Star Formation Newsletter#341 7-12
Yusuke Aso (KASI)

514v1 [astro-ph.SR] 2 May 2021

lative distributions and fitting a Gaussian radius(CSF), which can be considered the probability that any
870 MICRON DUST CONTINUUM OF THE YOUNGEST PROTOSTARS IN OPHIUCHUS
for our Class I and Flat sources. This allows us
given system in the sample has companions (MF) and11
24Frankie J. Encalada,
LeslieLooney, John J. theTobin,
Sarah I.
Sadavoy,3ofDominiqueSegura-Cox,
tter include the unresolved
sources intheW.
estimate.
averagenumbercompanions(CSF).
The quanti565+1.8+2.8 Erin Cox, Zhi-Yun Li, and Giles Novaknd mean radii of 23.5 1.2 1and 16.5 0.9 au, respecties are defined by
Department of Astronomy, University of Illinois, 1002 W. Green St., Urbana, IL 61801, USA. This is in comparison to
Tobin
et al. (2020),
2 National
Radio Astronomy Observatory, 520 Edgemont Rd., Charlottesville,VA 22903, USA+5.8 3+4.9+3.7
B + T + Q + ...
原始星期の円盤質量、半径の領域ごとの違い。
h finds Ø
44.9
28.5Engineering
forAstronomy,
Class Queen’s University, 64 BaderM
Departmentof Physics,
Lane,
Kingston, ON, K7L 3N6, Canada
3.4 , 37.0
3.0 , and
2.3 au andF=(3)
4 Center for Astrochemical Studies, Max Planck Institute for Extraterrestrial Physics, 85748 Garching, Germany
S
+
B
+
T
+
Q
+
...
Ø VANDAM
Orionで質量（降着かダスト成⻑）、半径（謎）が縮⼩。質量はOri,
Per, Tau > Oph。
and Flats,
respectively.
We only have one Class 0
5 Center for Interdisciplinary Exploration and Research in Astrophysics (CIERA), 1800 Sherman Avenue, Evanston, IL 60201, USAØ 連星率は重いほど、若いほど⾼い。Ori,
Perで75au（円盤）と3000
au（乱流）にピーク。
e (Oph-emb1), which has an6 Department
image Gaussianfitandof Astronomy, University of Virginia, Charlottesville, VA 22904, USA p
Ophの円盤は1/3がM
Multiplicity
Companion
Star
FigureThesources fromFraction=0.29,
our sample. Dotted contours are
at -3 and solid contours
are at 3Fraction=0.41。連星の⽅⇥ 2 for n 2 W.
dust>10M
E, multiplicity
s of 15Øau,so we did not use
the 3.CDF
methodOtherwise the same as Figure 1.
B + 2T + 3Q + ...
が軽い円盤。
ssed above.
We note that 15 au is much smaller ABSTRACT
CSF
=
,
(4)
00
Ø 1 Class
Class 00,radii
23
ClassI,a70.15
Flat。ALMA,Band
7, 0.15"-7.2"。
We
present
resolution
(21
au)
ALMA
870
µm
continuum
surveyof
25
pointings
containing
31
S
+
B
+
T
+
Q
+
...
the meanin
Orion,
but
with
so
few
Encalada et al.
young stellaretobjects
in the Ophiuchus molecular clouds. Using the dust continuum as a proxy for dust
（Class
II/IIIはCox
al. 2017）
0 sources in Ophiuchusitisdifficultto draw any
S,a B,
andQ+2.1
refer to +9.2
the number of single,
mass and circumstellar disk radius in our sample,where
we report
meanT,
massof 2.8
1.3 and 2.5 1.1 M+1.8 radii and+2.8
conclusions. However,diskbinary, triple, and quadruple systems.
and a overall,mean radiitheof 23.5
1.2 and 16.5 0.9 au, for Class I and Flat spectrum protostars, respectively.
mass for OphiuchusInprotostars
tothe
bemultiplicity
smaller statistics ofaddition, we seem
calculatethe dust surrounding
young
stellarobjectsare
in derived assumUncertainties
for thetwo
metrics
Ophiuchus.
multiplicity
fractionevo(MF) and companion
star fraction (CSF) of the combined Class
0.5
ビジビリティの例。
those in Orion without
any Theclear
trend withing
binomial
distributions
as
=
(R(1
R)/N
)
,
R
sys
I and Flats based連星サンプル。
solely on this work is 0.25 ± 0.09 and 0.33
± 0.10, respectively, which are consistent
振動が円盤。
n, whereas Orion andconsistent
with While wewhereR di↵erences
is either inthe
MF
CSFbetween
value and Nsys is the
withPerseus
the values arefor Perseus
and Orion.
see clear
mass
andorradiusØ
円盤半径はdeconvolved
2σ。
other.
the Ophiuchus and Perseus/Orion protostellar surveys,
we doofnotsystems.
see any significant
di↵erencesin thenumberWith thoseassumptions,the values
円盤質量は光学的に薄いとして
multiplicities of the various regions. We posit there are some di↵erences inØthe
conditions for starweobtainfor
only
thecombinedClass
I and F sources
formation in Ophiuchus that strongly a↵ects disk size (and consequently disk （VANDAMmass), but doesOrionに合わせて
notin Ophiuchus
surveyedin thispaper are MF = 0.25 ±a↵ect system multiplicity, which could imply importantvariation in planet
formation
1.84processes.
cm2 g-1）。
5.2. Multiplicity
0.09 and CSF = 0.33
± 0.10, see Table 3. These areØ Tdust=43 K (Lbol / 1Lsun)1/4。
Keywords:
protoplanetarydisks — stars:— stars: pre-main
sequence— radiointerferomfind multiplicity in
our sample,
we visually
in- formationconsistent
with themetricsof thecombined Class I andetry
ed our maps out to the FWHM of the primary
F sample of Perseus with a separation range of 20 to
Figure 4. An example of the real and imaginary components of the binned visibilities of a resolved source (Oph-emb 9) inof ALMA. We define a companion
protostar as pany
1000
al. a(2021,
inwhich
1. INTRODUCTION
tle auwithfrom
manyTobin
authorset
using
T The <
70 preparation),
K for Class
our sample. This is shown solely
to demonstrate the typical shape of the curve that shows evidence of a disk. bol
imaginarymultiplicity sources from our sample. Dotted contours are at -3 and solid contours are at 3 ⇥ 2 for n 2 W.
component(bottom) is expected
to be around zero.
The0real
component(top) =
shows
the extended
onand
the left As(thesources(e.g.,
Dunhamal. 2014).
the =protoehee samewitha 1.S/N &Molecular
4 or if theimfit
taskcores
required
two
MF
0.28
±etenvelope
0.08
CSF
0.28 ± 0.08,
cloudsform dense
and eventuallypro- reported
as Figure
2

2

2n

2n

Oph-emb 27 A - B

229.9

1

J162636.8-241552

4.6

2

J162715.4-242640

35.4

2

870 MICRON DUST CONTINUUM OF THE YOUNGEST
PROTOSTARS
IN OPHIUCHUS
J162648.5-242839
590.8
2
1
2
3
Embedded Protostars
in Ophiuchus
11Frankie
J. Encalada,
Leslie W. Looney,1 John J. Tobin,
Sarah I. Sadavoy,
Dominique
Segura-Cox,4J163152.1-245616
416.8
2
6
5
Erin Cox,5 Zhi-Yun Li,
Novak167.1
VLAand1623 AGiles-B
3,4

514v1 [astro-ph.SR] 2 May 2021

VLA 1623 A1 - A2
26.5
4,5
1 Department of Astronomy, University of Illinois,1002 W. Green St., Urbana, IL 61801, USA
IRAS 16293-2242 A - B
734.2
6
2 National Radio Astronomy Observatory, 520 Edgemont Rd., Charlottesville,VA 22903, USA
IRAS 16293-2242 A1 - A2
40.3
6
3 Department of Physics, Engineeringab ON, K7L 3N6, Canada
Queen’s
University,
64A BaderLane,companion.
Kingston,
HigherandorderAstronomy,
systems were reduced
to binaries
with sourceas the common
All distances adjusted to 140 pc.
1
2
3
4
5
4 Center for Astrochemical Studies,This paper.
et al. Institute(2019). Looneyal. (2000). Harris Physics,
et al. (2018).85748
Hsieh etal. (2020). 6Germany
Wootten (1989)
MaxCieza
Planckfor etExtraterrestrial
Garching,

Ø 同じ条件でもOph円盤が軽い。
5 Center for Interdisciplinary Exploration and Research
Astrophysics (CIERA), 1800 Sherman Avenue, Evanston, IL
60201, USAØ in可能性１）Ophサンプルは進化したClassI。
6 Department of Astronomy, University of Virginia, Charlottesville, VA 22904, USA２）ダストが違う。３）星形成の初期条件が違う。
１）はOri
Flatより軽いし、２）は報告がないので３）。
ABSTRACT
カプランマイヤー推定でnon-detectionも考慮して平均質量。
We present a 0.1500 resolution (21 au)ØALMA
870 µm continuum survey of 25 pointings containing 31
Classclouds.
I 2.8 M
2.5 continuum
ME (cf. Ori.
11.6
ME)。
young stellar objects in the Ophiuchus molecularUsing
the dust
as a 14.9,
proxy fordust
E, Flat+2.1+9.2
mass and circumstellar disk radius in our
sample, we report a meanmass
2.8 Ori.
2.5 28.5Ø 半径も同様に23.5,
16.5
auof(cf.
37.0,
1.3 and
1.1 Mau)。

+2.8
and a mean radii of 23.5+1.8
for3. Class
I and
Flat
spectrum
respectively.
1.2 and 16.5 0.9 au,
Table
Multiplicity
Fractionand Companion
Star protostars,
Fraction
In addition, we calculate the multiplicity
statistics of the dust surrounding
young
stellarobjectsinSamplea
S:B:T:Q
MF
CSF
Ref.
Ophiuchus. The multiplicity fraction (MF)and companion star fraction(CSF)
of the combined1 Class
Ophiuchus Class I + Flats (Paper)
18:4:2:0
0.25 ± 0.09 0.33 ± 0.10
I and Flats based solely on this work isOphiuchus
0.25 ±Class
0.09
0.33 +±Cieza)
0.10,72:7:2:0
respectively,which
I + and
Flats (Paper
0.11 ± 0.04
0.14 ±are
0.04 consistent
1,2
PerseusClass I +Flats
21:8:0:0 in 0.28
± 0.08and
0.28radius± 0.083
with the values for Perseus and Orion.
While
we
see clear di↵erencesmass
between
Orion
Class
I
+
Flats
182:32:2:0
0.16
±
0.02
0.17
±
0.03
theet Ophiuchus
andPerseus/Orion
we do not see any significant di↵erences3 in theure 5. Cumulative distribution from Tobin
al. (2020) with data
addedfrom this paper (andprotostellar
the non-overlappedsurveys,
sourcesa1
2
maximum
separation
range
is 1000
This paper.
Cieza et al. in(2019).
et al. (2021, in
Williams et al. (2019)). 連星距離（三重は主星からのみ）。
The key di↵erenceis that the Ophiuchus
Classvarious
I andThe
flat sampled
spectrum
masses
are
even
lower
than
multiplicitiesof theregions.
We
posit
there
areau. somedi↵erencesthe3 Tobin
conditions
forpreparation)
starØ
OphとOriはMFとCSFが近い。
inthe
20-1000
au
range
using
their
probabilistic
weightingwhich
corrects
for more crowded regions.
original figure.
formation in Ophiuchus that strongly a↵ects disk size (and consequently disk mass), but does notØ 今回⻑距離連星は観測できないがPerの双峰に
a↵ect system multiplicity, which could imply important variationin planet formation processes.

⽭盾しない。
Keywords: protoplanetary disks — stars: formation — stars:
sequence — pointに重みがあるAD検定でradio interferomØ pre-main
KS検定よりend
etry
同じ分布を棄却できない。
Ø 近接連星と円盤は関係していそうだが、違う。
1. INTRODUCTION
tle with
many authors using a Tbol < 70 K for Class

Molecular clouds form dense cores and eventually pro-0 sources (e.g., Dunham et al. 2014). As the proto-

1.5

3

4

5

1.0

1.0

1.0

0.5

0.5

0.5

0.0

0.0

−0.5

−0.5

−1.0

−1.0

−1.0

−1.0

−0.5

0.0

0.5

1.0

−1.5
−1.5

1.5

−1.0

−0.5

x (pc)

∗

1.5

0.0

0.5

1.0

ASURA, eps=10000au, 10Myr

0.0

−0.5

−1.5
−1.5

Received ; Accepted

1.5

ASURA, eps=1000au, 10Myr

y (pc)

2

1.5

Initial Conditiony (pc)1y (pc)

arXiv:2101.05934v2 [astro-ph.GA] 15 Apr 2021

SIRIUS project. II. a new tree-direct hybrid code Ø クラスター形成シミュレーションは0.1 Msunのシンクだと1000 Msun (~散開星団)まで。
for smoothed particle hydrodynamics/N -bodyØ さらに⼤質量はシンクにクラスターを⼊れる、星simulations of star clustersの衝突なし。しかし⼤質量星の移動は重要。
Michiko S. F UJII1 , Takayuki R. S AITOH2, 3 , Long WANG1 ∗ and Yutaka
6
Publications of the Astronomical Society of Japan, (2014), Vol. 00, No. 0
H IRAI4, 5*
Ø IMFから確率的に星質量を決めて、N体も⼊れてDepartment of Astronomy, Graduate School of Science, The University of Tokyo, 7-3-1>104 Msunさらに球状星団、銀河スケールへ。
Hongo, Bunkyo-ku, Tokyo 113-0033, Japan
Department of Planetology,
Graduate Schoolof Science, Kobe University, 1-1 Rokkodai-cho,Ø ASURA:重⼒softening(Plummer)有りのN体SPH。
Nada-ku, Kobe, Hyogo 657-8501, Japan
Ø Bridge:
惑星軌道＝太陽周りのケプラー＋惑星によるN体摂動。
Earth-Life Science
Institute, Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro-ku,
Tokyo 152-8551, JapanØ N体星団を含むN体銀河で銀河粒⼦による星団粒⼦（星）を摂動にBridgeを使って星団内の衝突を扱う。
Astronomical Institute, Tohoku University, 6-3, Aramaki, Aoba-ku, Sendai, Miyagi 980-8578,
Publications of the Astronomical Society of Japan, (2014), Vol. 00, No. 0
Japan
Ø AMUSE+BridgeのSPH/N体ではコード間のデータコピーが重いので、ASURA+Bridge。
RIKEN Center for Computational Science, 7-1-26 Minatojima-minami-machi, Chuo-ku, Kobe,
Ø PETAR:N体で短距離⼒をparticle-particle、⻑距離⼒をparticle-tree、⾼離⼼連星にも対応。
Hyogo 650-0047,
Japan
3
W
E-mail: fujii@astron.s.u-tokyo.ac.jp,
saitoh@people.kobe-u.ac.jp,
long.wang@astron.s.u-tokyo.ac.jp,
Ø CELIB:
⾼密度、⾼温、divv<0でKroupa
IMFから質量を決めて星。0.2 pc以内のガス質量が半分以下なA
yutaka.hirai@astr.tohoku.ac.jp
10
ら決め直す。
d
AG3S3
−1.5
−1.5

1.5

−1.0

−0.5

x (pc)

1.5

ASURA+BRIDGE, dt=50yr, eps=1000au, 10Myr

1.5

ASURA+BRIDGE, dt=200yr, eps=1000au, 10Myr

1.0

1.0

0.5

0.5

0.5
0.0

−0.5

−0.5

−1.0

−1.0

−1.0

−0.5

0.0

0.5

1.0

1.5

−1.5
−1.5

−1.0

−0.5

x (pc)

0.0

0.5

1.0

1.5

−1.5
−1.5

|∆E/E0 |y (pc)y (pc)y (pc)

0.0

−0.5

−1.0

1.0

1.5

ASURA+BRIDGE, dt=400yr, eps=10000au, 10Myr

Abstract10
テスト１：ガスに埋もれた星団（Bridgeの有無）
Star clusters form via clustering star formation inside molecular clouds. In order to understand5粒Ø theスケール⻑0.5pc,
SFE=0.3,
844 inM
ガスを10dynamical evolution of starclustersin their early phase,whichclusters are still
emsunstarbedded in their surrounding gas, we 3need an accurate integration of individual stellar orbits
⼦、Salpeterから10 個の計362 M andの星。
10
without gravitational softening in the systems including both gassunstars, as well as modelingstars with a realistic mass function. We develop a new tree-direct hybrid smoothedØ individual
断熱進化で星形成も無し。
particle hydrodynamics/N -body code, ASURA+BRIDGE, in which stars are integrated usØ ing
bridgetimescaleが⻑いと星団中a direct N(摂動を計算する)-body scheme or P E TAR, a particle-particle
particle-tree scheme code, without
10
ASURAのみだと近いエン
gravitational softening. In ASURA+BRIDGE, stars are assumed to have masses randomly
⼼がガス中⼼からずれる、⾃由落下の1％がよい。
drawn from a given initial mass function. With this code, we perform star-cluster formation カウントでエネルギー誤差
starting from molecular clouds without gravitational softening. We find that artificial
Ø simulations
ガス、星の密度分布はBridgeを⼊れても同じ。
がジャンプ（保存しない）。 10
−1.5
−1.5

0.5

−1

1.0

0.0

0.0

x (pc)

−1.0

s
C
t
H
b
s

BG3S3T20, dt=200yr, !s =1000au
BG3S0T20, dt=200yr, !s =0
BG3S0T10, dt=100yr, !s =0
BG3S0T05, dt=50yr, !s =0

−2

−3

−0.5

x (pc)

0.0

0.5

1.0

1.5

t
d
t
a

x (pc)−4
Fig. 1. Gas and stellar distributions of initial condition (top left) and models AG3S3 (top middle), AG4S4 (top right), BG3S0T05(bottom left), BG3S0T20(bottom middle) and BG4S0T40 (bottom right) at 10 Myr. The density plots are a slice through the mid-plane of the gas density. Points show the stars and themarker size depend on the stellar mass.

Hermiteintegrator
cansoftenmaintain the energy error sufficiently
dense cores in star-cluster centers due to the softening disappear the
when
we do
not usesmall ofeven
withoutclumps.
softeningStar
length.
ing. We further demonstrate that star clusters are built up via mergers
smaller!g = 1000 au−5

thanks to the larger softening length. 0For comparison
between
2
4
6
ASURA and ASURA+BRIDGE, we also show the snapt (Myr)

8

10

a
s
t

SIRIUS project. II. a new tree-direct hybrid code テスト２：星団形成
for smoothed particle hydrodynamics/N -body Ø 乱流分⼦雲、k=-4 (Bonnell+03), M =1000 M ,simulations of star clusters
R =0.5 pc, |E |=|E |, 20 K, P=(γ-1)ρu、uは内部E。
Ø 質量分解能をhigh0.002, middle0.01, low0.1 M とし
Michiko S. F UJII1 , Takayuki R. S AITOH2, 3 , Long WANG1 ∗ and Yutaka
4, 5*
H IRAIて⽐較。1 Myr=4 ffまで。
乱流場の種の違い
Department of Astronomy, Graduate School of Science, The University of Tokyo, 7-3-1
ØSFE
per Tokyo
ff=c113-0033,
Hongo,
Bunkyo-ku,
Japan*=0.02、桁で変えても最終質量の差は10%。
Department of Planetology, Graduate School of Science, Kobe
University,
1-1 Rokkodai-cho,
3, 1.4x10
4 auも最終質量は変えない。
MiddleWe also4show
softening
length andØ
gas softening length=2.8x103, 7x10
F16× 10 the results
0.1 with
10a larger
3800
0.83
1
0.07
0.0
0.02
0.2
200
∆t /1024 10
600 5 × 10
Nada-ku, Kobe, Hyogo 657-8501, Japan
High
lower threshold density for star formation (models Middle-l andEarth-Life
Science Institute, Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro-ku,
200Ø
乱流場の種は最終質量を2倍変える。
Low-l). In these models, star formation starts in earlier time (seeTokyo 152-8551, Japan
400
Fig. 6). However, the final stellar mass does not change muchØ
MFは与えたものになったし、星密度分布も分解能に依らない。
AstronomicalInstitute, Tohoku University, 6-3, Aramaki, Aoba-ku, Sendai, Miyagi 980-8578, from those with a smaller softening length especially in the case
2

−3

B

Ms (M!)

From the left: model name, initial cloud mass (Mg ), gas-particle mass (mg ), initial cloud radius (Rg ), initial cloud density (nini ), initial free-fall time (tff,ini ), initial virial
ratio (αvir ), softening length for gas ("g ) and stars ("s ), star formation threshold density (nth ), the maximum search radius (rmax ), timestep for Bridge (∆tB ), timestep for
PETAR (∆tsoft ), cut-off radius for PETAR (rout ).

3

4

5

5

N

5

N

1

1025

Japan
10 F16of the lowest mass resolution.
RIKEN Center for Computational Science, 7-1-26 Minatojima-minami-machi, Chuo-ku, Kobe,
Hyogo 650-0047, Japan

10010
200

10
テスト３：複合星団
3.2.4 Structures of5 formed star clusters
05
10
100
10
10
10
10
01
10
E-mail: fujii@astron.s.u-tokyo.ac.jp, saitoh@people.kobe-u.ac.jp,
long.wang@astron.s.u-tokyo.ac.jp,
5
0.0
0.2
0.4 a (au)
0.6
0.8
1.0
We
investigate
thestructuresoftheformedstar
clusters
in
our
Ø
半径10
pc,
質量4x10M,
分解能0.1M。
sunsunyutaka.hirai@astr.tohoku.ac.jp
t (Myr)simulations and the dependence on the resolution. In Fig. 8, we0
100Ø 0.75 Myrで星の数が~8000になるのでHermite present the stellar massfunctions obtained from our simulations
at 1 Myr. We confirm that formed stars follow the Kroupa mass10
schemeからPTEARに。星100個以上の星団が25個。
function we gave.−5Although we set the maximum mass of the−5
Abstract
m
10
1
mass
function
to
be
100M
,
we
did
not
find
such
a
massivestar
Ø
星団でできた⼤質量星が⾶び出す。（今はないtimStar clusters form via clustering star formation inside molecular clouds. In order to understand=1.0
Myr
set. This is simply due tot theprobabilitytoth
the dynamical evolution of star clusters in their early phase, in which star clusters are still em- in this simulation
−10
10
−10
が）フィードバックを弱める。
form such a massive−10
star in ∼−5500M starlow.
0 cluster
5 is too 10
5 the clusters
10 form via−10a sub-cluster(Model −5
Low, seed07). Since
bedded in their surrounding gas, we need an accurate integration of individual stellar orbits
x (pc)x (pc)
0.1
the mergerof smaller clumps, we sometimes find a surviving lo
without gravitational softening in the systems including both gas and stars, as well as modeling We compare the density profiles of formed clusters. In theØ
平均⼒学エネルギーの0.1%以上の束縛エネルギーで連星とする。
stsub-cluster at 1 Myr.
individual stars with a realistic mass function. We develop a new tree-direct hybrid smoothed density profiles, we find a small difference in the resolution.
of the runs with theWparticlehydrodynamics/N-body code, ASURA+BRIDGE, in which stars are integrated us- In Fig. 9, we present the density profiles
Ø a=100-10
auまでできる。遠い連星系が崩れて、連星率は0.7-0.8
Myrでsame random seed, but different mass resolution and softeningcl
0.7Myring a direct N -body scheme or P E TAR, a particle-particle particle-tree scheme code, without at t = 1 Myr.
3.30.01
Test 3: Star cluster complex model 0.8Myr
lengths. In the lowest resolution with a larger softening length下がる。
pegravitational softening. In ASURA+BRIDGE, stars are assumed to have masses randomly
0.9Myr(model Low-l), the formed cluster has a density slightly lowerWeinadditionpresent
theresultsofalarger-scale
simulation,
drawn
from a given initial mass function. With this code, we perform star-cluster formationis one or two-order of magnitude higher than the1Myr
observed ones
Ø 重いほど短周期で質量差⼤。観測では質量差は⼩さいが、1Myr以降の30000
comparedwith the others.
This is due to the large softening
in which multiple star clusters, so-called star clustercomplex, ca
simulations starting from molecular clouds without gravitational softening. We find that artificial(seefigure2
inPortegies
Zwart
et
al.
2010).
Including any
10length of gas, although we did not assume gravitational softenform in0.1
a turbulent molecular
cloud. 10
1100ul
segregationが効きそう。
densemasscores in star-cluster
centers due to the softening disappear when we do not use soften25000
feedback, the density later mwill(Mdrop) due to the gas expulsion.
ing for stars. For the other models, the density profiles of thebuing. We further demonstrate that star clusters are built up via mergers of smaller clumps. StarThis implies that star clusters were born with a density muchformed20000
clusters are identical.
Ø
次の論⽂でフィードバックも⼊れてより⻑く。
3.3.1 Initial condition and integrationclclusters formed in our simulations include some dynamically formed binaries with the minimum
higher than currently observed ones and became less dense after
24

3

4

5

6

Σ (cm−2 )y (pc)

Received ; Accepted22y (pc)

23∗

Fig. 14. Semi-major axis (left) and period (right) distributions of binaries at 0.7, 0.8,

21 Fig. 7. Run-to-run variation in stellar mass evolution for star cluster formation

models (model Low, Middle, and High). Thick curves in models Middle and

20 Low indicate the three models which have the three random seeds for the

turbulence used in model High.

!

19

fbin

!

Fig. 12. (Left) Snapshot at t = 1 Myr for Model F16. White points indicate stars more massive than 1M! , blue points are massive stars with > 20M! . Red
crosses indicate the positions of binaries with a semi-major axis of < 1000 au. (Right) Snapshot of clusters detected using HOP algorithm

−3

1

!)

arXiv:2101.05934v2 [astro-ph.GA] 15 Apr 2021

Ms (M

Low
for the initial turbulent velocity field up to 1 Myr. We show the400
Low-c0.1
time evolution of the total stellar mass of these runs in Fig. 7.
Low-l
As shown in this figure, one of them forms only ∼ 400M! stars
200
in total by 1 Myr, although the others form more than 600 M!
stars. We in addition perform 5 and 10 runs for models Middle
and Low with different random seeds. The maximum mass is Publications0of the Astronomical Society of Japan, (2014), Vol. 00, No. 0
g0.2
0.0
0.4
0.6 sun0.8
1.0
1.2
twice as large as the minimum one. Irrespective of the mass rest
(Myr)olution, models with the random seed which results in the forgkp500
mation of less number of stars always form stars less than those
0.7theMyr
Fig. 6. Time evolution of
total stellar mass for the same random seedwithsun(2014),
0.8 MyrSociety
14 the other seeds. Thus, the forming stellar mass slightlyPublicationsof theturbulent
Astronomicalof Japan,00,cluster
No. 0
for the initial
velocity
field of themolecular
cloud Vol.
for stardepends on the mass resolution, but we can reproduce the vari0.9
Myrformation400
model.
1 Myr
ation
onstarthe cluster
turbulentvelocity
fields. We summarize
Tabledepending
3. Model forcomplexformation
800
theNameaveraged M
total
massesinmgat 1 Myr
Rg with nstandard
tdeviations
αvir
"g
"s
nth
c!
rmax ∆tB∆tsoftroutgini
ff,ini−3table 2.
300 (cm−3 ) Low(M! )(M! ) (pc) (cm ) (Myr)(pc) (pc)
(pc)
(yr)
(pc)

We also investigate the run-to-run variations in the density!

Fig. 15. Binary fraction as a function of the primary mass (m1 ) for model

We adopt an initial condition similar to Fujii & Portegies Zwart ou(MCMC) random sampler, emcee (Foreman-Mackey et al. 2013), w in Eq. 1 equals one everywhere. Each run takes awhich efficiently
walks
a vastof parameters
to deteron anda 48-coremachine with a clock rate of 2.3 GHz
The
DiscoverMiner
I: Arange
statisticalframework
to detectquantify
mine the subset ofthem that best
reproducesdriven
the projected
lineplanets
Table
1 presents a summary of the functional
kinematical
perturbations
by young
in discs
intensity of the input disc,often encoded in a three-dimensional F. vaneredfor the model attributes and the best-fit param
A. F. Izquierdo1,2 , L. Testi1,3,4 , S. Facchini1 , G. P. Rosotti2,5 , and E.
Dishoeck2,62
ppv cube. We configure the MCMC sampler to maximise afor the three snapshots. Figure 4 shows the bes
European Southern
log-likelihood definedas, Observatory, Karl-Schwarzschild-Str. 2, 85748 Garching bei München, Germanycomputed for the 0.3 MJup snapshot, deprojected
e-mail: andres.izquierdo.c@gmail.com
Leiden Observatory, Leiden University, P.O. Box 9513, NL-2300 RA Leiden, The Netherlands
and lower emitting surfaces of the disc, and theØ 原始惑星は円盤内構造を作り得るが、直接観測例は少ないし、それ以外の検出⽅法は天体ごと。
Cluster Origin and Structure of the Universe, Boltzmannstr. 2, 85748 Garching bei München, Germany
n
n
pix
ch
X XExcellencehiINAF – Osservatorio Astrofisico di Arcetri, Largo E. 2
5, 50125 Firenze, Italytion
of-2boththe surfaces
Ø FARGO3D
disk,
RFermiau,
g cm, γ=1,
α=10-3(see, M*parameter=1Msun。walkers in
2 流体シミュレーション。tail2 Physics and Astronomy, University
Schoolofof Leicester,
Leicester LE1
7RH,Σ
UKc=0.03
c,=100
= 0.5
w
I
(r
,
)
I
(r
,
)
(1)
m
i
j
d
i
j
i
Max-Planck-Institut
für extraterrestrische Physik, Gießenbachstr. 1 , 85748 Garching bei München,WeGermany
notice1.0,
slightvariationsin theau。
best-fit paramete
R=15-700 au, |z|=0-240
au。0.6Hでスムージングした2D重⼒場,
M
=0.3,
3.0
M
,
R
=100
p
J
p
j
iReceived 10
March 2021 / Accepted 21 April 2021
two (1.0 and 3.0 MJup ) snapshots due to the comb
円筒状の速度と静⽔圧平衡。
planetmass and the depth of the gap it carves.
ABSTRACT
where
the index
runs=4000
over theK。ダストはintrinsic
pixel location in(x, y)skycoordi-3.5the-3, q=-3.5, 5 nmから1 mm。
Ø POLARIS,
L*=1.9
Lsun,i T
densityg
cmeff
the astellar
Context. The
study of disckinematics
recently opened
up as a promising
method to detect
planets. However,systematic, mass (M? ) retrieved by the model inates, and the index
jmeaningful
runs
over
thehasinput
velocity
channels.
Theunseen cally,-5
statisticallyanalysisof such z/R<0.3
an approach remainselse
missing in5x10
the field. -11。
kink
Ø CO J=2-1,
XCO=5x10
if
T>20
K
and
2
planetmass(M?1.0= km/s
1.013, 1.025 M , respectivel
The aim of this work
is to devise anthe
automated,
statisticallyrobust, techniqueand quantify
kinematical
perturbations
1.6,
1.3,
kernel is theAims.
di↵erence
betweenmodel
intensity,
Imtoofidentify,theandinducedbythepresenceofplanets
in
a
gas
disc,
and
to
accurately
infer
the
locationplanets.
Ø d=100 pc,
i=45°でチャンネルマップ。
emission
are shallower due to the increa
Methods.
produce hydrodynamical
simulations isof planet-discinteractions with di↵erent
0.3, 1.0 andsurfaces
3.0 M , atthe input
intensity,
IdWe, whose
uncertainty
encompassedin theplanet masses,a radius of R = 100 au in the disc, and perform radiative transfer calculations of CO to simulate observables for a disc inclination
Ø φP=-90°から90°。
gaps (e.g., z0 = 15.59, 14.35 au for the 1.0 and 3
45 , andfor 13 planet
We then fit thesynthetic residual
data cubes with aintensity
Keplerian model of the channel-by-channel emission
weighting factorof w
which
is azimuths.
computedusing
using the discminer package. Lastly, we compare the synthetic cubes with the best-fit model to extract
deviations fromNevertheless,
Kepleriansurfaces).
these variations have negl
Ø -5から5fromkm/s,
0.1 km/s,
pix=3.1
au。and localised intensity, linewidth and velocity fluctuations triggeredline-freevelocity
channels.
rotation andquantify
both large-scaleby the embedded planets,
and to provide strong constraints on their location in the disc. We assess the statistical significance ofon
the detections
using the peak and quantification of planet pertur
the detection
and variance of the planet-driven velocity fluctuations.
Results. Our findings suggest that a careful inspection of line intensity profiles to analyse gas kinematics in discs
is a robust
methodTaken
together,the parameters retrieved by th
to reveal embedded, otherwise unseen planets, as well as the location of gas gaps. We claim that a simultaneous study of line-of3.2.
Line
intensity
profilesight velocities and intensities is crucial to understand the origin of the observed velocity perturbations.
In particular,
the combinedØ フィットツールDisc Minerを開発。
verge
notably
well to the features we knew beforeh
contribution of the upper and lower emitting surfaces of the disc plays a central role in setting the observed gas velocities. This jointe↵ect is especially prominent and hard to predict at the location of a gap/cavity, which can lead to artificial
deviations
from Keplerian
input
simulations.
The mass of the star (1 M ), theØ 上下２層z=Rの冪。
In this work, werotation
assumefunction
theofline
depending aon generalised
how the disc velocities bell
are retrieved.
Furthermore,forregardlesstheir origin, gas gaps alone are capable of
ギャップがあるだけでキンク。
producing kink-like features on intensity channel maps, often attributed to the presence of planets.tion
Our technique,line the height of the lower surface tra( 45based), onandØ 光学的に厚いラインプロファイルにも対応。
profile kernel, centroid di↵erences, takes all this into account to capture0.3
only thestrongest,
localised, planet-driven perturbations. It does not getM
は⽬では厳しい。
Jby axisymmetric velocity perturbations that may result from non-planetary mechanisms. The
methoddetectCO
all threesideofcanthefreeze-out region, are all closely rep
Ø IpとLwはRとzの冪。Lconfusedsは⼀様。
simulatedplanets, at all azimuths, with an !
average accuracy of ±3 in azimuth and ±8 au in radius. As expected, velocity fluctuations
driven by planets increase in magnitudestructureaandline-of-sightwhen
planetis present. This makes the discmine
2L s as a1function of the planet mass. Furthermore, owing to discl.o.s
projection e↵ects,chplanets at kazimuths
close to ±45 yield the highest velocity fluctuations, whereas those at limiting cases, 0 and, typically range within 40 70 m s , 70(2)Im (R, z; ch ) = I±90
+the lowest. The observed peak velocities
studies of the three-dimensional stru
170 m s also
and 130 for
450 mfutures for 0.3, 1.0
p ,1driveand 3.0 M planets, respectively.
near planets due to organised
Lw Our analysis indicates that the variance of peak velocities is boostedusing
molecular
gas motions prompted by the localised gravitational well of planets. We propose an approach that exploits this velocity
coherence to line emission.
provide, for the first time, statistically significant detections of localised planet-driven perturbations in the gas disc kinematics.
Ø 与えたパラメーターは再現。
where I p is theKeypeak
is thechanneland– radiative transfer
words. intensity,
planet-disc interactions
and satellites:
detection velocity– protoplanetary discs
ch– planets
kl.o.s is the Keplerian line-of-sight velocity. The choice of this
4. Finding the kinematical footprints of p
1. Introduction
Gabellinione
et al. 2019;
Facchini et al. 2020). However, planet-disckernel is motivatedby the fact that at the cost of only
addi1

Planet2
3
4

Xiv:2104.09596v2 [astro-ph.EP] 22 Apr 2021

5
6

Jup

p

1

1

Jup

12

1

A&A proofs: manuscript no. main

Ø ケプラーモデルとの差
PlanetMostly(gap edges)

Ø 左右の差（folding）

Planet(gap edges)

Gap

MostlyPlanetSymmetry axis

Gap

ギャップは⾒えなくなり、Mpが⼤きいほど差も⼤きい。
ギャップでは上下層で輝度が違うので中央速度がずれる。
Ø 半径ごとのピーク⾓とピーク値。

A. F. Izquierdo et al.: The Disc Miner I

A&A proofs: manuscript no. main

Gap

中央値

PlanetFig. 8. Left panel: peak residuals and their azimuthal location for the (1.0 MJup , p = 45 ) snapshot. Right panel: peak residuals rearranged as afunction of their radial distance. The subpanel on the right shows the peak residual distribution (green line) and 1, 2 and 3 significance thresholds(dotted lines). The solid black lines are the azimuthal and radial location of the global peak, while the dashed black lines show the actual locationof the planet.

Ø 惑星周囲もずれる
Fig. 7. Folded (magnitude of) line centroid residuals as a function of azimuth, varying the analysis radius(in colours), for di↵erent planet masses
ので、K-meanで(0.3, 1.0 and 3.0 M , from left to right), and azimuths (0 , 45 , and 90 from top to bottom). The solid black contour runs along the projecteddistance of the planet (R = 100 au), and the dashed black line corresponds to the azimuth of the planet.
コヒーレント部分を取り出してより
Now, is it possible to determine not only the presence, but di↵erent planet masses andazimuths (see Sec. 2). Both methods
also the actual location of the planets out of the just-gathered in- make use of the folded centroid residuals presented in Sec. 4.3.
formation? From Fig. 7 it appears so, especially for those planets
有意に検出。
From inspection of Fig.
7, a first natural approach would be
at azimuths smaller than 90 , where the perturbation is obvious
Juppto the eye. However, line centroid residuals from less massive
planets, such as the 0.3 MJup included in our analysis, may beparticularly challenging to detect just by visual inspection because the localised signature, intrinsic of the planet, is small andcan easily get confused with other asymmetric, although more
extended perturbations. Also, planets of any mass, near the symmetry axis of the disc, = 90 , are equally challenging because
their signature is no longer localised to the eye. In the next subsection, we provide two statistical methods that conveniently examine line centroid residuals and do not rely on visual inspection
to robustly infer the location of embedded planets in discs.

to determine the location of the global maximum centroid residual, which is apparently well correlated with the azimuthal andradial location of the planet. Following this idea, we extract thepeak velocity residual for each radius in the disc, as a functionof the azimuth where it occurs. From the resulting distributionof velocity residuals, we compute a 3 threshold below which
residuals are discarded from the analysis. The inferred locationof the planet is assumed to be the median azimuth and the median radius of the leftover peak residuals, as illustrated in Figure 8. This procedure is what we call the Global Peak detection
method. In Appendix B.2, we show the dependence of peak fluctuations as functions of planet mass by fitting= aM b powerlaws, for each planet azimuth, with a and b free parameters. Notethat there are substantial di↵erences between the observed fluctuations at negative and positive planet azimuths. As explainedlater in Sec. 5.3, the reason is that the disc vertical structure andØ 速度分解能0.05 km/sでも同じ。
Gap
Planet?
Ø 下層からの輝度は上層より弱いので、惑星の位置が同じでもどちらから⾒るかで左右差が変わることも。
Ø 鉛直⽅向の流れや観測雑⾳、観測ビームは今後。

all cases is between 3 and 7 , might be a weak spot in noisy planet6 . This method, the Variance Peak, is equally accurate as
scenarios. This leads us to introduce a rather di↵erent method, the Global Peak method when it comes to determining the locabased on the assumption that the strongest velocity fluctuations tion of planets. However, it strikingly boosts the significance ofdriven by planets should be, at the same time, localised and co- the detection thanks to the coherent nature of velocities aroundin other
implies thatthe velocity
field, as(bottom)
planets.
Fig.herent,
6. Linewhichcentroid(top),words
peak intensity(middle)and linewidth
residuals for the (1.0 MJup , p = 45 ) snapshot. The azimuthal scans
as itsshouldsmoothly
as oneapproaches
The absoluteof peak residuals
in theGlobalPeakon wellthe leftrunobservables,
along constantradiivary
contours,
in disccoordinates,
whose colours
representvalue
their closenessto the radial
locationof theplanet, withtotheplanet.
If
this
is
true,
besideslarge
velocity
residuals
there
method,andin thePeak method
are shows theblue being the closest. The solid black contour runs along the projected distanceof thethe peakplanetvariances(R p = 100au) Varianceand the dashedblack line
shouldbealsoahigherdensityofpeak
velocity
residuals
aroundpresented
inFigure10.
Unsurprisingly,
regardlessoftheplanetazimuth of the planet.
the planet compared to undisturbed regions of the disc. This azimuth, peak fluctuations and peak variances increase steadily
is pictured in Figure 9, where we find clusters of peak residu- with the mass of the planet. Also, due to projection e↵ects, theals, independently for each spatial coordinate (R, ), using a K- highest peaks are reached at intermediate angles (30 , 45 , 60 ).

4.4. Statistical methods to detect planets

original form ZZZrXiv:2105.05798v1 [astro-ph.SR] 11 May 2021

versity of Edinburgh, Edinburgh, UK
omy, Bucknell University, Lewisburg, PA 17837, USA
on and Research (IISER) Tirupati, Rami Reddy
Road,
Mangalam(P.O.), Tirupati
517 507,
Indiathe W-band filter
ANagar,NovelKarakambadi
Survey for
Young
Substellar
Objects
withanètes, Université de Montréal, Québec, Canada
III: Searching for very low-mass brown dwarfs in Serpens South
University of Bern, Gesellschaftsstrasse 6, 3012 Bern, Switzerland
and Serpens Coreof Hawaii at Manoa, Honolulu, HI 96822, USA
National Central University, Zhongli, Taiwan
Sophie Dubber1,2? , Beth Biller1,2 †, Katelyn Allers3 , Jessy Jose4 , Loïc Albert5 ,), Hosur Road, Bengaluru 560029, India

Blake Pantoja3 , Clémence Fontanive6 , Michael Liu7 †, Zhoujian Zhang7 ,Wen-Ping
Chen8 , Bhavana Lalchand8 , Belinda Damian9 , Tanvi Sharma8
1
SUPA, Institute for Astronomy, Royal Observatory, University of Edinburgh, Blackford Hill, Edinburgh EH93HJ, UK

2 Centre for Exoplanet Science, University of Edinburgh, Edinburgh, UK

3 Department of Physics and Astronomy, Bucknell University, Lewisburg, PA 17837, USA

4 Indian Institute of Science Education and Research (IISER) Tirupati, Rami Reddy Nagar, Karakambadi Road, Mangalam (P.O.), Tirupati 517 507, India
5 Institut de recherche sur les exoplanètes, Université de Montréal, Québec, Canada

6 Center for Space and Habitability, University of Bern, Gesellschaftsstrasse 6, 3012 Bern, Switzerland
7 Institute for Astronomy, University of Hawaii at Manoa, Honolulu, HI 96822, USA

Serpens South/Coreの低質量（<~0.1 Msun）メンバー候補の
We present CFHT photometry and IRTF spectroscopy of low-mass candidate members 可視⾚外撮像と分光。
Accepted XXX. Received YYY; in original form ZZZABSTRACT

8 Graduate Institute of Astronomy, National Central University, Zhongli, Taiwan
9 CHRIST (Deemed to be University), Hosur Road, Bengaluru 560029, India

of Serpens South and Serpens Core (⇠430 pc, ⇠0.5 Myr), identified using a novel combination of photometric filters, known as the W-band method. We report SC182952+011618,
SS182959-020335 and SS183032-021028 as young,ABSTRACT
low-mass Serpens candidate members,
We present CFHT
photometry and
IRTF spectroscopyof low-mass candidate memberswith spectral types in the range M7-M8, M5-L0 and ofM5-M6.5
respectively.
Best-fit
effective
Serpens South and Serpens Core (⇠430 pc, ⇠0.5 Myr), identified using a novel combitemperatures and luminosities imply masses of < 0.12M
all threecandidatememnation offor
photometricfilters,
known as the cluster
W-band method.
We report SC182952+011618,
SS182959-020335 and SS183032-021028 as young, low-mass Serpens candidate members,
bers. We also present Hubble Space Telescope imaging
datatypes(F127M,
andandF850LP)with spectralin the rangeF139M
M7-M8, M5-L0
M5-M6.5 respectively. Best-fit effective
temperatures
luminosities
imply
masses of <0.12M for all three candidate cluster memfor six targets in Serpens South. We report the discovery
ofandthebinarysystem
SS183044bers. We also present Hubble Space Telescope imaging data (F127M, F139M and F850LP)
020918AB. The binary components are separated byfor⇡45
AU,inwithtypestheof
M7-M8
six targets
Serpensspectral
South. We reportdiscoveryof the binary system SS183044020918AB.
The
binarycomponents
areseparated
by
⇡45
AU,
and M8-M9, and masses of 0.08–0.1 and 0.05–0.07 andM M8-M9,. Weanddiscuss
the effects of high dustwith spectral types of M7-M8
masses of 0.08–0.1 and 0.05–0.07 M . We discuss the effects of high dust
attenuation on the reliability of our analysis, as well attenuation
as the presence
ofofreddened
backgroundon the reliability
our analysis, aswell as the presence of reddened backgroundstars in our photometric sample.
stars in our photometric sample.
Key words: brown dwarfs – stars: low-mass – binaries: general

Key words: brown dwarfs – stars: low-mass – binaries: general
1 INTRODUCTION

describes the distribution of stellar masses at formation. Discov-

1 Department of Astronomy & Astrophysics

Pennsylvania State University
525 Davey Laboratory
X-ray Super-Flares
From Pre-Main Sequence Stars: Flare Energetics And Frequency
University Park, PA 16802, USA
2 Center for Exoplanetary and Habitable Worlds 1
Konstantin V. Getman and Eric D. Feigelson1, 2

68v1 [astro-ph.SR] 11 May 2021

Department
Astronomy & Astrophysics(Accepted for publication1in
ApJ, May of2021)
Pennsylvania State University
525 Davey Laboratory
ABSTRACT
Park, their
PA 16802,
Solar-type stars exhibit their highest levels of magnetic University
activity during
early USA
convective pre2 Center for Exoplanetary and Habitable Worlds
main sequence (PMS) phase of evolution. The most powerful PMS flares, super-flares and mega-flares,

Chandraで40の星形成領域内
24000の若い星（t<5 Myr）のhave peak X-ray luminosities of log(LX ) = 30.5 34.0 erg s 1 and total energies log(EX ) = 34 38 erg.
for40
publicationin ApJ, May2021)from 1086のX線スーパーフレア。主Among > 24, 000 X-ray selected young (t . 5 Myr)(Accepted
members of
nearby star-forming
regions
our earlier Chandra MYStIX and SFiNCs surveys, we identify and analyze a well-defined sample of 系列より強い。原始星から円盤
ABSTRACT
1,086 X-ray super-flares and mega-flares, the largest sample ever studied.
Most are considerably more
無しまで含み、重いほどフレア
powerful than optical/X-ray
super-flares
detectedon mainsequence
This study
presentsduring
energytheir early convective preSolar-typestars exhibit
theirhighestlevelsstars.
of magnetic
activityがある。円盤の有無で違いはな
estimates of these main
X-ray sequence
flares and(PMS)
the properties
their host stars.
These
events areproducedphase ofofevolution.
The most
powerfulPMS
flares, by
super-flares and mega-flares,
young stars of all masses
over X-ray
evolutionary
stages ofranging
protostars to diskless
1 stars, with the く、円盤ー星の磁場は関わってhave peakluminosities
log(Lfromand total energies log(EX ) = 34 38 erg.
X ) = 30.5 34.0 erg soccurrence rate positively correlated with stellar mass. Flare properties are indistinguishable for diskいない。エネルギーの分布は進Among > 24, 000 X-ray selected young (t . 5 Myr) members of 40 nearby star-forming
regions frombearing and diskless stars indicating star-disk magnetic fields are not involved. A slope ↵ ' 2 in the
our earlier Chandra MYStIX and SFiNCs surveys, we identify and analyze a well-defined sample offlare energy distributions dN/dEX / EX↵ is consistent with those of optical/X-ray flaring from older 化した星と同じ。

X-ray super-flaresand erg)mega-flares,
the largest
ever studied.
Most are considerably morestars and the Sun.1,086
Mega-flares(log(EX ) > 36.2
from solar-massstarssample
have occurrence
rate ofpowerful than optical/X-ray super-flareson X-raymain sequence
1.7+1.0
20% to thedetected
total PMS
energetics.stars.
TheseThis study presents energy
0.6 flares/star/year and contribute at least 10
estimatesof theseastrophysical
X-ray flarese↵ects
and theproperties of their
host stars. These events are produced by
explosive events may
have importanton protoplanetary
disk photoevaporation,
young
stars ofofall
masses overevolutionary
from protostars
ionization of disk gas,
production
spallogenic
radionuclides
in diskstages
solids,ranging
and hydrodynamic
escape to diskless stars, with theof young planetaryoccurrence
atmospheres.
Our
followingcorrelated
paper details
plasmaandmagnetic
modelingare
of indistinguishable for diskrate
positively
withstellarmass.
Flareloop
properties
the > 50 brightest bearing
X-ray mega-flares.
and diskless stars indicating star-disk magnetic fields are not involved. A slope ↵ ' 2 in the↵
1. INTRODUCTION(2018), Sciortino
et al.
and Ar- flaring from older
flare energy distributions dN/dE(2017),
with thoseof (2019)
optical/X-ray
X / EFeigelson
X is consistentstars and theMega-flares giroffi(log(E(2019).
X ) > 36.2 erg) from solar-mass stars have occurrence rate of
1.1. Pre-Main Sequence Super-flaresand Sun.
Mega-flares+1.0
Thoughsurprising
at first,
the energetics. These
1.7 0.6 starflares/star/year
and contributeatPMS
leastX-ray
10 flares
20% wereto thetotal PMS
X-ray
X-ray imaging studies of nearby
forming regions,existenceof strong magnetic
dynamosin the interiors
have important
astrophysicale↵ects on
protoplanetary
disk photoevaporation,such as the Taurus clouds explosive
and Orionevents
Nebula,may
typicallyof
fully
convective,
rapidly
rotating
stars,
followedbyionizationof diskgas, production of spallogenic radionuclides in disk solids, and hydrodynamic
escapeshow that highly variable X-ray
emission
is a ubiquitous
eruption of field lines and violent magnetic reconnection

1

Centro de Astrobiología (CSIC-INTA), ESA-ESAC Campus, 28692, Villanueva de la Cañada, Madrid, Spain
School of Physics and Astronomy, University of Leeds, Leeds LS2 9JT, UK
3
Joint ALMA Observatory, Alonso de Córdova 3107, Vitacura, Santiago 763-0355, Chile
4
Homogeneous
study ofVAHerbig
Ae/Be stars from spectral energy
National Radio Astronomy Observatory, 520
Edgemont Road, Charlottesville,
22903, USA
5
Spanish Virtual Observatory
distributions and Gaia EDR3
6
Departamento Física Teórica, Facultad de Ciencias, Universidad Autónoma de Madrid, Campus de Cantoblanco, 28049, Madrid,
Spain
J. Guzmán-Díaz1 , I. Mendigutía1 , B. Montesinos1 , R.D. Oudmaijer2 , M. Vioque2, 3, 4 , C. Rodrigo1, 5 , E. Solano1, 5 , G.
2

Received September 24, 2020; accepted April 9, 2021

Meeus6 , and P. Marcos-Arenal1

1

Centro de Astrobiología (CSIC-INTA), ESA-ESAC Campus, 28692, Villanueva de la Cañada, Madrid, Spain
School of Physics and Astronomy, University of Leeds, Leeds LS2 9JT, UK
3
ABSTRACT
Joint ALMA Observatory,
Alonso de Córdova 3107, Vitacura, Santiago 763-0355, Chile
4
National Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, VA 22903, USA
5
SpanishVirtual studied
Observatory
Context. Herbig Ae/Be stars (HAeBes) have 6so
far been
based on relatively small samples that are scattered throughout the
Departamento Física Teórica, Facultad de Ciencias, Universidad Autónoma de Madrid, Campus de Cantoblanco, 28049, Madrid,sky. Their fundamental stellar and circumstellar
parametersandstatistical properties were derived with heterogeneous approaches
Spain2

iv:2104.11759v1 [astro-ph.SR] 23 Apr 2021

Gaiaと⾚外SEDを使って
Herbig Ae/Beの星と星周のパラメーターを統計的、等質的
before Gaia.
Received September 24, 2020; accepted April 9, 2021
Aims. Our main goal is to contribute to the study of HAeBes from the largest sample of such sources to date, for which stellar andに調べる。
circumstellar properties have been determined homogeneously from the analysis of the spectral energy distributions (SEDs) and Gaia
ABSTRACTEDR3 parallaxes and photometry.
28%が遷移円盤を持ち、⽳はMethods. Multiwavelength photometry was Context.
compiled
forAe/Be
209 stars
bonafide HAeBes
EDR3
distanceswere
Herbig(HAeBes)have so farfor
beenwhichstudied Gaia
based onrelatively
small samples
thatestimated.
are scattered throughout the
Their fundamental
stellar andcircumstellar
parametersand statisticalproperties
were stellar
derived withheterogeneousBeが34%、Aeが15%。
approaches
Using the Virtual Observatory SED Analysersky.
(VOSA),
photosphericmodelswere fit
to the optical
SEDs
to derive
parameters,
before Gaia.
and the excesses at infrared (IR) and longer Aims.
wavelengths
werecharacterized
deriveseveralA tostatistical
Our main goal
is tocontribute to the tostudy
of HAeBes
fromcircumstellar
the largest sampleproperties.
of such sources
date, for whichAeは年齢が⾼いが⽳が⼩さい
stellar andanalysis was carried out to show the potentialcircumstellar
use of sucha largehave
dataset.
properties
been determined homogeneously from the analysis of the spectral energy distributions (SEDs) and Gaia
EDR3 parallaxes
and photometry.
Results. The stellar temperature, luminosity, radius,mass, andage were derived for each star based on optical photometry. In addition, ので光蒸発では説明できない。
Methods. Multiwavelength photometry was compiled for 209 bona fide HAeBes for which Gaia EDR3 distances were estimated.
their IR SEDs were classified according to twoUsingdi↵erent
schemes,andSED
theirmass (VOSA),
accretion
rates, disk
masses,sizesof theinner
the Virtual Observatory
Analyserphotosphericmodelswere fitandto thetheopticalSEDs
to derive
stellarBeは光蒸発と⽭盾しない。
parameters,
dust holes were also estimated uniformly. Theandinitialmass function
fits andthe longer
stellarwavelengths
mass distributionof thetosample
within
2 < M⇤ /M< A statisticalthe excesses
at infrared (IR)were characterized
derive several
circumstellar
properties.
analysis was carried
to show thepotentialusecanof suchlarge fordataset.
12. In this aspect, the sample is therefore representativeof theoutHAeBeregimeandbe ausedstatistical purposes when it is taken
Results. The stellar temperature, luminosity, radius, mass, and age were derived for each star based on optical photometry. In addition,
into account that the boundaries are not well probed.
Ourstatistical
study
does
not
revealany
connection
the SEDshapeandfrom
their IR SEDs were classified according to two di↵erent schemes, and their massbetween
accretion rates,
disk masses,
the sizes of the innerthe Meeus et al. (2001) classification and thedust
presencetransitional
disks, The
which
based
the SEDs
show
holes wereofalso
estimated uniformly.
initialare
massidentified
function fits here
the stellar
massondistributionof thethat
samplewithin 2 < M⇤ /M <In this aspect, theInsampleis therefore
the HAeBe
regimeand transitionalcan be used for statistical
purposesan IR excess starting at the K band or longer12.wavelengths.
contrast,
onlyrepresentative⇠28% ofofthe
HAeBeshavedisks, andthe when it is taken
into account that the boundaries are not well probed. Our statistical study does not reveal any connection between the SED shape from
related dust disk holes are more frequent in HBes
thanetin(⇠ 34% vsand
15%).
The relativelysmallholes andold stellar
the Meeusal. HAes(2001) classification
the presenceof transitional
disks,innerwhich disk
are identified
here basedon the SEDs that show
ages estimated for most transitional HAes indicate
thatstarting
photoevaporation
be the main
mechanism
drivingdisk dissipationin disks, and thean IR excess
at the K band orcannot
longer wavelengths.
In contrast,
only ⇠28%of the HAeBes
have transitional
aretransitional
more frequent in
HBes arethan consistentin HAes (⇠ 34%
vs 15%).
The relatively small innerdisk holes and old stellar
these sources. In contrast, the inner disk holesrelated
and dust
agesdisk
ofholes
most
HBes
with
the photoevaporation
scenario,ages estimated for most transitional HAes indicate that photoevaporation cannot be the main mechanism driving disk dissipation inalthough these results alone do not unambiguously
discardother disk
dissipation
mechanisms.
these sources.
In contrast,
the inner
disk holes andages of most transitional HBes are consistent with the photoevaporation scenario,
Conclusions. The complete dataset is available
onlinethrougha Virtual
Observatory-compliant
archive,
representing
althoughthese
results alone
do not unambiguously
discard other disk
dissipation
mechanisms. the most recent
Conclusions.
The complete
dataset is availabletool
onlinea Virtual
Observatory-compliant
archive,identified
representing the most recent
reference for statistical studies on the HAeBe regime.
VOSA
is a complementary
forthroughthe futurecharacterizationof newlyreference for statistical studies on the HAeBe regime. VOSA is a complementary tool for the future characterization of newly identified
HAeBes.
HAeBes.
Key words.
Protoplanetarydisksvariables:– Stars: pre-main
sequence– Stars:Ae/Bevariables:
Tauri, fundamental
Herbig Ae/Be – Stars:
fundamental paramKey words. Protoplanetary disks – Stars: pre-main
sequence– Stars:
T Tauri,
Herbig– TStars:param-

eters – Astronomical
eters – Astronomical data bases – Virtual observatory
tools data bases – Virtual observatory tools