiusr,
andadiabaticγ=ppAssumingthat
theaccretionenergy
of
the
core
is
thus resolving both gas and dust
dynamics
withingravitationalit. Nevertheless,it energyEfrom
Figure1. Liberated
potential
core formation
ρc as relefb
(Kuiper 1951;
Boss, etWetherill
Haghighipour
2002;
Helled& theorem
GI protoplanets. Observations(Tychoniecal.
2018)&
suggestthat computational= 1.3
The ofVirialtells
that
the totalenergyofa donepolytrope• まず、コアからのフィードバックなしで計算を行い、コア成surrounding
gas instantaneously,
as
is
in
the
CA
remains
beyondcurrentmeans
to resolvetheus
growinga7/5.
fraction
protoplanet
bindingenergyfromequation(4).
The
protoplanet−3Etot
5 g cmSchubert discs
2008; contain
Boley etasal.much
2010;asNayakshin
2010, 2011).
young Class 0 protoplanetary
hundreds
−3
)
is
half
of
its
gravitationalbindingenergyandistherefore
givena et a(E
tot
Stevenson
1982;
Pollack
et
al.
1996;
Mordasini
has
a
mass
of
5
M
and
a
solid
core
centraldensity
of
5
g
cm
.
It
takes
solid
core
in
such
simulations,
and
so
we
introduceadust-sinkinJ
長への影響が最も高いペブルサイズを決める
Nayakshin grains.
(2016) showed
via 1D
population
synthesis
modelsof Earth masses in pebble-sized
These
will
be
focusedalternatively,by
totalenergycoreof
17 core.
M⊕ toAs
unbindthiscorresponding
protoplanet
has ora radiusof 1 (Eau.fb ) increases a
the protoplanet centre to model thepebbles
accrete onif itfeedback
insideyoung
GI 2010)
protoplanets
may in fact releaseinside protoplanets asthat
theycore
areformation
born (Boley&
Durisenand"#2! energy
to the core, we pass the liberated gravitational
potential
to
2
t
3GMp
Mc
rp
GMc Ṁ c ′ Efb• 1サイズのペブルを用いてコアからのフィードバックありで計
enough heatmigration
to remove(Johansen
some or &all Lacerdaof their gaseous envelopes.
accreted during any subsequent
.(1)
Ethe
dt , for GI=
Etot grow= − faster.simulations
the surrounding gas in the protoplanet.
Ourthanthereforefb =cancoolingtime-scale
protoplanets
then
This
process
could therefore
downsize
GI protoplanets, making
Etot
Mp
rc
rc
5rp
0
2010;
Ormel
&
Klahr
2010;
Baruteau,
Meru&
Paardekooper
also aim to explore the effects
ofhavecorethefeedback
astointroduced
inunbind them.
算を行う
they
potential
completelya (Nayakshin
more physicallyplausible
scenario for hatching ALMA
2011). Both analyticalGI
work
2015a)
and3D
simulations
where
Mc andrc aretheLetmass
andof the lu
Nayakshin (2016).
us
protoplan
After
its birth, this protoplanet
is ableto accrete
gasconsider
andradius
dusta fromp:
原始惑星、c:コア
planets.2018)
This idea
is also attractive
because of parallels with other(Humphries & Nayakshin
demonstrate
that
protoplanets
respectively.
we assumethat core
theitrate
dust Maccre
The paper3 is structured theas follows.
In Section
2 ablewe briefly
withofrapidly
mass
disc. If accreted
dust
istoIfgrowlargerocky
enough,will
c =systems,losingaccrete the majority ofastrophysicalmm and above
sized e.g.
dust galaxies
grains thatentermuch of their initialcore is constant and we take a constant core densit

On the origin of wide-orbit ALMA planets: giant protoplanets disrupted by their corescs may be much larger, but we take 100 au in order to reduce thecal cost of the simulations.

4/5187/5565062 by National Astronomical Observatory of Japan user on 13 November 2019

distribution in fewer than four orbits.

Downloaded from https://academic.oup.com/mnras/article-abstract/489/4/5187/5565062

dot. Notice that the protoplanet has already carved a deep gap in the dust
tions are identical to those in Section 4.1, save that core
distribution in fewer than four orbits.
ckOn
is disabled.
the origin of wide-orbit ALMA planets: giant protoplanets disrupted by their cores
2 shows the surface1 density of the gas disc1 for a simulation
J. Humphries & S. Nayakshin
cm grains after 1273 yr, the protoplanet core is marked with a
コアからのフィードバックなし、1cmペブル(灰色の点)
Department
of50
Physics
Astronomy,
University
LE1 7RH, UK
dot1 andis orbiting at
au in anandanticlockwise
direction.
The of Leicester, Leicester
nE-mailof pebbles
in the discis marked
by the over-laid grey colour
contact:
rjh73
at le.ac.uk
ヒル半径内のダスト質量The protoplanet has carved a deep gap in the grains in lessur orbits. In agreement with our previous work (Humphries &shin 2018), this shows that pebble accretion is very efficient26
protoplanets. Fig. 3 shows a zoom view of the gas density
最大密度のthe protoplanet in the plane of the disc from the simulation
原始惑星からの距離
2. Recall that pre-collapse protoplanets are initially very
ed with radii of up to a few au. The core is marked with a
dot in the centre of the image while dust particles are marked
te. For the parameters of this protoplanet, 1 cm grains are
ed to sediment on time-scales of a few thousand years and so
原始惑星の半径↑
jority remain suspended between 0.4 and 0.8 au away fromntre of the protoplanet.
now examine these simulations in more detail for a range
in sizes, a = 0.1, a = 1, and a = 10 cm, which roughly
pond to Stokes numbers of 0.04, 0.4, and 4 in the disc. The
nel of Fig. 4 shows that the total mass of pebbles accreted
he disc inside the half Hill sphere of the protoplanet is very
for all of the three grain sizes. This result is consistent withhries & Nayakshin (2018) who found that pebble capture
he disc into the Hill sphere is almost 100 per cent efficiento Stokes numbers of 0.1. Despite similarities in the accretion
1 cmペブルは原始惑星の周りの上層部分にある
however,dashedline showscore is only
able
to in Figure 3. A zoom view on the protoplanet from the same simulation as
Figure 2.theGlobalgas surface
densitythat
for aaprotoplanetwith nofeedback
Figure 4. Dynamics inside the protoplanet for different grain sizes. Top:or athe10cm
grains.
simulation with 1 cm pebbles. The pebbles are plotted with overplotted Fig. 2, the core sink particle is marked with a black dot. The sedimentation Pebble accretion rates on to the protoplanet are comparable for mm,
middlepanelof Fig.
showsoftheradial isdistanceof thea black time for 1 cm grains is thousands of years and so the majority of accreted cm, and 10 cm grains (Humphries & Nayakshin 2018). Cores only growgreyscalepoints.
The 4location
themeanprotoplanet
marked withdot. Noticethat the protoplanet
has alreadycarved (defineda deep gapasinthe
the dust grains remain suspended in the upper atmosphere of the protoplanet between for decoupled 10 cm grains while 1 mm and 1 cm grains remain in theinsidethe protoplanetwith respect
to its centre
0.4 and 0.8 au from the centre.

4

atmosphere. Middle: Mean radial distance of dust from dust sink. Note therapid settling of a = 10 cm grains. The dotted line shows the contractionof radius that contains the initial mass of the protoplanet (rP (M0 )). Bottom:
total protoplanet and atmosphere only metallic compositions. Small grains
remain suspended in the atmosphere, causing it to become metal enriched.

コアが破壊される

tacc = 104 yr
Z0 = 0.3 per centZ0 = 3 per centHigh res.

10 000
1000
1000
1000

0.01
0.003
0.03
0.01

2
2
2
16

0.1
0.1
0.1
0.8

Downloaded from https://academic.oup.com/mnras/article-abst

grains not incorporated into the core are redistribThis restocks the reservoir of large grains for peb
On the origin of wide-orbit ALMA planets: giant protoplanets disrupted by their coresto subsequent protoplanets. In these simulations, dJ. Humphries1 & S. Nayakshin1
原始惑星の1
release between 10 and 40 Earth masses of meta
Department of Physics and Astronomy, University of Leicester, Leicester LE1 7RH, UK
軌道長半径
disc.
ガス円盤に、0.1倍のスケールハイトで10cmのペブルを足し、50auの• contact: rjh73 at le.ac.uk
E-mail
Our setup
is conservative in that we assume a mas
タイプI軌道移動
Makingsub-Jovian
planetsvia
GI5191
位置に半径3 auで5木星質量のポリトロープ球を原始惑星として置いた
metal free protoplanet in order to reduce uncertai
26 for the models presented in Fig. 5.
Table 1. Summary of run parametersconditions; in reality we expect a factor of two
ヒル半径内のNdisc and NPP represent the total number of SPH particles in the initial dischancementoftheprotoplanet(Boley&Durisen20
質量and protoplanet, respectively, there is typically one dust particle for every
disruption occurs in this conservative case demon
three SPH particles.
コア破壊
growth feedback provides a powerful mechanism f
Nametacc (yr)Z0
Ndisc (106 )NPP (106 )protoplanets.3 On the other hand, our large grains an
熱+ time-scales (choices made for numerical convenien
No feedback
–
0.01
2
0.1
tacc = 103 yr
1000
0.01
2
0.1
up the disruption. We address these assumptions fur
重力ポテンシャル
エネルギー

Figure 5. The effect of including feedback from solid cores built throughthat
theTopinitialmassof theseparation
protoplanetof(reach)), whichdustcontains
accretion.
panel:
orbitalcoresweare released
P (M0core,plot
as a dottedlineafterin themiddle panelof Fig. Upper
4. Themiddle
orangepanel: gas
on eccentric
orbitsprotoplanet
disruption.
line
grainsofare
ableprotoplanet.
to sedimentMiddle
inside thismassshowsinsidethat
theonly
half 10
Hillcmsphereeachpanel: sum ofboundary,smallergrains remainsuspended
atmospherethermal andgravitational
potential
energy inofthe
gasupperwithinthe half Hill sphere.
ofthe protoplanet.
Essentially,small grains
remaincoupledtheto core must
Lower
middle panel:core luminosity;
for this
protoplanettheaccretedwhilsta 10
cm grainsBottom
penetrate
intothethesolid lines
reachalmostdisc
10−3material,L⊙ to cause
disruption.
panel:
protoplanet.
plot the total metal mass inside the half Hill sphere. Dashed lines plot only
The bottom panel shows the fractional metallic composition of thethe core mass.
total protoplanet and of only the atmosphere material. 1 mm grains5
cause a significant enhancement to the metallicity of the atmosphere,

4.2 Inside the protoplanetsFigs 6–8 show different views of the protoplanet
1000 yr simulation in Fig. 5 at 681 yr. Figs 6 andコア光度
slices from this protoplanet in the x–y and x–z pla
that the density rises from 10−11 to 10−9 g cm−3 b
and inner regions. The overplottedblackarrowsin
原始惑星内
the velocity field. Fig. 7 shows that low-density gheated via feedback is escaping through the mid-pl
ダスト質量
a zoom view of internal temperature, centred ocentral density of the protoplanet.
core
コア質量 =Including
dust sinkgas in the centre of the protoplanet: once this h5192

J. Humphries and S. Nayakshin

On the origin of wide-orbit ALMA planets: giant protoplanets disrupted by their cores
J. Humphries1 & S. Nayakshin1
1of Physics and Astronomy, University of Leicester, Leicester LE1 7RH, UK
重力不安定で原始惑星が形成されると、ペブル集積•Department

によりダストが円盤から原始惑星の周りにやってき
E-mailcontact: rjh73 at le.ac.ukて、ダストが10 cmくらいまで大きくなると原始惑星に集積し、コアを作る

26• ペブル集積が速ければコアがすぐ作られ、周りのガスを加熱して原始惑星が破壊される• 例えば重力不安定による原始惑星形成から1700年で15地球質量くらいのコアが残る• 土星は100地球質量くらいなので、ガスの解放↑コア質量が一部であれば土星も形成可能• ペブル集積が抑えられたりダスト成長が遅けれ

←コア質量

↓コア質量

ば、コア成長によって原始惑星は破壊されない• 円盤初期で観測されているギャップを作る惑星を形成することは可能

Figure 5. The effect of including feedback from solid cores built throughdust accretion. Top panel: orbital separation of each core, cores are released
6
on eccentric orbits after protoplanet disruption. Upper middle panel: gas

http://arxiv.org/pdf/1909.04056

34

Exploring the conditions for forming cold gas giants through planetesimal accretion
Anders Johansen1 and Bertram Bitsch2
1

Lund Observatory; 2 Max Planck Institute for Astronomy

E-mail contact: anders at astro.lu.se
The formation of cold gas giants similar to Jupiter and Saturn in orbit and mass is a great challenge for planetesimaldriven core accretion models because the core growth rates far from the star are low. Here we model the growth andmigration of single protoplanets that accrete planetesimals and gas. We integrated the core growth rate using fits inthe literature to N -body simulations, which provide the eﬃciency of accreting the planetesimals that a protoplanet
migrates through. We take into account three constraints from the solar system and from protoplanetary discs: (1)the masses of the terrestrial planets and the comet reservoirs in Neptune’s scattered disc and the Oort cloud areconsistent with a primordial planetesimal population of afew Earth masses per AU, (2) evidence from the asteroid
28 planetesimal diameter is 100 km, and (3) observations
belt and the Kuiper belt indicates that the characteristicof protoplanetary discs indicate that the dust is stirred by weak turbulence; this gas turbulence also excites theinclinations of planetesimals. Our nominal model built on these constraints results in maximum protoplanet masses of
0.1 Earth masses. Ignoring constraint (1) above, we show that even a planetesimal population of 1,000 Earth masses,
corresponding to 50 Earth masses per AU, fails to produce cold gas giants (although it successfully forms hot andwarm gas giants). We conclude that a massive planetesimal reservoir is in itself insuﬃcient to produce cold gas giants.
The formation of cold gas giants by planetesimal accretion additionally requires that planetesimals are small and thatthe turbulent stirring is very weak, thereby violating all three above constraints.
Accepted by Astronomy & Astrophysics
http://arxiv.org/pdf/1909.10429
7

Exploring the conditions for forming cold gas giants through planetesimal accretion12

Anders Johansen and Bertram Bitsch

微惑星によるコア集積で木星や土星のような冷たい巨大ガス惑星を形成することは、コア成長速度が中心星から遠•Lund
Observatory; 2 Max Planck Institute for Astronomy

1

いほど遅いので難しい
E-mail contact: anders at astro.lu.se

Theformation of cold gas giants similar to Jupiter and Saturn in orbit and mass is a great challenge for planetesim• 微惑星とガスを降着する1つの原始惑星の成長と軌道移動を考えた
driven core accretion models because the core growth rates far from the star are low. Here we model the growth a
コア成長速度は先行研究のN体計算の結果を用いた•
migration of single protoplanets that accrete planetesimals and gas. We integrated the core growth rate using fitsthe •literature
to N -body simulations, which provide the eﬃciency of accreting the planetesimals that a protoplan
太陽系と原始惑星系円盤の観測から3つの制限を考慮した
migrates through. We take into account three constraints from the solar system and from protoplanetary discs:• (1) 地球型惑星と彗星の総質量が初期の微惑星総質量と一致すること(1 auあたり数地球質量)kmであること• (2) 小惑星帯とカイパーベルトから、微惑星の特徴的な直径は10028• (3) 円盤観測から、ダストは弱い乱流によりかき乱されていること(同時に微惑星の傾きも増加させる)• 考慮したのは、原始惑星のタイプI軌道移動、微惑星の降着、ガス乱流と微惑星同士の散乱による離心率と傾きの進化、ガス降着(10地球質量で開始)、円盤進化(粘性降着率、サイズ)• 円盤の粘性加熱は無視できるものとする

8

is
used
as
a
3.4.
Gas
accretion
etesimals
andhence
reinforces the shepherding behaviour (theionsequally
well
(Yang
et
al.
2012).
A&A 631, A70 (2019)
f Exploring
thebecause
growth theassumed
that
a
core
mass
of
10
M
marks
the
conditions
forformingcold
gas
giantsthroughplanetesimalaccretionE
run-away
shepherding
effect).
drag
turbulentstirring
parameterthat
appearsin
Eq.
(12)
ccretion. Our
Johansenetal.
(2019)
allowedgas
accretiontocommenceafter
γ = 10γ by= 10accretion.
This
simplified
approachguaranteesThe
accretion
rate
of
the
protoplanet
is
given
1
2
tricities
and0
heAnders
scalingfactor of theexcitation of 0the −1
Johansenand random-walk
Bertram Bitsch19),
who prethe core reached the pebble isolation mass;
here we insteadiscomparable
tothoseinferred
from
the
gia
1• 降着効率2 et al.
simal
eccentricity.
(2008)
suggested
that
⇠
↵
mall
pebblestṀassumed
= Ida
E Ṁ
=
E2⇡rṙ⌃.
(7)
Lund
Observatory;Max
PlanckInstitutefor
Astronomy
pla that a corepla mass of 10 M marksthetransition
to
gas
solar
system(Guillot
2005;
Wahl
et
al.
2017),
w
E
in
Eq.
(12)
o↵coreaccredenotes the turbulent diffusion coefficient, int hereE-mailcontact:
andersat
astro.lu.se
accretion.
This
simplified
approachguarantees
a
core
mass
that
turbulentstirring
parameteremphasisethecoreformation
processratherth•eation
to
those
of
Here
E
is
the
accretion
efficiency
and
Ṁ
is
the
flux
of
planof ↵thet to thevalue that controls the viscous evolution
pla of
1/2
iscomparable
tothoseinferred
from
the
giant
planets
inthegas
accretion.
The
expressions
that
we
use
for
sions
(and
in
The
formation
of
cold
gas
giantssimilarto
JupiterandSaturn
in
orbit
and
mass
is
a
great
challengeforplanetesim
etesimals
that
cross
the
orbit
of
the
migratingprotoplanet.
The
dtoplanetarythat ⇠ ↵disc),or ⇠ ↵t H/r (where H/r is the disctもしくはmigration
speed
isparameterised
throughtheinverse
of
the
solar
system(Guillot
2005;
Wahl
et
al.
2017),
whileitallowsrate
arediscussedindetail
in
Ida
et
al.
(2018)
andrivencore
accretionmodelsbecausethecore
growth
far
from
the
star
are
low.
Here
we
the
growtha
0 us
0rates1 model2
1
ratio).
Observationsof the thickness of the dust layer in
efficient,innormalised
migration
timescale,
木星、土星形成領域に注目すると、
migrationof single
protoplanets
that
accreteplanetesimals
and(2019).
gas.
WeWe
integrated
thecore
growth
rate
using
fits
•
chose
a
low
opacity
in
the
gas
enve
emphasisethecoreformation
processratherthanthephysicsoftoplanetaryofdisc around the young star HL Tau suggest a
evolution
the 降着効率は0.0001から0.03程度
literature togas
N -body
simulations,which 4provide
theweeﬃciency
ofthe
accretingthe planetesimals
that atimescale
protoplan
rapid
gas
accretion;
thegasaccretionaccretion.
The
expressions
that
use
for
gas
accretion˙b̃ = ⌧˜ 1 . of a few times 10 (Pinte et al.
ntcoefficientr migrates
isdiffusion
the disc(8)pthrough.
We
take
into
account
three
constraints
from
the
solar
systemandfromprotoplanetary
discs:
mig
planetisgenerallyafewhundredthousand
year
rate
are
discussedindetail
in
Ida
et
al.
(2018)
and
Johansenetal.
but• theinterpretation
depends on the assumed Stokes
降着効率を上げるには
dust
layer
in
(1999),
based(2019).
We
chose
a
low
opacityinthegas
envelope
tofacilitateThe
accretion
efficiency
isthenparameterised
as
of
the
sedimented
pebbles.
The
radialscaleheightofplanetesimal
au suggesta0
乱流を弱くする(左→右)23.5. Evolution of0 the protoplanetary
•
1
-1
disc
rapid
gas
accretion;
thegasaccretion
timescaleofaJupiter-mass
ons
here.
The
rings
observed
in
the
Disk
SubstructuresatHigh
Angular28(Pinte et al. E = ↵ b̃˙ pla 1 .
(9)pla
ynamicsnearpplanetisgenerallyafewhundredthousand
years
in
this
model.
ion
Project(DSHARP)
indicatesadiffusioncoefficient星への降着率を下げる(下→上)•Stokesmed
The protoplanetary disc is evolved as a standarar.
Distances
Dullemondetal.
2018).
The
physicalorigin
of
this
weak
The
parameters
↵
and
are
fits
to
numerical
simulations.
leprotoplanet,height
of
pla
pla
with
↵ = 0.01. The disc surface density profile
原始惑星の軌道移動が遅くな•
nceis
unknown;thevertical
shear
instability
andactiveTanaka&Ida(1999)
foundtheexpressions
3.5.
Evolutionoftheprotoplanetary
discigh
Angulartive) power law out to a characteristic radius b
T K /hp . Here
るから
layer
turbulence
both
yield
turbulent
diffusioncoeffisFig.
1. Accretion efficiency
of a protoplanet accreting planetesimals of radius 50 km and for weak turbulent stirring ( = 10 , left plots) and
1/3
n/(3M
coefficientprofileisexponentiallytapered.
We
used
analy
is
? )]
very
weak
turbulentstirring
(
=
10
,
right
plots)
asafunctionofthedistance
from
the
star
and
the
protoplanet
mass.
The
three
rows
show the3
The
protoplanetarydiscisevolvedasastandard
viscous↵
discR̃nexcessof
10
inthe
mid-plane(StollKley
2016;paccretion&
efficiency
for threedifferent increasing (top to bottom) values of the mass accretion rate through the disc; the gas surface density in turne
protoplanetfortheevolutionoftheaccretionof this weak ↵pla
defines the equilibrium inclination of the planetesimals
and thetransition of the protoplanet
from anmass
accreting predator(exterior of therate
red line)o
=
2.5
,
(10)
with
↵
=
0.01.
The
surfacedensity
profilefollowsa(nega2 disctoanon-accreting
shepherd(interior
of
the
red
line).
The
parameter
regionof
core
formation in the solar system (the box made by 5 and 20 AUtedal.
2018).
We
here
assumedarelatively
low
value
of
1
+
0.37
ĩ
/
R̃
the
contripthecharacteristic
discradius,anchoredby
0y4and active tive) power law out
and by 1 and 10 M ) is marked with dashed oflines. The
accretion
efficiency beyond 5 AU lies
at a valueof a few percent
or lower for the case
of c
to
a
characteristic
beyondtheweakturbulentstirring, whilewith
loweringradius
the turbulent stirring
to = 10 which
increases the efficiency
by approximately a factor 3. The efficiency overall⌘ thisas the nominal value ⇣becauseisconsistent
e protoplanet
0.17
2increases with decreasing gas accretion throughthe disc
becausemore slowly
migratingplanets accrete
astar,
higher fractionof, theplanetesimals
that
tial
gas
mass
flux
onto
the
Ṁ
and
the
fin
sion
coeffi0
=
0.79
1
+
10
ĩ
.
(11)
pla
they
migratethrough(eventhoughtheplanetesimal
inclinations
increaseslightly as the gas drag is reduced).
profileisexponentiallytapered.
We
usedanalyticalexpressions
ly
only
above
0
ations and likely represents a lower limit to the realistic after 3 Myrofdisc
evolution,
Ṁ
,
after
which
th
Kley
2016;
3
adii
of 50 kmfor ˜theof planetesimalthemassaccretion
rate
onto
the
star
and
9
values.
This valueof evolution
also yields
inclimore slowly migrating protoplanetsa larger o
gas accretion rate, Ṁ = 10 M yr , Ṁdisc= 10 isM assumed
yr , and rate because
to photoevaporate.
Theaccrete
choice
-3

10−1

10-1

10−2

10-2

Shepherd / Predator

102 10-1

101

100

101

r [AU]

101

100

100

10-1

10-1

10-2

10-2

10-3
10-1

10-3
102

101

102 10-1

100

r [AU]

-4

-3

102
101

dM*/dt = 10 -6MO• yr-1

-5dM*/dt = 10 MO• yr

-1

-2

-6

101

-5

100

102

10-1

-6

10-1

100

10-2

10-2

10-3
10-1

10-3
102

100

101

102 10-1

r [AU]

100

101

r [AU]

4

5

E

5

6

1

7

1

M [ME]r [AU]

101

-4

100

M [ME]

-4

102

-5

-2

-3

101

dM*/dt = 10 -7MO• yr-1

-4

dM*/dt = 10 -7MO• yr-1

102

10-3
102

-3

100

M [ME]

-2

-1

101
100

r [AU]M [ME]

102

100

10−3
10−1

M [ME]

dM*/dt = 10 -8MO• yr-1

-3

M [ME]

101

−4

−2

102

-5

dM*/dt = 10 -8MO• yr-1

−3−4

Mpla = 977 ME

3.0e-09
1.5e-08

100.0

Mpla = 209 ME

1000.0

30.0

300.0

100.0

0.10001.2e-070.01006.7e-07

1

10
10
10.0
10.0

55

11

0.00010.00011.0e-09
5.2e-09
1.2e-070.01001.0e-06

1.0e-09
5.2e-0930.0

0.1000dM*/dt [MO• yr-1]

3.0

1.5

300.0

10
10
r0 [AU]

300.0
1000.0100.0
300.0
1000.0r0 [AU]

9.8e-079.8e-07
0.00100.00101.0e-06100
100

1.0e-06
100

Mpla =10 209 ME

1000.03.00001.000010
r0 [AU]

30.0
20
30
.0
100.0
20
300.0
1010.0000.0

30.0

3.00001.000010

13.0e-09
1.5e-0896.7e-07dM*/dt [MO• yr-1]

81.0e-06
100

10.0

7

3.0

6

9.1e-079.9e-079.9e-07

1.5

5

1
5
10

10.0

-4

20

10

20

9

1
5
10

10.0
3.0

8

1.0e−07 0.0001
100
1

30.0
20
100.0
10.0
30
0.0
10
00.0
30.0

10.0

10

20

10

7

-3

2.6e-07

3.0

1
5

1

6

10
1.0e−07 0.0001
1.0e−07 0.0010
10
100
1

2.6e-07

10
r0 [AU]10-2

58.1e-071

4.7e-0759.6e−08 0.0100
10-4

8.1e-070.00019.1e-073.0

11.0e−07 0.0010

0.01000.001010

1.0e-08
4.7e-08

10.0

10-3

10.0
t0 [Myr]

3.0

0.00101-2

4.7e-0730.0

6.7e−08 0.1000

1t0 [Myr]

9.6e−08 0.0100

0.10000.0100100.0

3.0e−09 3.0000
1.2e−08 1.0000

30.0

20

10

10.0

0.10006.7e−08 0.1000

10Mpla = 434 M3E0.0
100.0

3.0
Mpla = 977 ME

1

300.0

3.00001.00001• 赤線はガスのToomre Q

20 10.030.0001• 黄線は円盤サイズ

105• 白の等高線は原始惑星の最終的な位

惑星が形成されてしまう

40 au↑

1
3.0
5

める位置

3.00001.0000• 左上以外は、hot & warmなガス巨大

−1
−2
−3

0.0010 R
disc
10
Q
rf r0 [AU]0.0001
1 Mpla = 434 ME 10
0.r00 [AU]00.0

• 横軸は原始惑星が微惑星を降着し始置

log(M/ME)5る星の降着率

0.01001

10.0010 R
disc
Q
rf
0.0001
1

t0 [Myr]

28
始める時間、縦軸(右)はそれに対応す

10.0

log(M/ME)

t0 [Myr]

•

1

5

3
The formation of cold gas giants similar to Jupiter and Saturn in orbit and mass is a greatchallenge for2 planetesimalMpla = 211 ME
0.1000driven coreaccretion models because the core growth rates far from the star are low. Here we model the
最終的な原始惑星質量
1 growth
3.0000 and
100.0
300.0
migration of single protoplanets that accrete planetesimals and gas. We integrated the core growth rate using
fits in
1.0000
4
0
the literature to N -body simulations, which provide the eﬃciency of accreting the planetesimals that a protoplanet 3
0.0100−1 discs: (1) 2
migrates through. We take into account three constraints from the solar system and from protoplanetary
縦軸(左)は原始惑星が微惑星を降着し−2 0.1000
1
−3
0

1

• 色は円盤が3 Myrで散逸したあとの

E-mail contact: anders at astro.lu.se

1.0e-08
4.7e-08

10.0

dM*/dt [MO• yr-1]

3.0e−09 3.0000
100.0
30.0
30.0
1.2e−08 cold
1.0000
A. Johansen and B. Bitsch: Forminggas giants through planetesimal accretion

100.0

dM*/dt [MO• yr-1]

Mpla = 211 ME

10.0

Exploring the conditions for forming cold gas giants through planetesimal accretion
3.0000Anders Johansen1 and Bertram Bitsch2
300.0
11.0000Lund Observatory; 2 Max Planck Institute for Astronomy
49.6e-07 9.6e-07
9.9e-079.9e-07
100

10
1
100
r0 [AU]
r0 [AU]
Fig. 2. Growth maps for protoplanet growth by planetesimal accretion, with gas accretion commencing after 10 ME is reached. The maps show thefinal formassprotoplanet(i.e. the massgrowthreached bywhen
the disc has dissipatedafter
3 Myr)
continuous
colours and theafter
final 10positionswhite contour
Fig. 2. Growth maps
planetesimalaccretion,
withgas asaccretioncommencing
ME isasreached.
The lines
mapsasshfunction of the starting position of the protoplanet and the starting time. We additionally show the size of the protoplanetary disc (yellow line)
final mass (i.e. theaand
massreached when the disc has dissipated after 3 Myr) as continuous colours and the final positions as white contour l
the Toomre Q of the gas (red contour lines). The panels show the results for four values of the initial mass accretion rate through the disc anda function of the thestarting
positionrateof afterthe protoplanet
and the startingtime.
We additionally
thepanel
protoplanetary
discmass accretion
3 Myr. The instantaneous
stellar massaccretion
rate is shownshow
on thetherightsize
axis.ofEachis also labelledwith(yellotheand the Toomre Qtotalof mass
the gas(red
contour lines).
The panels
showhere
thethatresultsvaluesof theinitialmassis accretion
rateplanetesimals
through theof theplanetesimals
in the models;
we assume1% offorthefour
initialgas surface
density
profileconverted intoat dt =rate0. Thenominalmodelleft) achievesstellar
protoplanet
only up
to 0.1
ME , approximately
the axis.
mass ofMars.panel
The threeothermodels, w
the mass accretionafter
3 Myr.
The(upper
instantaneous
massgrowthaccretion
rate
is shown
on the rightEachis alsolabelled
planetesimals
formmodels;at the gravitationallyunstablephase1%of protoplanetaryevolution,
form hotisandwarm gas into
giantsplanetesi
interior
total mass of the whereplanetesimals
in thewe assume
here thatof the initialdisc
gassurfacesuccessfully
density profile
convertedof 2 AU, but are not able to form any cold gas giants in the 5–10 AU region (see inset in the upper right panel).

• 左上: 標準モデル、右上: 降着率↑、
下2つ: 初期降着率↑円盤サイズ↓

t = 0. The nominal model (upper left) achieves protoplanet growth only up to 0.1 ME , approximately the mass of Mars. The three other mwhere planetesimals form at the gravitationallyunstable phase of protoplanetary disc evolution, successfully form hot and warm gas giants
10
The growth maps in Fig. 2 indicate the final mass of the plan- power-law gas surface density transitions to an exponential taperof 2 AU, but are not able to form any cold gas giants in the 5–10 AU region (see inset in the upper right panel).

giants.

Exploring the conditions for forming cold gas giants through planetesimal accretion
1 current positions of the giant planets
2 in the solar system. Start-

104

Anders Johansen ing
andBertram
Bitschat 5 AU and with planetesimals of 50 km, we can only form

γ = 10−4, Rpla = 50 km
γ = 10−5, Rpla = 50 km
γ = 10−6, Rpla = 50 km
γ = 10−4, Rpla = 5 km
γ = 10−5, Rpla = 5 km
γ = 10−6, Rpla = 5 km

32
10Jupiter-massplanets
endingin(warm)
1
AU
orbits.
Starting
far太陽系と原始惑星系円盤の観測からの3つの制限を考慮すると、原始惑
•Lund
Observatory; Max Planck Institute for Astronomy
ther out at 10 AU, smaller planetesimals help in promoting the2星の最大質量は0.1
地球質量にしかなれない
10
growthoftheprotoplanet,butthenominalplanetesimal
radiusE-mail contact: anders at astro.lu.seof 50 km and the smaller planetesimals with radius 5 km both1
10
to produce
1 ME inandmass.Saturn& warmな巨大The
formation of coldfailgas
giantsprotoplanets
similar beyond
to Jupiterin orbit and mass is a great challenge for planetesim• 微惑星総質量を増やして星への降着率も増やすと、hotIn order to explore the role of the migration rate on the
0are low. Here we model the growth adrivencore
accretionmodelsbecausethecore
growth
rates
far
from
the
star
10
ガス惑星が作られてしまう
outcome of planet formation, we produced additional growth
tracks with the migration
multipliedplanetesimals
by a reduction factor
fmig gas.
=migration of single protoplanetsthat accreteandWe integrated
the core growth rate using fits
−1
10
0.2
for
all
protoplanet
masses.
We
also
ignoredthetransition
微惑星総質量を増やして、乱流を非常に弱くして、微惑星サイズも小さ• literature to N -body simulations, which provide the eﬃciency of accreting the planetesimals that a protoplan
theto shepherding in this experiment because shepherding might−2
Mpla = 965 ME
10く(半径5km)すれば木星や土星のコアを形成することができる
migrates through. We
take intowhen
account
constraints
from the solar system and from protoplanetary discs:
be suppressed
multiple three
protoplanets
scatter planetesimals
t0 = 0 Myr−3between them (Kokubo & Ida 2000; Thommes et al. 2003). The
10
result is shown with yellow curves in Fig. 3. The reduced migra• この結論は3つの制限を全て無視してしまっている1
10
100
tion has a weak positive effect on the growth tracks starting at
rf [AU]
5 AU; they now produce gas giants in 2–3 AU orbits.28
The effect
成長中の原始惑星同士の相互作用が必要か？
•
is overall negative farther out, however. Lowering the migration Fig. 4. Final planetary mass as a function of the final planetary orbit
speed leads to a higher efficiency of accreting planetesimals, but for protoplanets starting to accrete at t = 0 in the massive planetesimal• 円盤の粘性加熱を考慮しないといけないか？
the flux of planetesimals that the protoplanets migrates over is population model with two different planetesimal radii (50 and 5 km)
4
5
and
three
valuesoftheturbulentstirring
coefficient
(10
,
10
, and
also reduced, which reduces the overall growth rate.
6
4
10
).
The
solid
blue
line
(
=
10
/Rpla = 50 km) corresponds to a cut
外側への軌道移動が必要か？
In Fig. 3 we show with dash-dotted blue lines the growth•
tracks when we started with a larger protoplanet mass of Mp,0 = at t0 = 0 of the upper right panel of Fig. 2. The planetesimal inclinations
are
here
set
as
the
maximumoftheinclinations
given
byturbulent
0.1 ME . These protoplanets experience rapid migration towards• ペブル降着によって形成されたのか？
stirring and by mutual planetesimal scattering. Planetesimals with athe star. This has the positive effect of bringing the protoplanets nominal radius of 50 km can only form Jupiter-mass planets interior
into the regions of higher accretion efficiency closer to the star, of 3 AU, even when the turbulent stirring is very weak with = 10 5
but increasing the protoplanet mass does not help to create giant or extremely weak with = 10 6 , because of inclination stirring by
planets in cold orbits.
mutual planetesimal scattering. Lowering the planetesimal radius to
We show the final planet mass as a function of the final posi- 5 km allows the formation of Jupiter analogues when the turbulenttion forming in the massive planetesimal population model in stirring coefficient is = 10 5 or = 10 6 .
Fig. 4 for lower values of the turbulent stirring .11We assumed

M [ME]

1

Gaia
19ajj:A
Young
Star
BrighteningDueto26 Enhanced Accretion + Reduced Extinction
Lynne A. Hillenbrand1 , Bo Reipurth2 , Michael Connelley2 , Roc M. Cutri3 and Howard Isaacson4,5
1

Department of Astronomy, California Institute of Technology, USA; 2 Institute for Astronomy, University of Hawaii
at Manoa, USA; 3 IPAC, California Institute of Technology, USA; 4 Astronomy Department, University of California,
Berkeley, USA; 5 University of Southern Queensland, Australia
E-mail contact: lah at astro.caltech.edu
We report on the source Gaia 19ajj, identifying it as a young star associated with a little-studied star-forming region
seen along a complex line-of-sight through the Gum Nebula. The optical lightcurve recently recorded by Gaia exhibitsa slow and unsteady 5.5 mag rise over about 3 years, while the mid-infrared lightcurve from NEOWISE over thesame time period shows a 1.2 mag rise having similar structure. Available color information is inconsistent with pureextinction reduction as the cause for the photometric brightening. Optical spectroscopic characteristics in the current
bright phase include: little in the way of absorption except for the hallmark LiI 6707 Å signature of youth plus weake.g. Ca I and notably Ba II; strong wind/outflow in Ca II, Mg I b, Na I D, Hα, K I, O I; jet signatures in [O I], [SII], [Ca II], [Fe II], and [Ni II]; and narrow rest-velocity emission in neutral species such as Fe I, Ni I, and Mg I. The
infrared spectrum is also characterized by outflow and emission, including: a hot He I wind, jet lines such as [Fe II]
and H2 ; and weak narrow rest-velocity atomic line emission. The 12 CO bandheads are weakly in emission, but there
is also broad H2 O absorption. Gaia 19ajj exhibited a previous bright state in the 2010-2012 time frame. The bodyof photometric and spectroscopic evidence suggests that the source bears resemblance to V2492 Cyg (PTF 10nvg)
and PV Cep, both of which similarly experience bright phases that recur on long timescales, with large-amplitude
photometric variations and emission-dominated spectra. We interpret the behavior of Gaia 19ajj as caused by cycles
of enhanced disk accretion accompanied by reduced extinction.
24
Accepted by Astronomical Journal (AJ)
http://arxiv.org/pdf/1910.05790

12

27

The New EXor Outburst of ESO-Hα 99 observed by Gaia ATLAS and TESS
Klaus W. Hodapp1 , Bo Reipurth1 , Bertil Pettersson2 , John Tonry3 , Larry Denneau3 , Patrick J. Vallely4 ,Benjamin J. Shappee3 , James D. Armstrong5 , Michael S. Connelley1 , C. S. Kochanek4 , MichaelFausnaugh6 , Rolf Chini7 , Martin Haas7 and Catalina Sobrino Figaredo7
1

University of Hawaii, Institute for Astronomy, 640 N. Aohoku Place, Hilo, HI 96720, USA; 2 Department of Physics
and Astronomy, Uppsala University, Box 516, SE-751 20 Uppsala, Sweden; 3 University of Hawaii, Institute for
Astronomy, 2680 Woodlawn Drive, Honolulu, HI 96822, USA; 4 Department of Astronomy, The Ohio State University,
140 West 18th Avenue, Columbus Ohio, 43210-1173, USA; 5 University of Hawaii, Institute for Astronomy, 34 Ohia Ku
St., Pukalani, HI 96768, USA; 6 MIT Kavli Institute for Space and Astrophysics Research, 77 Massachusetts Avenue,
37-241, Cambridge, MA 02139, USA; 7 Astronomisches Institut, Ruhr-Universität Bochum, Universitätsstraße 150,
44801 Bochum, Germany
E-mail contact: hodapp at hawaii.edu
We report photometry and spectroscopy of the outburst of the young stellar object ESO-Hα 99. The outburst was
first noticed in Gaia alert Gaia18dvc and later by ATLAS. We have established the outburst light curve with archival
ATLAS “Orange” filter photometry, Gaia data, new V -band photometry, and J, H, and Ks photometry from IRIS
and UKIRT. The brightness has fluctuated several times near the light curve maximum. The TESS satellite observed
ESO-Hα 99 with high cadence during one of these minor minima and found brightness fluctuations on timescales ofdays and hours. Imaging with UKIRT shows the outline of an outflow cavity, and we find one knot of H2 1 − 0 S(1)emission, now named MHO 1520, on the symmetry axis of this nebula, indicating recent collimated outflow activity
from ESO-Hα 99. Its pre-outburst SED shows a flat FIR spectrum, confirming its early evolutionary state and its
similarity to other deeply embedded objects in the broader EXor class. The pre-outburst luminosity is 34 L⊙ , amuch higher luminosity than typical EXors, indicating that ESO-Hα 99 may be a star of intermediate mass. Infraredand optical spectroscopy show a rich emission line spectrum, including HI lines, strong red CaII emission, as well as
infrared CO bandhead emission, all characteristic EXors in the broadest sense. Comparison of the present spectrawith an optical spectrum obtained in 1993, presumably in the quiescent state of the object, shows that during thepresent outburst the continuum component of the spectrum has increased notably more than the emission lines. The
Hα equivalent width during the outburst is down to one half of its 1993 level, and shock-excited emission lines aremuch less prominent.
Accepted by The Astronomical Journal

13

28

Chronology of Episodic Accretion in Protostars - an ALMA survey of the CO and H2Osnowlines
Tien-Hao Hsieh1 , Nadia M. Murillo2, Arnaud Belloche3 , Naomi Hirano1 , Catherine Walsh4 , Ewine F.
van Dishoeck2,5 , Jes K. Jørgensen6 and Shih-Ping Lai7
1

Institute of Astronomy and Astrophysics, Academia Sinica, P.O. Box 23-141, Taipei 106, Taiwan; 2 Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA, Leiden, the Netherlands; 3 Max-Planck-Institut für Radioastronomie,
Auf dem Hügel 69, 53121 Bonn, Germany; 4 School of Physics and Astronomy, University of Leeds, Leeds LS2 9JT,
UK; 5 Max-Planck-Institut für extraterrestrische Physik, Giessenbachstraße 1, 85748, Garching bei München, Germany; 6 Niels Bohr Institute, University of Copenhagen, Øster Voldgade 57, DK 1350 Copenhagen K., Denmark; 7Institute of Astronomy, National Tsing Hua University (NTHU), Hsinchu 30013, Taiwan
E-mail contact: shawinchone at gmail.com
Episodic accretion has been used to explain the wide range of protostellar luminosities, but its origin and influence on
the star forming process are not yet fully understood. We present an ALMA survey of N2 H+ (1 − 0) and HCO+ (3 − 2)
toward 39 Class 0 and Class I sources in the Perseus molecular cloud. N2 H+ and HCO+ are destroyed via gas-phasereactions with CO and H2 O, respectively, thus tracing the CO and H2 O snowline locations. A snowline location at
25 luminosity suggests that an accretion burst has occurreda much larger radius than that expected from the current
in the past which has shifted the snowline outward. We identified 18/18 Class 0 and 9/10 Class I post-burst sourcesfrom N2 H+ , and 7/17 Class 0 and 1/8 Class I post-burst sources from HCO+ . The accretion luminosities during thepast bursts are found to be ∼ 10 − 100 L⊙ . This result can be interpreted as either evolution of burst frequency or
disk evolution. In the former case, assuming that refreeze-out timescales are 1000 yr for H2 O and 10,000 yr for CO,
we found that the intervals between bursts increases from 2400 yr in the Class 0 to 8000 yr in the Class I stage. Thisdecrease in the burst frequency may reflect that fragmentation is more likely to occur at an earlier evolutionary stage
when the young stellar object is more prone to instability.
Accepted by ApJ
https://arxiv.org/pdf/1909.02706

14

29

Understanding the origin of the magnetic field morphology in the wide-binary protostellar system BHR 71
Charles L. H. Hull1,2,11 Valentin J. M. Le Gouellec3,4 Josep M. Girart5,6 John J. Tobin7,8,9 Tyler L.
Bourke101

National Astronomical Observatory of Japan, NAOJ Chile, Alonso de Córdova 3788, Oﬃce 61B, 7630422, Vitacura,Santiago, Chile; 2 Joint ALMA Observatory, Alonso de Córdova 3107, Vitacura, Santiago, Chile; 3 European Southern Observatory, Alonso de Córdova 3107, Vitacura, Santiago, Chile; 4 AIM, CEA, CNRS, Université Paris-Saclay,
Université Paris Diderot, Sorbonne Paris Cité, F-91191 Gif-sur-Yvette; 5 Institut de Ciències de l’Espai (ICE-CSIC),
Campus UAB, Carrer de Can Magrans S/N, E-08193 Cerdanyola del Vallès, Catalonia; 6 Institut d’Estudis Espacials
de Catalunya, E-08030 Barcelona, Catalonia; 7 National Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, Virginia 22903, USA; 8 Homer L. Dodge Department of Physics and Astronomy, University of Oklahoma,
440 W. Brooks Street, Norman, OK 73019, USA; 9 Leiden Observatory, Leiden University, P.O. Box 9513, NL-2300RA
Leiden, The Netherlands; 10 Square Kilometre Array Organization, Jodrell Bank Observatory, Lower Withington,
Macclesfield, Cheshire SK11 9DL, UK; 11 NAOJ Fellow

E-mail contact: chat.hull at nao.ac.jp
We present 1.3 mm ALMA observations of polarized dust emission toward the wide-binary protostellar system BHR
71 IRS1 and IRS2. IRS1 features what appears to be a natal, hourglass-shaped magnetic field. In contrast, IRS2
exhibits a magnetic field that has been aﬀected by its bipolar outflow. Toward IRS2, the polarization is confined
mainly to the outflow cavity walls. Along the northern edge of the redshifted outflow cavity of IRS2, the polarized
emission is sandwiched between the outflow and a filament of cold, dense gas traced by N2 D+ , toward which no dust
polarization is detected. This suggests that the origin of the enhanced polarization in IRS2 is the irradiation of theoutflow cavity walls, which enables the alignment of dust grains with respect to the magnetic field—but only to adepth of ∼ 300 au, beyond which the dust is cold and unpolarized. However, in order to align grains deep enough inthe cavity walls, and to produce the high polarization fraction seen in IRS2, the aligning photons are likely to be in themid- to far-infrared range, which suggests a degree of grain growth beyond what is typically expected in very young,Class 0 sources. Finally, toward IRS1 we see a narrow, linear feature with a high (10–20%) polarization fraction anda well ordered magnetic field that is not associated with the bipolar outflow cavity. We speculate that this feature
may be a magnetized accretion streamer; however, this has yet to be confirmed by kinematic observations of dense-gas
tracers.
Accepted for publication in the Astrophysical Journal
http://arxiv.org/abs/1910.07290

15

Accepted by MNRAS

31
http://arxiv.org/pdf/1909.04395
New four-color optical photometry of the young stars V521 Cyg and FHO 27
Sunay Ibryamov1 and Evgeni Semkov2
1

Department of Physics and Astronomy, Faculty of Natural Sciences, University of Shumen, 115, Universitetska Str.,
9700 Shumen, Bulgaria; 2 Institute of Astronomy and National Astronomical Observatory, Bulgarian Academy of
Sciences, 72, Tsarigradsko Shose Blvd., 1784 Sofia, Bulgaria
E-mail contact: sibryamov at shu.bg
Results from new BVRI photometric observations of the pre-MS stars V521 Cyg and FHO 27 collected during theperiod from February 2015 to July 2019 are presented. These stars are embedded in the LDN 935, dubbed “The Gulf
of Mexico” in the active star-forming complex NGC 7000/IC 5070. Previous studies revealed them as variables withlarge amplitude brightness variations. Our new observations show that the stars continue to exhibit strong photometricvariability and fading events in the brightness. On the basis of the new data received and the data available in theliterature, we were able to specify the periodicity in the brightness variations of V521 Cyg. In the case of FHO 27, wefound a significant peak in its periodogram corresponding to 893 days period.
Accepted by Bulgarian Astronomical Journal
http://arxiv.org/pdf/1909.01458

ALMA observations of water deuteration: A physical diagnostic of the formation of
protostars
16

32

ALMA observations of water deuteration: A physical diagnostic of the formation of
protostars
S. S. Jensen1 , J. K. Jørgensen1 , L. E. Kristensen1 , K. Furuya2 , A. Coutens3 , E. F. van Dishoeck4,5 , D.
Harsono4 and M. V. Persson6
1

Niels Bohr Institute & Centre for Star and Planet Formation, University of Copenhagen, Øster Voldgade 5-7, DK1350 Copenhagen K, Denmark; 2 Center for Computational Sciences, University of Tsukuba, Japan; 3 Laboratoired’Astrophysique de Bordeaux, Univ. Bordeaux, CNRS, B18N, alle Geoﬀroy Saint-Hilaire, 33615 Pessac, France; 4Leiden Observatory, Leiden University, PO Box 9513, 2300 RA Leiden, The Netherlands; 5 Max–Planck Institute für
extraterrestrische Physik (MPE), Giessenbachstrasse, 85748 Garching, Germany; 6 Department of Space, Earth andEnvironment, Chalmers University of Technology, Onsala Space Observatory, 439 92 Onsala, Sweden
E-mail contact: sigurd.jensen at nbi.ku.dk
How water is delivered to planetary systems is a central question in astrochemistry. The deuterium fractionation ofwater can serve as a tracer for the chemical and physical evolution of water during star formation and can constrain
the origin of water in Solar System bodies.
The aim is to determine the HDO/H2 O ratio in the inner warm gas toward three low-mass Class 0 protostars selected
to be in isolated cores, i.e., not associated with any cloud complexes. Previous sources for which the HDO/H2 O
ratio have been established were all part of larger star-forming complexes. Determining the HDO/H2 O ratio toward
three isolated protostars allows comparison of the water chemistry in isolated and clustered regions to determine theinfluence of local cloud environment.
27
We present ALMA Band 6 observations of the HDO 31,2 –22,1 and 21,1 –21,2 transitions at 225.897 GHz and 241.562 GHz
along with the first ALMA Band 5 observations of the H18
2 O 31,3 –22,0 transition at 203.407 GHz. The high angularresolution observations (0.3′′ -1.3′′ ) allow the study of the inner warm envelope gas. Model-independent estimates for
the HDO/H2 O ratios are obtained and compared with previous determinations of the HDO/H2 O ratio in the warm gas
toward low-mass protostars. We successfully detect the targeted water transitions toward the three sources with S/N ¿5. We determine the HDO/H2 O ratio toward L483, B335 and BHR71–IRS1 to be (2.2 ± 0.4)×10−3, (1.7 ± 0.3)×10−3,
and (1.8 ± 0.4)×10−3, respectively, assuming Tex = 124 K. The degree of water deuteration of these isolated protostars
are a factor of 2–4 higher relative to Class 0 protostars that are members of known nearby clustered star-forming
regions.
The results indicate that the water deuterium fractionation is influenced by the local cloud environment. This eﬀect
can be explained by variations in either collapse timescales or temperatures, which depends on local cloud dynamics
and could provide a new method to decipher the history of young stars.
Accepted by A&A
https://arxiv.org/pdf/1909.10533

17https://arxiv.org/pdf/1909.10533

33

A stellar relic filament in the Orion star forming region
Tereza Jerabkova1,2,3 , Henri M.J. Boﬃn1 , Giacomo Beccari1 , and Richard I. Anderson1,4
1

2

European Southern Observatory, Karl-Schwarzschild-Straße 2, 85748 Garching bei München, Germany; Helmholtz
Institut für Strahlen und Kernphysik, Universität Bonn, Nussallee 14-16, 53115 Bonn, Germany; 3 AstronomicalInstitute, Charles University in Prague, V Holesovickach 2, CZ-180 00 Praha 8, Czech Republic; 4 Départementd’Astronomie, Université de Genève, 51 Ch des Maillettes, 1290 Sauverny, Switzerland
E-mail contact: tjerabkova at gmail.com
We report the discovery of the oldest stellar substructure in the Orion star forming region (OSFR), the Orion relicfilament. The relic filament is physically associated with the OSFR as demonstrated by Gaia DR2 photometry andastrometry, as well as targeted radial velocity follow-up observations of a bright sub-sample of proper-motion selectedcandidate members. Gaia DR2 parallaxes place the Orion relic filament in the more distant part of the OSFR, ≈430pc from the Sun. Given its age, velocity dispersion, spatial extent, and shape, it is not possible for the Orion relicfilament to have formed as a single stellar cluster, even taking into account residual gas expulsion. The relic filament isalso too young to be a tidal stream, since Galactic tides act on much longer time scales of order 100 Myr. It therefore
appears likely that the structure formed from a molecular cloud filament similar to Orion A in the OSFR and retained
its morphology despite decoupling from its natal gas. Hence, the Orion relic filament bears witness to the short-lived
evolutionary phase between gas removal and dispersion due to shears and tides, and provides crucial new insights into
how stars are formed in molecular clouds.
Accepted by MNRAS
http://arxiv.org/pdf/1909.04056
18

35 Origin of hydrogen fluoride emission in the Orion Bar: An excellent tracer for CO-dark
H2 gas clouds

Umit Kavak1,2 , F.F.S. van der Tak2,1 , A.G.G.M. Tielens3 and R.F. Shipman2,11

Kapteyn Astronomical Institute, University of Groningen, P.O. Box 800, 9700 AV Groningen, The Netherlands;
2
SRON Netherlands Institute for Space Research, Landleven 12, 9747 AD Groningen, The Netherlands; 3 Leiden
Observatory, Leiden University, PO Box 9513, NL-2300RA, Leiden, the Netherlands
E-mail contact: kavak at astro.rug.nl

The hydrogen fluoride (HF) molecule is seen in absorption in the interstellar medium (ISM) along many lines of sight.
Surprisingly, it is observed in emission toward the Orion Bar, which is an interface between the ionized region aroundthe Orion Trapezium stars and the Orion molecular cloud. We aim to understand the origin of HF emission in theOrion Bar by comparing its spatial distribution with other tracers. We examine three mechanisms to explain the HF
emission: thermal excitation, radiative dust pumping, and chemical pumping. We used a Herschel /HIFI strip map of
the HF J = 1 → 0 line, covering 0.5′ by 1.5′ that is oriented perpendicular to the Orion Bar. We used the RADEX
non-local thermodynamic equilibrium (non-LTE) code to construct the HF column density map. We use the Meudon
PDR code to explain the morphology of HF. The bulk of the HF emission at 10 km s−1 emerges from the CO-darkmolecular gas that separates the ionization front from the molecular gas that is deeper in the Orion Bar. The excitationof HF is caused mainly by collisions with H2 at a density of 105 cm−3 together with a small contribution of electrons
in the interclump gas of the Orion Bar. Infrared pumping and chemical pumping are not important. We conclude
that the HF J = 1 → 0 line traces CO-dark molecular gas. Similarly, bright photodissociation regions associated withmassive star formation may be responsible for the HF emission observed toward active galactic nuclei.
Accepted by Astronomy and Astrophysics (A&A) on September 23, 2019.
http://arxiv.org/pdf/1910.00707
19