Star Formation Newsletter#329 36-42
Yusuke Aso (KASI)

pressureradial(Flocketisal.related to the pressure profile by2w
observed.
mps (Pinillaetdistribution
al. 2012). The
Recenthigh-resolution
anddirection.
latterthe vanishes. Assuming the limit case fV = fE = 1, this happens ataccumulatedmaterialshould
form
mas2
The
formation
of
a
dust
ring
in
the
disk
is
most
likelydue
thes much
more
with
ALMA
do
indeed
show
that
most
proto2. Dust where
dynamics
incentera ringto
andits
accretionbya 1w(precisely:
relationshipthelocationwherea⇠0.95),
meaning
⇠3w
/r0
r
is
the
of
the
pressuremaximum
andits
width.
0.sdrifts2019).
The
corollary
is
that
every
star
a pressure
bumpin theformation
disk,whateverthecause
havetowardring-likeappearance
structures(ALMA
Partnershipp ofbeyondpressurebump.
instead fV = 1, but fE = 0, theplanetPlanetbygas
pebbleaccretioninringeddisks
Because
of
drag,
the
dust isthe
attracted
towardtheIfpressure2
ssiveplanets,
which
is
not
observed.
p(r)
=
c
⌃
/(
2⇡H),
(6)
for sthisbump (e.g.,etMHD
instabilities,withdust-gas
instaet However,al. 2018; Longet gal.pressure
2018; Dullemondal. maximum,
9).
solutionof[.] = 0 equation occurs for a = 2.55, meaningaradialspeed
thatis theA.
Morbidelli
The
formation
of
a
dust
ring
in
the
disk
is
most
likelyto the2dbesolved
if
the
dust
grows
much
more
ng
bottlenecks
dust-drift
exist ofat multiple
bilities,
the existenceother planets).
As inforDullemondet al.
wlythat(because⇠0.5w /r0 inward dueof thepressure bump (which corresponds to
where
cfor
Assuming,
simplicity,s = H⌦ is the sound speed.
2appearanceofapressure
bump
in
the
disk,
whateverthe
cause
coagulation
models,andsodriftstoward⌦⌧we assume
thatDépartement
radialpressure
isAntipolis,
Gaussian,
Lagrange,
University of profile
Nice –H
CNRS,aObservatoire
la Côtenote
d’Azur, that,
Nice, Franceof disks at= 2).deWe
if the vortensity-driven
torque were also
H寡占成⻑では100/ (2018),
r (flat disk),
one gets⌃the/thisr2 p(r),with
p(r)Sophia
given
by
g (r)
v
(r)
=
(r
r
)
(2)
e-mail:morby@obs-nice.fr
r
0Økmから成⻑しないのでドリフトでペブルを集めて成⻑。
forpressure
bump(e.g.,
MHD
instabilities,
dust-gas
instaescale(Johansenetal.
2019).
However,22namely:rings is still debated. Magnetoxistofin thesedisks.
w 2020(1 r+2 ⌧at)the strongly saturated ( fV ⇠ 0), the equation [.] = 0 would not haveEq.
(1).more
It is slowly
trivial (because
to see that⌃bilities,(r)
is 2020
proportional
g19
Received
March/ Accepted
9 April totheexistenceof other planets).
As in thisDullemond
et al.
would
also
grow
Ø simulations
数auでガスに対してダストがなくなっていない。降着しないペブルが内側で巨⼤惑星になっていない。
l (MHD)show that the2 !surface
todust-drift
a solution:means that
there would be no planet trap and thelocationofthepressuremaximumr,
whereasithasamaximum
0
(r
r
)
(see
Formula44
inDullemondetal.
2018),
whereH
is thepres0(2018),
weassumethattheradialpressure
profileisGaussian,r).
observationsof disks
at function
ution
disks
can
becomeaqradial
wavy
A&A
638,
A1 (2020)
ce Moreover,ofofmultiplanet
would
migrate
out
of
the
ring.
Because this study only
Ø 圧⼒バンプで⽌まるとうい解。リング分布のペブルを検証。
ABSTRACTp(r)= p0 exp
,
(1)
2
2
2
sure
scale-height
of
the
disk,
⌦
is
the
keplerianorbitalfrequency,
namely:that
mm-size
dust
does
exist
in
disks.
2w
at
r
=
r
(1
+
1
+
8w
/r
)/2.
For
w/r
⌧
1,
the
locationof. 2017), leading
multiple
M to 0 the formation of0Context.
0 is expected to be the dominant process for the
gh-resolution
makes
the
planetconfined
Pebble accretion
formationsenseplanets,such as is
the coresof giant in the ring, we assume that
1of massiveifsolid
and
⌧
is
the
dust’s
Stokesnumber
.
!
hat
there
are
impediments
todust-drift
and super-Earths. So far, this process has 2been 2studied under the assumption that dust coagulates and drifts throughout thesmost
and protohence
dust rings
et planetsal.
2019).
the structured
vortensity-driven
corotation
torque is strongly desaturated
theofmaximum
ofthe⌃(Riols
can beapproximatedbyOn
r Mobservations= rother2w/rmany(rr0+)showfull
protoplanetary
disk. maximum
However,thatdisks areinマゼンタの範囲に惑星がいる
rings
that may be due
to pressure maxima,
g center0and
0 . the
where
r
is
of
the
pressure
whand,
its
width.
0
the
dust
also
undergoesturbulent
diffupossibly
due
to
the
presenceofmultiust-gas
instabilities
canalso
generate
dust
rings
preventing
theglobalradial
drift
of
the
dust.
p(r) = p0 by
expa parabolic,(1) so that the planet trap is at a disPartnership( fV ⇠is confined
1, asinjustified
above),
Approximating
⌃of
near
its
maximum
profile,2 paradigm
g (r)Because
gas
drag,
the
attracted
towardthepressure
Aims.
Wedustaim to is
study
how thepebble-accretionchanges if the dustfora ring.is (Youdin & Lithwick
2w
sion.
The
diffusioncoefficient
the
dust
2
Ø
ガウシアンリング圧⼒
illa
et
al.
2012).
Recenthigh-resolution
l. 2019).etOn
the other hand, massive planets
can2
Methods. Our
approach is mostly analytic.
We derived a formula that tanceprovides an upperbound
to the growtha planet as a function
emond
al.
w /r
theofpressure
bump, with between 0.5 and0 from
we
then
find
that
that
⌃(r)
/
1/r
at
r
=
r
+
3w
/r
.
maximum,
witharadial
speed
that
is
1
0
0
2007):
of
time.
We
also
numerically
implemented
theanalytic
formulæ
tocomputethegrowthofaplanetlocatedinatypicalring observedMAdoindeed
show
that
most
protohegas
distribution
andgeneratepressurebumps
Ø ガスドラッグと粘性拡散のバランスでダスト⾯密度。
depending andon thet at multiple
where
r0 isaswith
theas center
pressurew itsactual
width.degree of saturation of the entropyin the DSHARP
survey,wellina putativeof
ringtherescaled
at 5 AU. 3 maximum
The planet’s radialmigrationin disks
arbitrarysurfaceResults.2020;
Planet Type I migration is stopped
in
a
ring,
but
not
necessarily
atits
center.
If
theentropy-driven
corotation torque is
2Yang & Zhu
ng-likePartnershipal. 2018;structures
Weber et(ALMA
al. 2017;
Ddrag,
drivencorotation
torque.
Thus,
is a measure of where the planet
H ⌦⌧Because
of
gas
dust
is
attracted
towardthepressure
desaturated,theplanet isinlocatedina regionwith thelow etdust
density,which
severely
limitsits accretionrate.
If theplanet is instead
density
profileshasbeenstudiedindetailPaardekooper
al.
D
=
,
(3)
d
v
(r)
=
(r
r
)
(2)
18;
Long
et al. 2018;
Dullemondal.nearndez& Baruteau
Thus,
whetherrings
arer 2020).
0themaximum,
ring’scenter, its accretion rate can 2be similar to the one it woulda classic (ringless)disk ofpressure
equivalent dustmaximum,
density.
traphaveisin relative
to thein units of w2 /r0 . The
2 (1 +
2 et. Magnetowith1afeels
speed
that is+isradial⌧ a by(2010,
2011).
They
demonstratedthetorquew⌧)
However, thatthe growth
rateplanetof the planet
limitedthe diffusion of dust in the ring, and the final planet mass is bounded by the total
ttlenecks
fordust-drift
at multipleringforformation
or exista consequenceofmass.
planet
The DSHARP rings are too far from the star to allow the formation
planetsofwithin
the disk’s lifetime.
radialof massive
profilespressure,gas,However,and dust surface densities, as welltheplanet
surface
from
the
disk
that
is:
2couldDlead
where
isturbulent
diffusion
intheassumedto 円盤と⽐べて外へasimilarring rescaled to
5 AU
tothe
the formationof a planet
incorporating thefull
ringgas,mass in(usuallyless than 1/2 My.
tavy
yetfunction
known. Similarly,
for
the
solar
system,there
H⌦⌧(see" Formula 44 in Dullemond2018),
thetopres!
as
the
range
of
possible
locations
Conclusions. et
The al.
existenceof ringswheremay
not beH
an isobstacleplanetformationby pebble-accretion.
However, foraccretion to beof the planet trap and correvrthe(r)resting= position(rhas
rbe0relatively
)
(2)
be
equal
itstoviscosity)
following↵-prescription
3of +effective,
the planet
near which,the ring’s center,and the ring theneeds tobe not too far from the of 指数的に減少
toftherings
dust
from
the
outer
disk
did
not
penetrate2 (12to
0 debated.
se
is
still
Magnetosure
scale-height
of
the
disk,
⌦
is
the
keplerianorbitalfrequency,w⌧)
sponding
valuesofareillustrated
multiple =centralstar. The
formationof planets ain rings can explain the existence of giant planets with core2masses smaller than the so-called in Fig. 1. For completeness,( 2.5leading
1.7b +
0.1a)+ 1.1
fv (⌫,
M)
Shakura& Sunyaev (1973), gives D = ↵H ⌦, where ↵ is the tur1
disk
in
large
quantities,totheobserved
pebbleisolationmass.
) simulations
the surface
2
and ⌧ isthat
the dust’s
Stokes number .
had weassumeda flared disk,
meaning with H/r / r2/7 andetal. 2019). show
bulentparameter.
Observations
of
the
dust
distribution
in
disks
(see
Formula44formationin Dullemondet2018),interactions
where H is the pres!# Keyheisotopic
properties
ofthe
inner
and
outer
solar
syswords.
planets
and undergoes
satellites:– protoplanetary diffudisksal.
– planet-disk
iskscanrings
becomea radialwavyfunction
On
other
hand,
the
dust
also
turbulentb= 0.15–2 (again,4
3 = 3/7, the planet trap would be located ate dust⇠suggestthat↵disk,= 10⌦–10(Pinte et orbitalal. 2016;
Dullemond et al.
surescale-height
of
the
is
the
keplerianfrequency,
als(Kruijeretal.
2017).
Whether
thisseparation
（a,
bはΣ,
Tの指数）
leading
totheformationofmultiple
+
f
(⌫,
,
M)
7.9
,
(7)
E The diffusion coefficient for the dust is (Youdin & Lithwickon the saturation of the entropy-driven torque), ande planets can sion.
1 depending2018).
Nevertheless, the models of pebble accretion developed so
and
⌧
is
the
dust’s
Stokesnumber
.
1.
Introduction
ner
and
outer
disk’s
dust
was
due
to
the
existencence ofbumps
dust rings(Riols et al. 2019).
2007):
theof thegas
surfacedensityfar
are maximum
quite simplistic, because
theyassumea continuous
flux distribution would be at
ssure
The
balancediffusion(whichtends
toan
dispersetheotherthebetween
dust
undergoesturbulent
diffuump
in thecan
disk(Brasser&
MojzsistoThe2020)formationof On
massive(coreshand,of giant planets,
super- also
2 or
2 the
2planetsof pebbles
as if their initialreservoir had
infinite radial extenstabilities
also
generate
dust
rings=However,12/7.
where 0 = [(M/M
is theandnominaltorque⇤ )(r p /Hp )] is⌃notg (ryetp )r&f JupiterZhu 2020;
p ⌦elucidated.
pdust)
Earths)fully
Theclassic
model
offorceolision.
protoplanetarydisks
are thenot infinitely
The
D
the
drag
(whichto &bring
dust wide.
back
sion.
The
diffusioncoefficient
for
the
dust
isoftends(Youdin
Lithwick
Ø
トルク：f
渦、f
エントロピー起源。
(Kruijer
et
al.
2017)
is
under
debate.
Fig.
1.
Radialprofilespressuredensityand disksdust surface density (blue) for the DHSARP ring B77, from Dul
V ,planets
E garchic
On rings
the othermassivecan growth
in a diskof planetesimals
has stardifficulties
produc- (3)median The
observed
sizes (black),(ingas)gasofsurface
Class-Iand (red),Class-II
Dd =planetaccretes
pebbles/dust
at a rate(Mhand,andM
themassesoftheplanetandtherespec⇤ beingherareetmaximum)al. (2018) – see Sect.gives
3. The verticalscaleof eachhas been state
renormalized for illustrative purposes. The dashed thin red curves sh
2
towards
thelifetime,pressure
thedust
theprofile
steady2007):
ing such massive
bodies within
the disk’s
particularlyare100
and and
200theirAU,tangential
respectively,
althoughsomecurve
exceptional1+⌧
paper places
itselfinthefirsthypothesis(ringsstribution
andgeneratepressurebumps
Ø
粘性が⼩さいとf
も⼩さいが磁場や惑星がリングを維持するのでprofilesr
and
1/r,
points
to
the
red
solid
shows
the locations where the gas surface density profile has local
V atoftively and r p the location
the planet,thedistribution:
disk
distancestypical of where
theradial
giant planetsof thehassolar heightsystemprofiles.
or onescanextend todelimit⇠1000
AU(Najita& Bergin2018).
Due at the migration trap, depending on the degree of saturatio
ceofplanetThese
twolocations
the
range
of
possiblepositionsoftheplanetṀ
= C⌃d (r(8)
netformation
takes
place)
andexploreshow
the otherwarm
jupiters
aroundstars(Levison
etcoefficiental. 2010; Morbidellip ), theWeber
et al.
Yang
Zhu
2020;
observationalbiases,real
medians
are certainly
smaller.
D
vortensity-driven
torque.
Theirdistances
relativeto the locationof the pressure maximum are given in terms of the parameter ,
where
D is&
the
turbulent
in
the
gas,
(usuallyto~1とする。
H pf2017;
androtationfrequency⌦2015;).
In
Formula(7),
thea toThere
0assumed
1 is corotation
pdiffusionVystem,
thereetal.
Johansen&Lambrechts
2017;
Johansen&
Bitschevidence
that
theprotoplanetary
disk
of
the
solar
sys2
maximum of the gas surface-density
distribution. The widths w and w in Eqs. (1) and (4) are also graphically represe
Dd =If,(3)n paradigm
inwhetherthepresenceof Oligarchic
rings.
BBB particuCCCdidof theBaruteau2020).
Thus,
rings
are
(ras therlocation)where2
0tem2019).
growthin⌧profile,planetesimaldisks
is also
be
equal
to
its
viscosity)
which,
followingthe↵-prescriptionofischanges
theexponentofthegas
surface-density
approximatednot ⌃exceed80
AUinradialextension (Kretke
et al. at the planet’s location,1+
Øfは断熱で0、等温で定値。トルク０のところに惑星が留まる。
(r
)
is
the
densityof dustnotpenetrateBCdp⌃dplanetesimals(r)a2= ⌃d,0,Given the expected rates of growth and radial
(4)
E
A
larly
inefficient canif the initialare exp
mostly@⇠100 km 22012).
sequenceplanetformation,
what
follows
driftofformationoforaconsequenceofplanetA turcaveat, however, needs to be mentioned. At ⇠ 0, the gas accretion
rate is capped by (assuming ↵ ⌧ ⌧ ⌧ 1):
Shakura &of
SunyaevD=↵H
whereis theas a power-law
type ⌃gin(1973),(r)= ⌃givesr p ) ⌦,
aroundr p↵; simi2w
g (rp )/(randthecoefficientCdepends
ontheaccretion
cross-section of0 2 2 1
size (Fortieret al.
2013),
as predictedby thestreaming
instabilheobserved
dust(Brauer
et
al.
2008;
Birnstiel
et
al.
2010,
2016),
in
disks
d
where
D
is
the
turbulent
diffusion
in
the
gas,
(usually
assumedtonear
theplanet’s
orbit
is
not
pressure
supported,
and
thus
it
has
a
descriptionofhowasecondgenerationofplanBBwn. Similarly,
for
the
solar
system,therew ⌧ CCC
ity
model(Johansen
et
al.
2015;
Simon
et
al.
2017)
and
suggested
of
these
sizes
the
flux
of
pebblesis
expected
to
last
much
less
ofsize-frequency
the dust distribution
distribution
into follow
disks
larly,bulent
b is theparameter.
exponent Observations
ofbythedisk’sof thetemperature
profile,ofwrittenC,horseshoe
streamlines
relative
tovelocity
the planet, eventhe Ṁ
⇥ 2⇡r p ⌃d (r pto)vr (rthe⌃d,0 exp BB@
theplanetandtheForofetif thedust
relativeplanet,ter solar
sysMax = 2p ) = A⌧
2 A
observations
theremainthan
thelifetime
of
the
disk.
instance,Satoal.
(2016)
be
equal
to
its
viscosity)
which,
followingthe↵-prescription
of
from
the
rings
inducedby
theformation
of
the
2↵r
b4
3
with
a
width
w
related
to
the
width
w
of
the
pressurebumpbyt from the as
outer
disk
did
not
penetrateplanet’smass
issmall.
The
half-widthofthehorseshoeregion0d
T (r)= T (rr ) .–10
Coherently
assumption
thatet al.
suggestthat)/(r↵ = 10(Pintewith
et al.our
2016;
Dullemond
2.55

d0planetary mass as aplanet.
an upperon theinthe locations
case of delimit
a ring,thetherangedensityof dustin these twoof possible
positionsof theplanetTo
at obtain
the migration
trap,bounddependingon thedegree of saturation of the22function
with
time,
we
assumethattheplanetalwaysaccretes
uniform.
Thus,theestimate
of
the
dust
accreTherefore,
grows
slow
rate if at↵r, as/⌧w
riven corotation torque. Their distances relative to the location of thepressure initiallymaximumthe
areplanetgiven intermsatofathe
parameterwell ⌧
0
the maximal
rate
(Eq.
(11)).
By
approximating
eondecreased
by a factor
x2s /2w2d ). In Planet
mostof the maximumof theexp(
gas surface-density
distribution.
The widths
and 2waccretion
in1Eqs.
and(4)arealso graphicallyformation
pebbleinringeddisks
d⌧1byorpossible⌧w3 /↵(we(1)
remindthereaderthatの続きit’s anrepresented.
upper bound);r
will be ⇠1, but it can be significantly smaller ifthenrate slows
A. Morbidelli
MringtheMring down ⌧as the mass of the planet approaches
init(smallw)
andtheplanetmassive(large
x
).
⌃
=
=
,
d
d,0
at, however,needs to be mentioned.Département
At ⇠sLagrange,
0, thegas
accretion
ratecappedby(assuming↵ ⌧ ⌧ ⌧ 1): (12)
the
limit
Mis
asymptotically.
3/2Antipolis,
3/2 rde law
ring
University of Nice
– Sophia
CNRS,
Observatoire
Côted’Azur,Nice,
France
↵
(2⇡)
r
w
(2⇡)
0 d
0
0 2 2 1he planetaccrete moreduste-mail:than
the it has
morby@obs-nice.fr
anet’sorbitcannot
is not pressure
supported,
and
thus
BBB
wto⌧ CCCthe
wethe2 ⇥
planet’smass
would
evolveaccordingng can deliverby diffusion
advection.
Received
19 March
2020if/ Accepted
9 April2020 =
BCA ,orseshoestreamlines
relativeand/or
to the planet,eventhefindØ
惑星が得られる質量の最⼤値を計算。
Ṁthat
2⇡r
⌃
(r
)v
(r
)
=
A⌧
⌃
exp
(11) in ringed disks@Maxp d p r p
d,0Morbidelli:
2A.
Planetformation
3.
An
application:
growthofaplanetina
DSHARPdifferential equation:2↵r0
density
profile(Eq.
(4)) is setofby
equilib- regionass
is small.
The
half-width
thethehorseshoeØ 円盤と違って前後からダストが来るが総量は少ない。
ABSTRACTr
0 2 2 1usion
and⇤advection,
wewhencan useofmassthemis small be the dominant
ring
2
6r
)1/2 (H/r) 1/2
theeitherplanet3p (M/M
Context. Pebble accretion is expected to withprocessfortheformationofmassivesolid
planets,
as
A
=
4⇡r
H
⌦
/r
.
We
note
first
factor of 2 at the
BBB suchthe
CCC of giantA⌧ pthe assumption
wthethroughout⌧coresp 0
pinit init
4 orbit. planetsand super-Earths.
So
far,
this
processhas
beenstudied
under
that
dust
coagulates
anddriftsthestdelivery
rate
to
the
planet’s
Using
the
whose
solution
is:
M(t)
=
M
[1
exp(Bt)],withB
CṀ If=observations(Minofexp
A.disk.
Morbidelli:Planetformation
ringeddiskso ⇠2.5RH Ø
when
M/M⇤ & 10 (Masset
et al. 2006b).
Fig. 2. EvoringinM)
right
hand
the
first
whichaccretionis (13)due totheringfull protoplanetary
However,show
thatmanysidedisks arestructured
rings
thatequals
may@be due tosign,pressuremaxima,常に最⼤値を取って進化すると2 Aplanet
3/2
To
make
a
quantitative
calculation
of
in
a
dust
↵
(2⇡)
r
w
2↵r
0
preventing
theglobalradial
drift
of
the
dust.
)tsandexpression
(9),
we
find
that
the
planet’sr
0
0 2 in2 afrom
1
the perturbations of the planet on Aims.
the Wegas,
the nom- fact paradigm
pebbles in t
thatwedust
both sidesof the planet’sorbit
aim to study how the pebble-accretionchangesthe delivered
dust is confinedring.
3 thatifis
ring,considered
the
rings
studiedin
Dullemondetal.
(2018).
BC
1/3 is mostly analytic.A⌧w⌧
Methods. Our approachWe
derivedaformulaprovides
an
upper
bound
to
the
growthofaplanetasafunction
BBBunlikeon radiusin theis:Hill
regimeis [1racc =(⌧/0.1) with
RBH .=
If dueues of ) reinit
to formulæ
diffusion,
inCCCA . locateda diskdominatedby(14)
advection,oftime. We also numerically
implemented
the analytic
to compute
thegrowth of a planet
in a typicalring observed
exp
whose
solutionM(t)= Mringexp(Bt)],@Ø
解はIn
particular,
we
chose
the
ring
denotedB77
inthat
paper,
as
2
3/2
in
the
DSHARP
survey,
as
well
as
in
a
putativering
rescaledat5AU.
↵material
r wthethe dust reachesrthe Hill radius of theplanet without(2⇡)wherecase (black2↵r the entropy-driven
flows
fromcorotation
one torqueside.is For the reasons
Results. Planet Type I migration is stopped
in a0ring,
but not
necessarily at its
center. If0 only
0
1
it is
thedensity,“medianring”
terms
ofplanet
distancedesaturated,
the planet
is located in a region withlow dust
which severely
limits its inaccretion
rate. If theis instead from the star andtion of theantially affectedby
the gasstreamexplained
above,
if
⇠
0,
the
dust
flux
shouldbe
BBB 2 w2 ⌧ CCCof
2
2 computedA⌧3 the distortions
near the ring’s center, itsTherefore,
accretion rate canbeinitiallysimilar to the one
it
would
have
in
a
classic(ringless)diskofequivalent
dust
density.
theplanet
grows
at
a
slow
rate
if
↵r
/⌧w
⌧
width.
The
properties
this
given
in0 Dullemondet al. (from ⇠0.2
2 ringCAHowever,growth rate ofBut
the planet is limitedby the diffusionof dustin the ring, and the offinal planetmass
is boundedby the total
Be =nominal
exp B@ is the2correct. theradius.
(14)
accretion radiusbysetting
=
max(r
,
x
)r
/w
.
For
comparison,themaxiacc
s
0
3/2
ring
mass.
The
DSHARP
rings
are
too
far
from
the
star
to
allow
the
formationofmassiveplanets
withinthedisk’slifetime.
However,
3
2
d
↵ the
r0floww0.1 of2↵r0 a similar
orcould⌧ lead
/↵mal
⌧
1 (we
remindthereaderthat
it’s1/2wanbound);
rincorporating
AU,
Tmass
=without
4.5 AU,
Mdust
= 40
M .
0 = 75d =upperring rescaled to 1
5 AU
to the(2018)
formation
ofare:
a planet
the
ring
in22
less K,thanMy.
the (2⇡)
specificgas becomes
important.
OnlyØM
accretion
rate
in
afull
disk
pressure
bump
and
dust
Eを種にしてDSHARPのメディアンリング。
Conclusions. The existence of rings may not be an obstacle to planet formation by pebble-accretion. However, for accretion to be
then
the
down
ascenter,thededuced
ofdirectlytheplanet2 rate2to beslows
arethe observations. The
the resting
positiontheplanethasThese
relatively
near
the ring’s
andthe ring =needs
tobe 2not
toofromfar fromapproaches
therried alonginitiallynarrowof gaseffective,thesepØaà3
Myrでは⼤きくならない。⼤きくなっていればリングを⾷べ尽くすので観測とは合う。
Therefore,
the stream
planet grows
atarounda slowrate
ifof↵r/⌧w
density⌃properties
wouldbeṀmass4⇡r⌧⌘⌃d,0existenceMaxd,0 , which is a facp⌦
init⌧
central star. The formation of planets0in rings can
explain theof giant planets withcore masses smaller
thanp the so-called
TheA&Aplanetlimit
Masymptotically.
maximalofthepressurebumpw
isestimatedfromtheA. Morbidelli:Planetformationwidth
in ringed
disks638, Athe
coorbital
motionisaccretedbytheplanet.
The
32Ø Σ∝r
‒1.5
pebbleisolation
mass. the
2
2
2
ring
でr0=5
1 or ⌧ /↵ ⌧ 1 (we remindtheauまでスケールすると3reader that it’s an upper bound);tor Myrで10⇠1/ exp[MEwになる。
⌧/(2↵r0 )] larger than the one in Eq. (11).
Ifthep is w = 17 AU. The radio
Key words. function
planets and satellites:
formation –distanceprotoplanetary disks– planet-disk interactions
betweenadjacent rings and
hhen
W(M,
of slows
this stream
complicatedofapproaches
solutionM(t)the= M planet[1 exp(Bt)],withthe ⌧)
rate
down isasawhosethemassis:of
Formula(11)
hasamaximum
for
=
↵/⌧r
=
2
speed. In t
0 /w, where
Max den5Ṁauにスケールrp0described
1
w
/w
sets
↵/⌧
=
7.7
⇥
10
.
Educatedguessesonthegas
sheandtheparticles’
Stokesnumber,
and
is
init
d
B
C
⌧
w ⌧3.
limit MØasymptotically.B = 1.AIntroduction
ameaning
planetina developed
DSHARPB
CCC An application:
DSHARP
Nevertheless,the models
of pebble accretionso
A ↵⌧r/w⌃d,0growth
exp(
1/2),
it is
proportional
exp BB@
2 ofring トルクはオパシティや磁場に敏感に依る。
A.
ra & Kurokawa(2020). On the(2⇡)
other
the
↵
sity(14)0(⌃
=
10
g
cm
)
and
dust
size
then
lead
to ⌧ = to3 ⇥the10 3 cyan curver w hand,2↵r shear
far
are
quite
simplistic,
becausethey
assumeacontinuous
flux
g
root
the
diffusioncoefficientin thedisk.
The formation of massive planetsring(cores ofsquare
giant planets,
super-of of
pebbles
as4 if their initialreservoir had an infinite
radialextenplanet in a
Therefore,
thes ⌦,
planetgrows
at a slow
rateclassic
if ↵r /⌧w
⌧ of olithis dust relative to the planet
isinitially
3/2x
which
is
Earths)is notyet
fully
elucidated.
The
model
and
↵
=
2.3
⇥
10
Even
if
there
large
uncertainties
on
some
sion.
However,. protoplanetarydisks areare
not infinitely
wide.
The
Another
to consider
pebbleaccretion
in aFig.
ring2. Evolution of the
1 or ⌧ /↵⌧ 1growth(we remind
theofreaderthat it’s anhas
upperbound);
garchic
in a disk
planetesimals
difficulties
produc- aspect
バンプがないとpebble-isolation質量に達する。。
median observed
sizes (in gas)aboutof Class-I
and Class-II
disks
3.3/2rAnacc
application:
growthofaplanetina
DSHARP⌦.ØUsing
numerical
simulations,Kuwahara
then
the
rate
slows
down
as
the
mass
of
the
planetapproaches
oftheseparameters,
wethem
as andfiducial
for the because x s
To themake
athat
quantitative
ofmassplanetaccretionina numbers
dust
ing such massive bodies within
disk’sis
lifetime,
particularly
arecalculation
100 andof
200solid
AU, adopted
respectively,
althoughsome
exceptionalthereservoirisfinite
equal
to
the
dust
limitatMdistances
asymptotically.
pebbles in the
B77we
DHA
of the giant
planets of the solar system or ones can extend to ⇠1000 AU (Najita & Bergin 2018). Due
これは>20
MtheEdensity
で⽊星コアより３倍は⼩さい。
awa
showed that,
if the dust
istypical
radiallyunicases
c
ring
ring,
weetmassconsidered
the
rings
studiedin
Dullemondetal.
(2018).
sakeofexample.
We
also
assumed,forsimplicity,thatthecentralwarm jupiters around other stars(Levisonal. 2010;
Morbidelli
to observationalbiases,the real
medians
are certainly
smaller.
of
the
ring.
As
the
planetis
growingbyaccretingdust,
theues of ) relative
to the
リングによって⼩さいコアを実現可能。
et al.
2015; other
Johansen&ofLambrechts
2017;
Johansen & Bitsch There is evidence that the protoplanetary disk of the solar systwo effects approximately
cancel
each
out,
and
maximum
3. An application:
growtha planetin a DSHARP
initinitiallyIn
particular,
we
chose
the
ring
denotedB77
inthat
paper,
as
star
has
the
mass
of
the
Sun,
even
though
the
real
star
is
0.8
M
.
Fig.
2.
Evolution
of
the
mass
of
a
planetof
0.1
M
accretingmasstemindidthe
Mring=M
M, where
Oligarchic
growthin planetesimaldisks
is dustalso particucase (black curve) wher
not ring
exceed decreases
80 AU in radial asextension(Kretke
et al.
To
make arate
quantitativeofplanetaccretion
intotala mostly
dust
ring
accretion
onto the calculation
planetringis2019).
almostasthe
Totrapadd
larly
inefficientthe
if the same
initialit
planetesimals
are
⇠100
km
2012).
Given
theexpected
rates
ofgrowthandradialdrift
ofis“median
ring”
in
terms
of
distance
from
the
star
and
pebbles
inofthe
B77
ring,
for
different
locations(i.e.,
valTheseparameters
were
used
to
draw
1,
illustrative
Fig.three
3.Fig.
Same asFig. for2, but rescaling
the geometryof theB77 DHSARP
init
tion
ofdensitytheplanetFig.
2. Evolution
themassof alsoa DHARPplanetofinitially
0.1
M
accretingin asize(Fortier
etcalculation
al. 2013), asofpredictedby thestreaming
instabilTo
makequantitative
planetaccretion
in
a
dust
dust
(Brauer
et
al.
2008;
Birnstiel
et
al.
2010,
2016),
in
disks
ing,
we
considered
the
rings
studiedin
Dullemondetal.
(2018).
Mistheinitial
dust
mass
in
the
ring,
and
M
is
the
mass
of
ring
at
5
AU
and
increasing
thedust
surface/
1/r.
ted, a priori incorrectly, using
theconsiderednominalaccretion
pebblesin the
B77
DHARP
ring,
forthethree different(i.e.,
val-pressure
mass,
ues
of
) thisrelative
positionofthemaximum
welltoas⇠0.4).
a themodel (Johansen
al. 2015;
et al.
2017)
and suggested
width.
The
properties
ofgiven
inlocations
Dullemondring, weity
the ringsetstudied
inSimonDullemond
et ringal.
(2018).
ofrelative
these
sizes
thering
fluxtoofthe
pebbles
is maximum
expected
towelllast
muchless et al.
purposes.
(fromas⇠0.2
ues
of
)
to
the
positionofpressureasasanparticular,
we
chose
the
ring
denotedB77
inthat
paper,
asbyobservationsofthesize-frequency
distributionoftheremainIn
particular,
we
chose
the
ring
denotedB77
inthat
paper,
asthanthelifetime
of
the
disk.
For
instance,
Sato
et
al.
(2016)
theplanet.
Toobtain
an
upper
bound
on
the
planetary
mass
asawever, in the case of a ring, the density of dust
in are: r = 75case(blackcurve)
theplanetfollows(mass-dependent)
locacase(blackcurve)
theplanetfollows
theshow
saturation
end
Sect. 2.
theof saturat(2018)
AU,
T=where22
K,
wwhere
= the
4.5
AU,
Mand
=the40M(mass-dependent). effect discussed at the locainit
ring

3

3/2

3

0

2 2

2
0

2
0

2

2

init
ring

3/2

arXiv:2004.07586v1 [astro-ph.GA] 16 Apr 2020

Publications of the Astronomical Society of Japan, (2018), Vol. 00, No. 0

45 m, COHiroyuki N AKANISHI1 , Shinji F UJITA,2 , Kengo TACHIHARA,2 , Natsuko
I ZUMI3 , Mitsuhiro M ATSUO4 , Tomofumi U MEMOTO4 , Yumiko OASA5 andPublications of the Astronomical Societyof Japan, (2018), Vol. 00, No. 05Tsuyoshi I NOUE2
without1 HI gas, referred to as naked H2 clouds, were not
1x107Kagoshima University, 1-21-35 Korimoto
Graduate School of Science and Engineering,
Mass of all HI gas
considered in our study.
Mass of HI clouds

KagoshimaØ ⽔素分⼦⽐は圧⼒、紫外線（光解離）、⾦属量（シールド）の情報。
Department of Astrophysics, Nagoya University, Furo-cho, Chikusa-ku, Nagoya, Aichi
5 Properties
of identified
464-8602,
Japan HI clouds
Ø 5.1コアの外殻でHI,
CI,HI gas
CIIがシールドと思われていたが、CI,
CIIは粒々VLA, HIMass
fraction ofof Science,
HI clouds toIbaraki
the totalUniversity,
CollegeBunkyo 2-1-1, Mito, Ibaraki, 310-8512, Japan
Nobeyama
Radio
Observatory,
National
Astronomical
ObservatoryofJapan,
National
While
analysiswiththedendrogram
algorithm
canidenという過去観測。⼀般にHI, H2雲はどうなのか。
tify clumpy
HI clouds,
cannot recover
all the diffuse
Institutes
ofitNatural
Sciences,462-2HINobeyama Minamimaki, Minamisaku, Nagano
by definition as mentioned
in the section 3.
Ø component
Four-beam
45
mとVLAで銀河⾯サーベイ。
384-1305, Japan
Therefore, we first discuss the fraction of HI mass recovered
Facultyof Education,
Saitama
University, 255 Shimo-Okubo, Sakura, Saitama, Saitama
byour analysisHI,
to the total
estimatedby integrating2 K, missing fluxのためにGreen Bank 100 m.
VLA:
1',mass1.56
km/s,
338-8570,
Japan
the original HI spectra.
Figurem:
2 shows the total20",
HI mass, which
estimated
45
1.3waskm/s,
0.2 K@1'.
E-mail:CO,
hnakanis@sci.kagoshima-u.ac.jp
by integrating the original HI spectra, and the HI cloud
which is estimated by summing
up masses of HIØ mass,VLAのV<‒20
km/sへデンドログラムでHI雲同定。閾値：5
K, 125clouds identified by the dendrogram algorithm, as a functionof the Galactocentric distance. The total HI mass wasvoxel。全ガスの20%が分⼦雲。
Abstract
HI mass within 20◦< l < 50◦ [M"]2

5s estimatedhe HI cloudasses of HI, as a funcHI mass was

1000

100

10

∗

Received 2019 December 19; Accepted 2020 March 17
5

10

12
14
16
Galactocentric Distance [kpc]18

20

Publications of the Astronomical Society of Japan, (2018), Vol. 00, No. 0

Fig. 2. Radial distributions of the total HI mass (black line) and HI cloud mass(gray line). The former was calculated by integrating the original HI spectra
in a Galactic longitude range of 20◦ ≤ l ≤ 50◦ . The latter was calculated bysumming up masses of HI clouds listed in table 1.

1x107

280

0.35

7
Fig. 1. Longitude-velocity diagrams of 12 CO(J = 1–0) (top)HI (bottom) data obtained from FUGIN and VGPS, respectively.
The data were integrated inline.
1x10and
Theoretical
the Galactic latitude ranges of |b| ≤ 1.0◦ for CO and |b| ≤ 1.3◦ for HI. This study focuses on theregions Curveof |b| ≤ 1.0◦ and VLSR First,≤ −20 kms−1 .
it should
be mention

雲が少ない。
àØ 最⼩雲質量14 M

HI Volume Density [cm-3]! b2we! l2 !analyze
Number Density
V2
Mass ofall HI gas
In this study,
molecular
gas formation in neutral Cloudatomichydrogen (HI) clouds
HI Volume Densityof HI clouds
MHI = XHI D2 (∆D)mp
TBMass(VLSR)dVLSRdldb,(3)
240
6
using
the
latestCOdataobtainedfromthefour-beamreceiversystemon a0.345-m tele1x10
b1
l1
V1
◦
◦
◦
◦
scope
(FOREST)
unbiased
Galacticplane
imagingsurveywith
theNobeyama
45-m
200
0.25 telescope
where b1 = −1 , b2 = +1 , l1 = 20 , l2 = 50 , D1 =100000D − ∆D/2,
D2 = Dand+ ∆D/2,V1 = taken
VLSR (D1from), andthe
V2 =Very Large Array (VLA) Galactic plane survey (VGPS).
(FUGIN)HI data160
0.2
VLSR (DWe
2 ).
applied a dendrogram algorithm to the HI data cube to identify HI clouds, and we calcu10000
The total mass of all HI clouds identified by the dendro0.15
lated the HI mass and molecular gas mass 120
by summing the CO line intensity withineach HIgram1000
listed in table 1 amounts to 3.2 × 106 M" while total
cloud.
On theof◦ theresults,
we createda catalog of 5,737 identified HI 0.1clouds with◦80
HI massin the rangesof 20basis≤ l ≤ 50
and R
=9–19 kpc
of rest the(LSR)velocityof VLSR 40≤ −20 km s−1 in Galactic longitude 0.05
and latitude
measured
1.6 ×standard
107 M" . Therefore,
fractionof the HI
100local◦ dendrogramemissionrecoveredanalysisrangesofby20the≤ l ≤ 50◦ and−1◦was≤ 20b ≤%.1◦ , respectively. We found that most of the HI clouds are0
10 fraction varies with the Galactocentric distance
The local
14 which
16 are 18
10
12
16
18
20
distributed
within 14a Galactocentric
distanceof0 1610kpc, 12
mostof
in the 20Cold
NeutralGalactocentric Distance [kpc]GalactocentricDistanceR, being as high as 23%
around R= 12 [kpc]kpc and decreasMedium (CNM) phase. In addition, we determined that the high-mass end of the mass HI funcing below 20% beyond R = 13 kpc. This variation seems
Fig. 3. Radial distributions of HI cloud number density (solid histogram) and
Fig. 2. Radial distributions of the total HI mass (black line) and HI cloud mass
tiontheformer
isphasewellfitted withthe power-law
functionwith an index of 2.3. Although two sequences ofrelatedtheoriginalcold
neutral
HI volume density (dashed line) in a Galactic longitude range of 20◦ ≤ l ≤(gray line).toThe
was transition
calculated by between
integrating theHIspectrasun◦◦
medium
(CNM)
the
warm
neutralmedium50◦ . The HIto
cloud
number density
was calculatedon H
the2 basisof information
self-gravitating
anddiffusecloudsareasbyexpected
appearin theMtot -Mdiagramaccordingin a Galactic
longitudeand
range
of 20
≤l≤
50
. The
latterwas(WNM)
calculated
given in table 1. The HI volume density was calculated by using the original
summing up inmassesof HI clouds
listed in table 1.
discussed
thenext
subsection.
to previous works based on a plane-parallelmodel, the observational data show only a sinspectra of l = 20◦ –50◦ , |b| ≤ 0.5◦ . The error bars were derived according
Cloud Number Density [kpc-2]

m can idene diffuse HIe section 3.
ss recovered
integrating

10000

6

1x10
K
of the HI clouds was 14.0 M! ,00
75
tectable mass was 8.1 M! . T
Galactocentric distance; Column (7) shows100000
the major andfully identified. T=
mass
of WNM withina sphereminor axis radii; Column (8) shows the position angle;
Similarly, because the minimum
heliocentric
distancewhich
is calculated
as la = cs t
Column (9) shows the velocity dispersion;10000
Column (10)of the HI clouds was 11.9 kpc, themassof the faintest
timescale
et ashows the HI mass; and Column (11) shoes the H2 mass.
10.5kpc cloud was estimated to coolingdetectable
be as small
as 8.1tcM(Inoue!
1000
12.5kpc
because the mass of WNM wit
A complete list of these parameters is available online as
The maximum, minimum, and mean values of theis estimated as M = mHI n(4π/14.5kpc
digital data.
100
mass of H2 gas were 1.37 × 104 M! , 0.300 M! , andan HI density of n = 1 cm−3 .
16.5kpc
2
The maximum and minimum values of the HI clouds
2.84 × 10 M! , respectively. The H2 gas was detected1018.5kpc
4
does
not match CNM mass w
mass were 1.35 × 10 M! and 14.0 M! , respectively. The
among HI clouds in the heliocentric distance range of 11.9–
20.5kpc
an acoustic scale of ∼ 0.1 pc omean mass was calculated to be 5.59 × 102 M!1. The near25.1
heliocentric
distance was 16.4 kpc,
0.001
0.01
0.1 kpc.
1 The10mean 100
1000
well as a WNM mass with
est and farthest HI clouds were located at the heliocentric HIwhich
Volume Density[cm-3] the same as that as
is almostof HI clouds. −1
The corlength of ∼ 10
pc. Therefo
distances of 11.9 kpc and 25.1 kpc, respectively. The mean
responding minimum, maximum, and mean values of theFig. 4. Density-pressure diagram. The thick curve is a theoretical model ofto be consistent with the scena
heliocentric distance was 16.4 kpc, which can be regarded
Galactocentric
distanceswere 9.1 kpc, 18.2 kpc, and 11.3
the thermal equilibrium calculated by using
the heating andcooling functions
shocked WNM owing to thermas the typical distance of the HIpresented
clouds byinInoue
our etsample.
al. (2006). The
dashedline shown in
the lowerthe
HI RMS noise in the convolved
kpc,
respectively.
Because−3
by Koyama, & Inutsuka (2002)
12
volume density
range ofn < 3 H cm
corresponds
to
the
WNM
phase,
The corresponding minimum, maximum,and meanvaluesCO data, velocity resolution, and the maximum helioand9.1
thekpc,higher18.2
range
of nand> 10 H cm−3 corresponds to the CNM phase.
The HI mass function plotof the Galactocentric distances werekpc,
centric
distance
of
HI
cloud
with
detected CO emission
The mean HI volume densities of each Galactocentric distance derived with
figure−1 6 can be well fitted with
11.3 kpc, respectively.
were, and Dmax = 25.1
CO = 0.2 K, ∆v = 0.824 km s
original spectra of b = 0 are plotted as
filled ∆Tsquares.
−αConsidering that minimum voxel number of 125 andkpc, respectively, HI clouds with molecular
N = N0 Mmass, less than√π 2 2mp18
minimum contour level of Tmin = 5 K were adopted as de125X) M# = 26 M! could
M = theCO Tminlow mean densities even thoughnumberis ∆v(Dsmall.maxA3600 180
in the mass range of M ≥ 100
scribed in section 3 and that the maximum heliocentric disnot
be counted
completely.
Similarly to the case of HI, thepossible interpretation is thatthese
are mixturesof CNMThe factor N0 and power α w
tance was Dmax = 25.1 kpc, HI clouds with mass less than
mass
of
the
faintest
detectableH
2 gas was estimated toand WNM-phase gases though such clouds are rare.
109.5±0.4 and α = 2.3 ± 0.1, r
18
π 2 mpØ 雲数と密度が⽐例。
ßclumpy-CNM,
diffuse-WNM.

はCNMではなくWNMの熱不安定（cstl内の質量）。

Thermal Pressure: p/kB [K cm-3]

4

estimated with

HI mass within 20◦< l < 50◦ [M"]s, were not

1x106

1000003018), Vol. 00, No. 0

HI gas4

FOREST unbiased Galactic plane imaging
survey with the Nobeyama 45 m telescope(FUGIN). VII. molecular fraction of HI clouds

WNM|CNMØ 内側の>~1000 K cm‒3は外的擾乱でCNMになって雲を作れる。

ds, which
ds in eachouds withed, as exHI clouds
ith a grayの続き

Hiroyuki N AKANISHI1 , Shinji F UJITA,2 , Kengo TACHIHARA,2 , Natsuko
I ZUMI3 , Mitsuhiro M ATSUO4 , Tomofumi U MEMOTO4 , Yumiko OASA5 andTsuyoshi I NOUE2
1

Graduate School of Science and Engineering,Kagoshima
1-21-35
Publications of the
AstronomicalUniversity,
Society of Japan,(2018), Korimoto
Vol. 00, No. 0
Kagoshima2
Department of Astrophysics, Nagoya University, Furo-cho, Chikusa-ku, Nagoya, Aichi
464-8602, Japan
3
College of Science, Ibaraki University, Bunkyo 2-1-1, Mito, Ibaraki, 310-8512, Japan
4
Nobeyama Radio Observatory, National Astronomical Observatory of Japan, National
Institutes of Natural Sciences, 462-2 Nobeyama Minamimaki, Minamisaku, Nagano384-1305, Japan
5
Faculty of Education, Saitama University, 255 Shimo-Okubo, Sakura, Saitama, Saitama
338-8570, Japan

Fig. 6. The number of the identified HI clouds in each mass interval(∆ log10 M) of 0.1. Considering that HI clouds with mass less than 36 M"
were not fully identified, only the mass function for M > 36 M" meets thecompleteness and is plotted with a black thick line. The number for the otheris incomplete and is plotted with a gray line. The line denotes the power-law
function of d(logdN M ) = 109.5 M −2.3 = (3.2 × 109 )M −2.3 .
10

1000

These im
scone
mode
100
at
vit
a
r
dominants
lf-g
seand cold HI
10
self-gravitat
CNM phase1of the CNM2tot
clouds are d
10
100
1000
10000
A schematic
Mtot [M!]Publications of the Astronomical Society of Japan, (2018), Vol. 00, No. 0
scribed abov

Received 2019 December 19; Accepted 2020 March 17

Abstract

diffuse cloud

70
0) 2

∗

d

tot /

Ø Diffuseとself-gravitatingが提唱されているが、今回はその間。à混合物。
E-mail: hnakanis@sci.kagoshima-u.ac.jp
Ø 3D MHD計算でMtot∝t0.9, MH2∝t1.8àMH2∝M .
全体として右の⽮印のような進化。

Fig. 8. Possible paths of molecular cloud evolution superposed on the image

In this study, we analyze molecular gas formation in neutral atomic hydrogen (HI)shown
cloudsin figure 7. The dashed lines labeled (Mtot /7)2 , (Mtot /70)2 , and(Mtot/700)2 denote paths predicted by Inoue & Inutsuka (2012).
using the latest CO data obtained from the four-beam receiver system on a 45-m
telescope (FOREST) unbiased Galactic plane imaging
2 survey with the Nobeyama 45-m telescope(FUGIN) and HI data taken from the Very Large Array (VLA) Galactic plane survey (VGPS).
We applied a dendrogram algorithm to the HI data cube to identify HI clouds, and we6.2
calcuHI and CO distributions of the four most massive
lated the HI mass and molecular gas mass by summing the CO line intensity within eachHIHIclouds
cloud. On the basis of the results, we created a catalog of 5,737 identified HI clouds with
Figure 9 shows the HI and H2 distributions in the four most
local standard of rest (LSR) velocity of VLSR ≤ −20 km s−1 in Galactic longitude and latitude
massive HI clouds, which are presented by gray images andranges of 20◦ ≤ l ≤ 50◦ and −1◦ ≤ b ≤ 1◦ , respectively. We found that most of the HI clouds arecontours, respectively. The thick dashed lines indicate the
distributed within a Galactocentric distance of 16 kpc, most of which are in the Cold Neutral
edges of the HI clouds identified by using the dendrogram.
Medium (CNM) phase. In addition, we determined that the high-mass end of the mass HI function is well fitted with the power-law function with an index of 2.3. Although two sequences ofself-gravitating and diffuse clouds are expected to appear in the Mtot -MH2 diagram according
Fig. 9. Distributions
HI gas (gray scale)and Hbased
the four most massivemodel,
HI clouds. The
levels are 1, 2, 3,...17
s−1 . The
whitea sintoofprevious
works
on ainplane-parallel
thecontour
observational
dataK kmshow
only
2 gas (contours)Ø HIガスは球状ではなく、分⼦ガスも点在。特に端は分⼦雲衝突など
外的ラム圧。
Ø ジャムドーナッツではなくチョコ
チップスコーン。

As show
the most t
chocolate-chmore molecu
as the centre
the cloud prrienced ram
gas in a trig

lou

c
ing

(M

10000

MH2/Mtot~17%

70 2
)

1000
Mass [M!]

10000

tot /

100

10000

Fig. 7. H2 mass versus total mass including HI and H2 gas masses. The
1.06
1.06
power-law function MH2 = 10−0.75 Mtot
= 0.17Mtot, obtained by leastsquares fitting, is plotted with a gray dashed line. The curves of the theoretical model based on Elmegreen (1993) are also plotted and labeled as
self-gravitating cloud and diffuse cloud.

(M

10

1000
Mtot[M!]

7) 2

10

100

ditions beca
the observat
quires modifi
molecular g
range of theclouds, as dular cloud li
gravitating

2

100

diffuse cloud10Ø 指数 ‒2.3
Ø GMCの‒0.5より軽い側に集中。

1000

1

1

tot /

in the mass range of M ≥ 1000 M! as shown in figure 6.
The factor N0 and power α were estimated to be N0 =109.5±0.4 and α = 2.3 ± 0.1, respectively. The resultant
power of α = 2.3±0.1 is much larger than that of the Giant
Molecular Cloud (GMC), which is known to be typically
α = 0.5 in the Galaxy (Solomon et al. 1987). Because
the HI clouds are much smaller in mass than GMCs, it
is understandable that they are very different. This steep
function seems to imply that most of the atomic mass iswithin small CNM clouds.
10

FOREST unbiased Galactic plane imaging
survey with the Nobeyama 45 m telescope(FUGIN). VII. molecular fraction of HI clouds

(M

(5)

MH [M!]

tical line deeshold value
d without CO
ctively.

N = N0 M −α ,10arXiv:2004.07586v1 [astro-ph.GA] 16 Apr 2020

small. Aof CNMare.

s

by Koyama, & Inutsuka (2002).
The HI mass function plotted with the black line infigure 6 can be well fitted with power-law function:

Nunmberthe lower HI
WNM phase,
CNM phase.
derived with

H コントアHIグレー

thick dashed lines denote the boundaries of HI clouds identified by using the dendrogram. The identification numbers correspond to those of the clouds listed
Fig. 10. Schematic diagram of the chocolate-chip scone model. The blackof self-gravi
diffuse clou
6.3 Possib
sistent withthe four mo
In addition
Considerservational d
as implied isystematic elar clouds fi
H2 conversio
collecting dure 7. As
self-gravitat
change the aAcknowled
affect our rebers for coto thank amanuscript.

Advance Access publication 2020 April 28

TW Hya: an old protoplanetary disc revived by its planet
Sergei Nayakshin ,1‹ Takashi Tsukagoshi,2 Cassandra Hall,1 Allona Vazan,3,4
Ravit Helled,4 Jack Humphries,1 Farzana Meru ,5,6 Patrick Neunteufel1,7
and Olja Panic 8287

Planet disruption in TW Hydra

1 School of Physics and Astronomy, University of Leicester, University Road, LE1 7RH Leicester, UK

2 Division of Radio Astronomy, National Astronomical Observatory of Japan, Osawa 2-21-1, Mitaka, Tokyo 181-8588, Japan

Institute for Computational Science, Center for Theoretical Astrophysics and Cosmology University of Zurich, CH-8057 Zurich, SwitzerlandØ 若い円盤のサブ構造があって、そんな円盤の質量は惑星形成には不⼗分
Department of Physics, University of Warwick, Gibbet Hill Road, Coventry CV4 7AL, UK
Centre for Exoplanets and Habitability, University of Warwick, Gibbet Hill Road, Coventry CV4 7AL, UK
à<1 Myrでの惑星形成。
Max-Planck-Institut fur Astrophysik, Karl-Schwarzschild-Str. 1, D-85748 Garching, Germany
School of Physics and Astronomy, University of Leeds, Leeds LS2 9JT, UKØ a few Myrの円盤にmmダストを観測
vs. ドリフト。 Md ~ 1 MJならば速く移動しない。
Accepted 2020 April 17. Received 2020 April 10; in original form 2020 March 13
Ø TW Hya, 60 pc, 10 Myr, 1.8x10‒9 Msun/yr。1, 24, 41 auにギャップ、しかA B S T R AII平均の100倍。
CT
し惑星候補は外側の崖にある。MdustがClass
4
5
6
7
8

Downloaded from https://academic.oup.com/mnras/article-abstract/495/1/285/5826181 by KASI user on 12 June 2

3 Racah Institute of Physics, The Hebrew University, Jerusalem 91904, Israel

0.5 x 2.2 au
半径x⽅位Dark rings with bright rims are the indirect signposts of planets embedded in protoplanetaryØ 結論：Steady-stateではなく惑星から質量放出。

Ø SS. 冪乗、α粘性、50 auで10 Myrのドリフト。
à Md~0.2 Msun, Mcpd > 10ME << Mp（✖） 。
外の⽅がドリフト時間が短いので元はもっと急（✖） 。
タイプI移動時間は0.04 Myr。惑星100個（✖） 。

３倍のズレ。
Ø Mpなどパラメーターを調整してもダメ 。à
崖で光学的に厚い。
（TW Hya; ∝ν2.6）
Key words: planets and satellites: formation – protoplanetary discs.

substructures are seen in very young discs, and that the masses of

Downloaded from https:/

discs. In a recent first, an azimuthally elongated AU-scale blob, possibly a planet, was resolved
with ALMA in TW Hya. The blob is at the edge of a cliff-like rollover in the dust disc rather
a dark ring. Here we build time-dependent models of TW Hya disc. We find that292than inside
S. Nayakshin
et al.
the classical paradigm cannot account for the morphology of the disc and the blob. We proposethat ALMA-discovered blob hides a Neptune mass planet losing gas and dust. We show thatradial drift of mm-sized dust particles naturally explains why the blob is located on the edgeof the dust disc. Dust particles leaving the planet perform a characteristic U-turn relative toit, producing an azimuthally elongated blob-like emission feature. This scenario also explains
why a 10 Myr old disc is so bright in dust continuum. Two scenarios for the dust-losing planet
are presented. In the first, a dusty pre-runaway gas envelope of a ∼ 40 M⊕ Core Accretion
planet is disrupted, e.g. as a result of a catastrophic encounter. In the second, a massive dusty
pre-collapse gas giant planet formed by Gravitational Instability is disrupted by the energy
released in its massive core. Future modelling may discriminate between these scenarios andallow us to study planet formation in an entirely new way – by analysing the flows of dust andgas recently belonging to planets, informing us about the structure of pre-disruption planetary
envelopes.

Advance Access publication 2020 April 28

TW Hya: an old protoplanetary disc revived by its planet
Sergei Nayakshin ,1‹ Takashi Tsukagoshi,2 Cassandra Hall,1 Allona Vazan,3,4
の続き
4
1
5,6
1,7
Ravit Helled, Jack Humphries, Farzana Meru , Patrick Neunteufel
Planet disruption in TW Hydra
and Olja Panic 8295

1 School of Physics and Astronomy, University of Leicester, University Road, LE1 7RH Leicester, UK

gas disc mass of a factor of 50 between the two models shows that
there is a certain degeneracy in the model parameters, e.g. a higher
3 Racah Institute of Physics, The Hebrew University, Jerusalem 91904,
value
of α could be compensated for by a higher v br .
Israel2 Division of Radio Astronomy, National Astronomical Observatory of Japan, Osawa 2-21-1, Mitaka, Tokyo 181-8588, Japan
4
5

8 G I P L A N E T D I S RU P T I O N

67

8.1 Physical motivation and constraints

8 School of Physics and Astronomy, University of Leeds, Leeds LS2 9JT, UKØ 惑星は質量も軌道も⼀定なので物理を追加。

296

S. Nayakshin et al.

ABSTRACTØ CA: 崩壊前Mp~20 MEはケプラー速度程度で壊れる。
Dark rings with bright rims are the indirect signposts of planets embedded in protoplanetary
吹き⾶ぶdiscs. In a recent first, an azimuthally elongated AU-scale blob, possibly
a planet, was resolved
落ちないようにMdを下げて𝑀̇ を変えないようにαを上げた。
with ALMA in TW Hya. The blob is at the edge of a cliff-like rollover in the dust disc rather
8.1.1 The need for a very rapid primordial disc dissipation

298

Downloaded from https://academic.oup.com/mnras/article-abstract/495/1/285/5826181 by KASI user on 12 June 2020

In Section 7.1 we argued that a core accretion planet may losea major fraction of its gas-dust envelope if two conditions are
satisfied: (i) the envelope is in the extended, pre-collapse state; (ii)a significant energy is injected in it, e.g. via collision with anotherAccepted 2020 April 17. Received 2020 April 10; in original form 2020
March
massivecore. 13
We now detail conditions under which a GI planet
disruption could be relevant to TW Hydra’s observations.

Downloaded from https://academic.oup.com/mnras/article-abstract/495/1/285/5826181 by KASI user on 12 June 2

Institute for Computational
Science, Center for Theoretical Astrophysics and Cosmology University of Zurich, CH-8057 Zurich, Switzerland‒6 M /yr。
Ø リング状放出。𝑀̇ Z=6x10
Department of Physics,
Universityof Warwick, Gibbet Hill Road, Coventry CV4 7AL, UK
E
Centre for Exoplanets and Habitability, University of Warwick, Gibbet Hill Road, Coventry CV4 7AL, UK
CA
Max-Planck-Institut fur Astrophysik, Karl-Schwarzschild-Str. 1, D-85748 Garching, GermanyØ 成⻑すると内側へ落下à⾃然に崖。光学的にも薄い。

S. Nayakshin et al.

8.1.3 GI planet disruption by a coreGI

Ø GI: 早く惑星ができるので円盤が> MJだと落ちてしまう。
最近のオパシティを使うと崩壊まで5-10 Myrかかる。
崩壊前のコアが>10 MEだとエンベロープが吹き⾶ぶ。
why a 10 Myr old disc is so bright in dust continuum. Two scenarios for the dust-losing planetØ 惑星からの質量放出によって⽅位⾓⽅向の伸びも説明。
ß放出されたダストがガスの速度になって離れる。

Key words: planets and satellites: formation – protoplanetary discs.

惑星軌道↙

Downloaded from https://academ

are presented. In the first, a dusty pre-runaway gas envelope of a ∼ 40 M Core Accretion
planet is disrupted, e.g. as a result of a catastrophic encounter. In the second, a massive dusty
pre-collapse gas giant planet formed by Gravitational Instability is disrupted by the energy
released in its massive core. Future modelling may discriminate between these scenarios andallow us to study planet formation in an entirely new way – by analysing the flows of dust andgas recently belonging to planets, informing us about the structure of pre-disruption planetary
envelopes.

Downloaded from https://academic.oup.com/mnras/article-abstract/495/1/285/58261

A number of authors have shown that pre-collapse GI planets may
develop massive solid cores via grain growth and sedimentation(Kuiper 1951; McCrea & Williams 1965; Boss 1998; Helled &Schubert 2008; Boley et al. 2010; Cha & Nayakshin 2011; Nayakshin 2011; Vorobyov & Elbakyan 2019). The time-scales on which
the core grows are a minimum of thousands of years but may bemuch longer, depending on convection and grain material/growthproperties, such as v b (Helled et al. 2008). If the core grows more
massive than ∼ 10 M⊕ , then the energy release during its formation
can be too large for the pre-collapse planet – its envelope expands
and is eventually lost (Nayakshin & Cha 2012; Nayakshin 2016;
Humphries & Nayakshin 2018). This scenario for the core-initiateddisruption of GI protoplanets is physically analogous to how coresof AGB stars eject their envelopes except for the energy source –the gravitational potential energy rather than the nuclear energy ofPlanet disruption
in TW Hydra299
the core – and the physicalscales of the systems.
2
At present, there exists no stellar/planet evolution code that takes
1 R
−1
4
=
5.4
×
10
yr
α
,
(26)
tvisc =
into accountthe source
relevantmodel
physics that we wish to explore here. For−1
Figure ⊕8. Same as Fig. 4 but now for the coreaccretionalldust
3α"K H 2
example, Helled et al. (2008), Helled & Bodenheimer (2011) present
at t = 6.8 × 104 yr.
where α −1 = (α/0.1) and the estimate is made at R = 51.5 AU.
models of grain growth and sedimentation in pre-disruption isolated
planets cooling radiatively. Vazan & Helled (2012) investigate
Numerical simulations show that the α parameter due to self-gravity
is slow. Further, due to its lower mass the type I migration timehow external irradiation affects contraction of these planets. Theseof the disc may reach values of order ∼0.1 in early massive discs
scale is longer. Therefore, there is no reason
to demanda complete
studies did
not include
the effects of the massive core energy release(Gammie 2001; Rice et al. 2005; Haworth et al. 2020), and evendisappearance of the primordial gas discbeforetheplanetdisruption
on
to
the
planet,which
is central for us here. On the other hand,α ∼ 0.2 in the magnetized discs (Deng et al. 2020). Additionally,
commences in the CA framework. Nayakshin (2015, 2016) include grain growth, sedimentation, core
disc depletion due to external photoevaporation may be faster thanformation, and the effects of the core energy feedback on to thepreviously thought (e.g. Haworth & Clarke 2019).
envelope, but use a simplified follow-adiabats approach to model
For the case at hand we must require that the primordial gas disccooling of the envelope, and a simpler equation of state
8.1.2 Why did the planet not collapseradiativein 10 Myr?
is long gone in TW hydra. This is because the migration time-scalethan Vazan & Helled (2012) do. Further, the opacities used by theof a GI planet with mass initially exceeding 1 MJ would be much
1 a few AU,
GI planets are born extended, with theirrp ∼two radiuscodes aredifferent.
shorter than 10 Myr. Indeed, if the planet did not open a gap andHere we shallthe codeof Nayakshin (2016) to understand theand cool: their central temperature is in hundredsof Kuse(e.g.
Helled,physicalconstraints
the pre-disruption
planet mass, metallicity,
migrated in the type I regime then its migration time is less than
Podolak & Kovetz 2008). If dust growth
insidethe GIonplanetisthe massof the core
responsible for the planet disruption.
1 Myr even for a disc as low mass as a few MJ . If, on the otherneglected, then it cools, contracts, andand
eventuallycollapses
dynamThese
willbe seen to place significant limitations onhand, the planet did open a wide gap and migrated
typeEvolution
II then of gas icallymoleculedissociation
whenradius,
the endothermic
reactionFigureof H2 constrains
Figure 9.inTop:
giant planet
rp , for different
11.
Same
as Fig.
4 but nowfor planet
the GImassmodel
the disrupted
GI planet
scenario(e.g. theis (Section
unlikely 8.3) at t =metallicities
dashed) aandwithout(thinofsolid)
core energy
the migration time-scale is (e.g. Lodato & Clarke2004)of the planet with (thickabsorbs
vastamountthermalof×the(Bodenheimer
6.8
104planetyr.
).
InAppendix B we compare for the first timetoexceed∼2
MJ
formation.
Bottom:Core massdiscs,
as a function
ofThe
time
for
different
metallicities
substructures
are
seen
in
very
young
and
that
the
massesof1974).
collapseterminatesinformationofaplanetthatisMp
the results of uniform planet contraction calculations (no dust

In the Gravitational Instability (GI; Kuiper 1951) theory for planet
formation, massive and very young gaseous discs fragment at
separations ∼50–100 AU on to Jovian mass gas clumps (e.g.
Rafikov 2005; Rice, Lodato & Armitage 2005). Hydrodynamical
simulations show that these planets migrate inward very rapidly
than inside a dark ring.
Here we build time-dependent models of TW Hya disc. We find thatin massive discs (e.g. Vorobyov & Basu 2006; Boley et al.
Baruteau,
Meru &for
Paardekooper2011), perhaps
the classical paradigm2010;
cannotaccount
the morphology
ofexplaining
the disc and the blob. We propose(Humphries et al. 2019) why wide-orbit separation gas giants arethat ALMA-discovered
blob
hides
a
Neptunemass
planetlosingso rare in direct imaging surveys (Vigan et al. 2012; Chauvin et al.gas and dust. We show that2015).
On theparticles
other hand, planet–planet
allowthesomeblob is located on the edge
radial drift of mm-sized
dust
naturally scatterings
explainsmaywhy
GI planets to survive on wide orbits (Vigan et al. 2017), especiallyof the dust disc. Dustif particles
leavingtheplanetperformacharacteristic
U-turn relative tothe primordial disc is dispersed rapidly. For the early massive
protoplanetary
discs, the primarytime-scaleon whichfeature.
its mass isThis scenario also explainsit, producing an azimuthally
elongatedblob-likeemissionlost (Clarke, Gendrin & Sotomayor 2001) is the viscous time,

78712-1205, USA
5 Institute for Astronomy, University of Hawai‘i
2MASS J04435686+3723033
B:
at Mānoa, 2680 Woodlawn
Drive, Honolulu, HI 96822, USA
A Young Companion at the Substellar Boundary with Potential Membership in the Pictoris Moving Group4
ApJ 7thApril2020) C. Liu,5 and Kimberly Sokal4
Caprice L. Phillips,1, 2 Brendan (Accepted
P. Bowler,3by
Gregory
Mace,Michael1 Department of Astronomy, The Ohio State University, Columbus, OH, 43210,USA

ABSTRACT

2 LSSTC DSFP Fellow
3 Department of Astronomy, The University of Texas at Austin, Austin, TX, 78712, USA

We 4present
a detailed characterization of 2MASS J04435750+3723031, a low-mass companion orbitMcDonald Observatory & Department of Astronomy, University of Texas at Austin, 2515 Speedway, Stop C1400, Austin, TX
78712-1205, USA
ing the young M2 star ” 2MASS J04435686+3723033
at 7.00 6 (550 AU) with potential membership in
5
for Astronomy, University of Hawai‘i at Mānoa, 2680 Woodlawn Drive, Honolulu, HI 96822, USA
the 23 Myr InstitutePictorismoving group ( PMG). Using near-infrared spectroscopy of the companion from
IRTF/SpeX we have found a spectral(Acceptedtype of M6
± 7th
1 and
of youth through age-sensitive
by ApJ
Aprilindications
2020)
absorption lines and a low surface gravity index (VL-G). A young age is supported by H↵ emission
ABSTRACT
and lithium absorption in the host. We re-evaluate
the membership of this system and find that it is apresent a detailed
characterizationof 2MASS
J04435750+3723031,
a low-masscompanion
orbit-velocities
marginally We
consistent
kinematic
match to thePMGusing Gaia parallaxes
andnew radial00
ing andthe youngM2 star ”If2MASS
J04435686+3723033
7. 6the(550 PMG,
AU) with
membershipin outlier
for the hostcompanion.
this systemdoes belongatto
it potential
would bea kinematicthe 23 Myr Pictoris moving group ( PMG). Using near-infrared spectroscopy of the companion fromand the companion
would be over-luminous compared to other similar ultracool objects like PZ Tel B;
IRTF/SpeX we have found a spectral type of M6 ± 1 and indications of youth through age-sensitive
this would
suggest lines
2M0443+3723
B could
be index
a closebrown
dwarfagebinary(⇡52+52
MJup
if equal-flux,
absorption
and a low surface
gravity(VL-G).
A young
is supportedby H↵
emissioncomparedandwith
99 ±absorption
5 MJup if
would makeit the sixthofsubstellar
lithiumin single),
the host. andWe re-evaluate
the membership
this system companion
and find that in
it isthisa group.
marginallyconsistent
the images
PMG usingparallaxes
and new
To test thishypothesis,
wekinematic
acquiredmatch
NIRtoAO
withGaia
Keck
II/NIRC2,
butradial
theyvelocities
do not resolve
for the hostanda companion.
If thisdoes belong limit
to the ofPMG,
it would
be a kinematic outlier
the companion
to bebinary downtosystemthe di↵raction⇠3 AU.
If 2M0443+3723
AB does notand the companion would be over-luminous compared to other similar ultracool objects like PZ Tel B;
belong tothis
anywould
movinggroup then its age is more uncertain. In this case it is still young (.30 Myr),suggest 2M0443+3723 B could be a close brown dwarf binary (⇡52+52 MJup if equal-flux,
and the impliedof ±thewould be between ⇠30–110 MJupcompanion.
comparedmass
with 99
5 Mcompanion
in this group.
Jup if single), and would make it the sixth substellar

250v1 [astro-ph.SR] 15 Apr 2020

04.07250v1 [astro-ph.SR] 15 Apr 2020

Department of Astronomy, The University of Texas at Austin, Austin, TX, 78712, USA
4 McDonald Observatory & Department of Astronomy, University of Texas at Austin, 2515 Speedway, Stop C1400, Austin, TX

To test this hypothesis, we acquired NIR AO images with Keck II/NIRC2, but they do not resolve
companion
to be –a stars:
binary down
to thedi↵raction
limit ofindividual⇠3 AU. If 2M0443+3723
AB does not
Keywords:
binaries:
close
brown
dwarfs,imaging,(2MASS J04435750+3723031,
βPic. theMovingGroupの候補が明るすぎるのでBD連星かと思ったが3auで分解できず。
belong to any moving group then its age is more uncertain. In this case it is still young (.30 Myr),
2MASS J04435686+3723033), low-mass
and the implied mass of the companion would be between ⇠30–110 MJup .

1. INTRODUCTION

ally defined based on the limit for the onset of hydrogenLynne A. Hillenbrand3, Selma E. de Mink4,5 , and David R. Soderblom2© 2020. The American Astronomical Society. All rights reserved.

1

Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA; imants@jhu.edu
2
Space Telescope
Telescope Science
Institute, in
3700
MartinNebula
Drive, Baltimore,
21218,ofUSA
Hubble 3Space
Astrometry
theSanOrion
Cluster:MD
CensusLow-mass
Department of Astronomy, CaliforniaRunaways
Institute of Technology, Pasadena, CA 91125, USA
4
Center for Astrophysics, Harvard-Smithsonian, 60 Garden Street, Cambridge, MA 02138, USA
5
1
2
Imants PlataisInstitute, Massimo
Robberto2 , Andrea
Bellini, Vera Kozhurina-Platais, Mario, Giovanni
Strampelli,
Anton Pannekoek
for Astronomy,
University
of2 Amsterdam,
Science Park
904,Gennaro1098 2XH,
Amsterdam,
The2 Netherlands
34,52
Lynne A.21;
Hillenbrand, SelmaApril
E. de15;
Mink, and2020
DavidApril
R. Soderblom
Received2019 Octoberrevised 2020
accepted
24; published 2020 May 21
1
Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA; imants@jhu.edu
2
Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA
3
Department of Astronomy, California Institute of Technology, Pasadena, CA 91125, USA
4
Center for Astrophysics, Harvard-Smithsonian, 60 Garden Street, Cambridge, MA 02138, USA
5
Anton Pannekoek Institute for Astronomy, University of Amsterdam, Science Park 904, 1098 XH, Amsterdam, The Netherlands
Received 2019 October 21; revised 2020 April 15; accepted 2020 April 24; published 2020 May 21

Abstract

We present a catalog of high-precision proper motions in the Orion Nebula Cluster (ONC), based on Treasury
Program observations with the Hubble Space Telescope’s
Abstract(HST) ACS/WFC camera. Our catalog contains 2454
<24.7,
thusNebula
probingstellarmassesobjects in the magnitude
14.2<mproper
F775W
We present a range
catalog ofofhigh-precision
motionsin the Orion
Clusterthe(ONC),
based onTreasuryof the ONC from2
HubbleTelescope’s
ACS/WFC
camera.
Our catalog
2454velocity dispersion∼0.02observations
M☉ overwithanthearea
ofSpace∼550
arcmin(HST). We
providea numberofcontains
internal∼0.4 M☉ down toProgram
objects in the magnitude range of 14.2<mF775W<24.7, thus probing the stellar masses of the ONC fromestimates for the ONC
that
indicateon2.stellar
locationmass.velocity
Thereis good agreement withdown
to ∼0.02 Ma☉weakover an dependence
area of ∼550 arcmin
We provide
a numberandof internal
dispersion∼0.4 M☉estimates for
the ONC that estimates,
indicate a weak althoughdependence onnearlystellar location
Thereis good agreement
the published velocity
dispersion
all andof mass.
them(including
ourswithat Tv, x  0.94 andthe−1published velocity dispersion estimates, although nearly all of them (including ours at Tv, x  0.94 andTv, y  1.25 masyrTv,y ) 1.25
might
be−1biased
thebyoverlapping
young
stellarpopulations
ofidentiﬁed
Orion four
A. We identiﬁed fourmasyr) might beby
biasedthe overlapping youngstellarpopulationsof Orion A. We
total ∼1 M☉. The total
new ONC
candidatebased
runawayson
based
on HST
the Gaia
Gaia DR
2 data,all withallmasseslessmasses
than ∼1 Mless
new ONC candidate
runaways
HSTandandtheDR
2 data,with☉. Thethan
census of known candidate runaway sources is 10—one of the largest samples ever found in any Milky Way open
census of known candidate
runaway sources is 10—one of the largest samples ever foundin any Milky Way open
them indeed
star cluster. Surprisingly, none of them have tangential velocities exceeding 20 km s−1. If most of−1
. If mostof them indeed
star cluster. Surprisingly,
ofitthem
havethetangential
20young
km stars clusters.
originated innone
the ONC,may compel
re-examinationvelocities
of dynamical exceeding
processes in veryItappears
that
the
mass
function
of
the
ONC
is
not
signiﬁcantly
affected
by
the
lost
runaways.
originated in the ONC, it may compel the re-examination of dynamical processes in very young star clusters. It
Uniﬁed Astronomy Thesaurus concepts: Hubble Space Telescope (761); Space astrometry (1541); Proper motions
appears that the massfunction of the ONC is not signiﬁcantly affected by the lost runaways.
(1295); Young star clusters (1833); Low mass stars (2050); Stellar dynamics (1596); Runaway stars (1417)
material:
machine-readable
table Space Telescope (761); Space astrometry (1541); Proper motions
Uniﬁed AstronomySupporting
Thesaurus
concepts:
Hubbleハッブル望遠鏡でONCの固有運動を⾼精度でカタログ。10個のrunawayはどれも20(1295); Young star clusters (1833); Low mass stars (2050); Stellar dynamics (1596); Runawaykm/s超えず。
stars (1417)
1. Introduction

Supporting material: machine-readable table

As one of the most recognizable objects in the astronomicalsky and one of the most popular in studies of star formation, theOrion Nebula Cluster (ONC) may not need a formal
introduction.
Its basic properties are reviewed by, e.g., O’Dell
1. Introduction(2001), O’Dell et al. (2008), and Muench et al. (2008). Overclusters and associations. Among them, the ONC was given
special attention due to its prominent role in our understandingof star formation. These authors found that the ONC is an
ordinary gravitationally bound cluster, while the known
asymmetry in the internal velocity distribution is re-conﬁrmed.
clusters
and wereassociations.
Among
them,
Similarconclusions
reached by Kim
et al. (2019)using the ONC was given
archival ONC images from various HST cameras and from

Konkoly Observatory, Research Centre for Astronomy and Earth Sciences, Konkoly-Thege Miklós út 15-17, 1121 Budapest,
Hungary
15
ELTE Eötvös Loránd University, Instituteof Physics, Pázmány
Péter sétány 1/A, survey
1117 Budapest, Hungary
Exocomets:
A spectroscopic
16
Aurora Technology B.V. for1 ESA, ESA-ESAC,
Villanueva de la Cañada, 28691, Madrid, Spain
I. Rebollido , C. Eiroa1, 2 , B. Montesinos3 , J. Maldonado4 , E. Villaver1 , O. Absil5 , A. Bayo6, 7 , H. Canovas8 , A.
17
Observatorio Astronómico9 Nacional10 (OAN-IGN)-Observatorio
de Madrid,
Alfonso XII,
11
12
6, 7
12 3, 28014
13 Madrid, 1Spain
Carmona , Ch. Chen , S. Ertel , Th. Henning , D. P. Iglesias , R. Launhardt , R. Liseau , G. Meeus , A.
Moór14, 15 , A. Mora16 , J. Olofsson6, 7 , G. Rauw5 , and P. Riviere-Marichalar17
1

Dpto. Física Teórica, Universidad Autónoma de Madrid,
Cantoblanco,28049 Madrid, Spain, e-mail: isabel.rebollido@uam.es
ABSTRACT
Observatorio Astrónomico de Calar Alto, CAHA, 04550 Gérgal, Almería, Spain
3
Centro de Astrobiología (CAB, CSIC-INTA), ESAC Campus Camino Bajo del Castillo, s/n, Villanueva de la Cañada, 28692,
Madrid,Context. While exoplanets
areSpainnow routinely detected, the detection of small bodies in extrasolar systems remains challenging. Since
4
INAF, Osservatorio Astronomico di Palermo, Piazza del Parlamento 1, 90134 Palermo, Italy
the discovery of sporadic
eventsinterpreted
exocomets(Falling
Evaporating
Bodies)around Pic in the early 80s, only ⇠20 stars
5STAR
Institute,
Université deas
Liège,
F.R.S.-FNRS,
19c Allée duSix Août, B-4000
Liège, Belgium
6Institutode Física y Astronomía,
Facultad de Ciencias, Universidad de Valparaíso, 5030 Casilla, Valparaíso, Chile
have been reported to host
exocomet-likeevents.
7
Núcleo Milenio de Formación Planetaria - NPF, Universidad de Valparaíso, Av. Gran Bretaña 1111, Valparaíso, Chile
Aims. We aim to expand
the sample of known exocomet-host stars, as well as to monitor the hot-gas environment around stars with
8
European Space Astronomy Centre (ESA), PO Box, 78, 28691 Villanueva de la Cañada, Madrid, Spain
9previously known exocometary
Université activity.
de Toulouse, UPS-OMP, IRAP, Toulouse F-31400, France
10
Space
Telescope
Science Institute,
3700spectra
San Martinof
Drive,
Baltimore, MD 21212,
USA of 117 main-sequence stars in the spectral type
Methods. We have obtained
high-resolution
opticala heterogeneous
sample11
Steward Observatory, Department of Astronomy, University of Arizona, Tucson, AZ 85721, USA
range from B8 to G8. 12The
data have beencollected
in 14
observing
campaigns
over 2 years from both hemispheres. We
Max-Planck-Institut
für Astronomie(MPIA),
Königstuhl
17, D-69117
Heidelberg,expanding
Germany
13of Space,and Environment,
Universityof Technology, Onsala Space
Observatory, originated439 92 Onsala, in the circumstellar
have analysed the Ca ii Department
K&H and
Na i EarthD linesin order Chalmers
to search
for non-photosphericabsorptions
Sweden14
environment, and for variable
thatResearch
could Centre
be causedby outgassingof exocomet-likebodies.
Konkolyevents
Observatory,
for Astronomy
and Earth Sciences,
Konkoly-Thege Miklós
út 15-17, 1121 Budapest,
Hungary
Results. We have detected
non-photospheric absorptions towards 50% of the sample, attributing a circumstellar origin to half of the15
ELTE Eötvös Loránd University, Institute of Physics, Pázmány Péter sétány 1/A, 1117 Budapest, Hungary
detections (i.e. 26% of16 thesample).
gas
is detectedin themetallic
lines inspected via narrow stable absorptions,
AuroraTechnologyHot
B.V. circumstellar
for ESA, ESA-ESAC,
Villanueva
de la Cañada,28691,
Madrid, Spain
17
Observatorioabsorption
Astronómico Nacional(OAN-IGN)-Observatorio
de Madrid,
Alfonso
XII, 3, in
28014
and/or variable blue-/red-shifted
events.
Such variable events
were
found
18Madrid,stars Spainin the Ca ii and/or Na i lines; 6 of2

084v1 [astro-ph.SR] 24 Mar 2020

arXiv:2003.11084v1 [astro-ph.SR]

14them are reported in the context of this work for the first time. In some cases the variations we report in the Ca ii K line are similar to
those observed in Pic. While we do not find a significant trendwith the age or location of the stars, we do find that the probability
ABSTRACTof finding CS gas in stars with larger v sin i is higher. We also find a weak trend with the presence of near-infrared excess, and with
Context. While exoplanets are now routinely detected, the detection of small bodies in extrasolar systems remains challenging. Since
anomalous ( Boo-like)butevents
thisinterpreted
would requireconfirmation
by expanding
the abundances,discovery of sporadic
as exocomets(Falling Evaporating
Bodies) around the
Picsample.
in the early 80s, only ⇠20 stars
have been reported to host exocomet-like events.

Aims. We
aim to expandsystems
the sample ofknown exocomet-host -stars,
as wellas to monitor
the hot-gas environment
Key words. stars: general- planetary- comets:general
ISM:
clouds- circumstellar
matter around stars with

1. Introduction

Ca II K, Ca II H, Na I D輝線で117個の主系列星で系外彗星を調査。

previously known exocometary activity.
Methods. We have obtained high-resolution optical spectra of a heterogeneous sample of 117 main-sequence stars in the spectral type
range from B8 to G8. The data have been collected in 14 observing campaigns expanding over 2 years from both hemispheres. We
have analysed the Ca ii K&H and Na i D lines in order to search for non-photospheric absorptions originated in the circumstellar
environment, and for variable events that could be caused by outgassing of exocomet-like bodies.

fied architectures (see e.g. The Extrasolar Planets Encyclopae-

1© 2020. The American Astronomical Society. All rights reserved.
Infrared Science Archive(IRSA), IPAC, 1200 E. California Blvd., California Institute of Technology, Pasadena, CA 91125, USA; rebull@ipac.caltech.edu
2
Spitzer Science Center (SSC), (IPAC), 1200 E. California Blvd., California Institute of Technology, Pasadena, CA 9112, USA
3
Bay Area Environmental
Research
Institute, 625
2nd inSt. Taurus
Ste. 209, Petaluma,
Rotationof Low-massStars
with K2CA 94952, USA
4
Astronomy Department, California Institute of Technology, Pasadena, CA 91125, USA
5
1
2
3
5
5
6
L. M.deRebull, J. R. Stauffer, A. M. Codyde, L. A. HillenbrandJ. BouvierGrenoble,, N. Roggero, andCNRS,
T. J. David
Université de Grenoble, Institut
Planétologie
et d’Astrophysique
Grenoble(IPAG),4, F-38000
France;IPAG,
F-38000 Grenoble, France
1
6
Infrared Science
Archive (IRSA), IPAC, 1200 E. California Blvd., California Institute of Technology, Pasadena, CA 91125, USA; rebull@ipac.caltech.edu
2Center for Computational Astrophysics, Flatiron Institute, New York, NY 10010, USA
Spitzer Science Center (SSC), (IPAC), 1200 E. California Blvd., California Institute of Technology, Pasadena, CA 9112, USA
3
Received 2020 February
revised Research2020 April
AprilCA
8;94952,
published2020 May 26
Bay Area24;
Environmental
Institute,7;625accepted
2nd St. Ste. 2020
209, Petaluma,
USA
4
5

Astronomy Department, California Institute of Technology, Pasadena, CA 91125, USA
Université de Grenoble, Institut de Planétologie et d’Astrophysique de Grenoble (IPAG), F-38000 Grenoble, France; CNRS, IPAG, F-38000 Grenoble, France
6
Center for Computational Astrophysics, Flatiron Institute, New York, NY 10010, USA
Received 2020 February 24; revised 2020 April 7; accepted 2020 April 8; published 2020 May 26

Abstract

We present an analysis of K2 light curves (LCs) fromAbstract
Campaigns 4 and 13 for members of the young (∼3 Myr)
Taurus association,Weinpresent
addition
to an
older
(∼30(LCs)
Myr)
population
of 13
stars
that isoflargely
an analysisof K2
light curves
fromCampaigns 4 andfor members
the youngin(∼3theMyr)foreground of the
Taurus association,
additionan older(∼30 Myr) population ofstars thatmembers,
is largely in theof the81% are periodic.
Taurus molecular clouds.
Out ofin156
of tothehighest-conﬁdence
Taurusweforegroundﬁnd that
Taurus molecular clouds. Out of 156 of the highest-conﬁdence Taurus members, we ﬁnd that 81% are periodic.
Our sample of young
foregroundstars is stars
biasedandandincomplete,stars(37/38)
Our sampleof young foregroundis biased
incomplete, but but
nearlynearly
all stars all(37/38)are periodic.
Theare periodic. The
distributionrates
of rotation
as a function
color (afor mass)
similar isto thatfound inoverall distributionoverallof rotationas aratesfunctionof ofcolor(aproxy
proxy
for ismass)similartootherthat found in otherclusters: the slowest rotators are among the early M spectral types, with faster rotation toward both earlier FGK andclusters: the slowestlaterrotators
are among the early M spectral types, with faster rotation toward both earlier FGK and
M types. The relationship between period and color/mass exhibited by older clusters such as the Pleiades islater M types. The alreadyrelationshipand population
color/mass
exhibitedolder clusters
such
in place bybetween
Taurus age. period
The foregroundhas veryfew stars butby
is consistent
with the USco
andas the Pleiades is
Pleiadesperioddistributions.
As
found
in
other
young
clusters,
stars
with
disks
rotateon
averageslower,andfew
already in place bywith
Taurusage. The foreground population has very few stars but is consistent with the USco anddisks are found rotating faster than ∼2 days. The overall amplitude of the LCs decreases with age, and higherPleiades period distributions.
foundlower
in otheryoung
clusters, stars.
starsStars
withaveragemass stars haveAs
generally
amplitudes
than lower-mass
withdisks
disks rotate
have on on
averagelarger slower, and few
amplitudes
thanfasterstars without
though thephysicalmechanisms
drivingof
thethevariability
the resulting
LC age, and higherwith disks are foundrotating
than disks,∼2 days.
The
overall
amplitudeLCs anddecreaseswithmorphologies are also different between these two classes.
mass stars have generally
lower amplitudes than lower-mass stars. Stars with disks have on average larger
Uniﬁed Astronomy Thesaurus concepts: Star formation (1569); Stellar rotation (1629); Starspots (1572)
amplitudes than stars
without disks, though the physical mechanisms driving the variability and the resulting LC
Supporting material: machine-readable tables
morphologies are also different between these two classes.
Uniﬁed Astronomy Thesaurus
concepts: Star formation (1569);
StellarStarspots
1. Introduction
preparation)
will rotation
explore the (1629);
range of LC
properties (1572)
of thedisked members.
Taurusの光度曲線。早期M型で遅い回転、早期FGKと後期M型で速い回転。
The Taurus-Auriga
star-forming region
has fundamentally
Supporting material:machine-readable
tablesWe have previously published our analysis of the K2 data forshaped our understanding of how low-mass stars form (see,光度曲線の強度は年齢とともに下がり、重いほど弱い。円盤があると強い。
the Pleiades (∼125 Myr; Rebull et al. 2016a, 2016b; Stauffere.g., Kenyon & Hartmann 1995). At an age of 3 Myr (e.g.,
Kraus & Hillenbrand 2009; Luhman 2018) and a distance of∼140 pc (e.g., Esplin & Luhman 2017), even the low-mass
1. Introduction
objects(mid-M to early-L) are bright enough to be wellstudied. A signiﬁcant fraction of the members have an infrared

et al. 2016b; Papers I, II, and III, respectively), Praesepe(∼790 Myr; Rebull et al. 2017; Paper IV), and Upper Sco/ρOph(∼8 and ∼1 Myr;et al. 2018;
V). Wepreparation)
willRebull
explorethePaperrange
ofhaveLC
deliberately performed our analyses of these clusters, now

disked members.

properties of the