星形成ゼミ20210205
SFN #335 15-21

片岡章雅

15. A new prolonged decrease event in the brightness of the young stellar object V2492 Cygni, Sunay
Ibryamov et al.
16. The Cloud Factory II: Gravoturbulent Kinematics of Resolved Molecular Clouds in a Galactic Potential,
Andres F. Izquierdo et al.
17. Evolution of the density PDF in star forming clouds: the role of gravity, Etienne Jaupart et al.
18. Star Formation Efficiency and Dispersal of Giant Molecular Clouds with UV Radiation Feedback:
Dependence on Gravitational Boundedness and Magnetic Fields, Jeong-Gyu Kim, et al.
19. The CARMA-NRO Orion Survey: Filament Formation via Collision-Induced Magnetic Reconnection –The Stick in Orion A, Shuo Kong, et al.
20. Probing the Cold Deep Depths of the California Molecular Cloud: The Icy Relationship between CO andDust, John Arban Lewis et al.
21. Characterization of very wide companion candidates to young stars with planets and disks, F. Z. Majidi, etal.

20. Probing the Cold Deep Depths of the
California Molecular Cloud: The Icy Relationshipbetween CO and Dust, John Arban Lewis et al.
We study the relationship between molecular gas and dust in
the California Molecular Cloud over an unprece- dented
dynamic range of cloud depth (AV = 3 – 60 magnitudes). We
compare deep Herschel-based measure- ments of dust
extinction with observations of the 12CO, 13CO, and C18O
J=2-1 lines on sub-parsec scales across the cloud. We directly
measure the ratio of CO integrated intensity to dust extinction
to derive the CO X-factor at over 105 independent locations in
the cloud. Confirming an earlier study, we find that no single
12CO X-factor can characterize the molecular gas in the cold
( Tdust ≤ 20) regions of the cloud that account for most of its
mass. We are able to derive a single-valued X-factor for all
three CO isotopologues in the warm ( Tdust >25 K ) material
that is spatially coincident with an H II region surrounding thestar LkHα 101. We derive the LTE CO column densities for
13CO and C18O since we find both lines are relatively
optically thin. In the warm cloud material CO is completely in
the gas phase and we are able to recover the total 13CO and
C18O abundances. Using CO abundances and deep Herschelobservations, we measure lower bounds to the freeze-out of
CO onto dust across the whole cloud finding some regions
having CO depleted by a factor of > 20. We construct the firstmaps of depletion that span the extent of a giant molecular
cloud. cores and discuss their physical nature.

アブストラクト
Av=3-60の大きなダイナミックレンジで、California
Molecular Cloudを観測した。
観測データ：ハーシェルのダスト減光と、12CO, 13CO,
C18O
目的：10^5箇所の独立な場所でCO X-factorを導出した。
結果：
Tdust < 20Kの冷たい領域では、CO X-factorは単純な1パラメータではない。
Tdust < 25K の暖かい場所では、single valueのCO Xfactorを導出した
LTEを仮定して、13COとC18Oの密度を出した。光学的に薄い。温かい領域では、COは全部気相にあると仮定して良さそう。
COの凝結についても調べた。場所によっては>20くらい
depletionしてる。

678

C. Lee et al.

イントロ
通常X-factorは12COを用いて、GMCの比較的面密度の低いところ(Av<4-5)で求められている。
先行研究：Kong et al.. Av~35までについてX-factorを出した。Tex>18Kのところではsingle X-factorを導出した。ただし、これらはHII regionのある場所と一致している。
この研究は Kong et al. 2015の拡張。まず、ハーシェル使ってAv=3-60のダイナミックレンジで観測。COで観測した領域は3倍大きい。星形成が活発じゃない領域も含まれる。温度はハーシェルを使った。

et al.
2014 clouds listed in Table 1. (left) The average relation calculated by combining the relgure 3. ICO as a function of AV combining theLee
24 localmolecular
the individual clouds. Each grey point here shows a bin from an individual cloud; the green profile shows the mean and scatter among the profiles

he California molecular cloud for pixels with SNR>3 for 12 CO and 13 CO,

観測
Arizona Radio Observatory (ARO) 10 m Heinrich-Hertz
Sub-millimeter Telescope (SMT), とハーシェル
PACS 160 μm and SPIRE 250, 350, and 500 μm をつかったが、分解能は500µm似合わせて36秒角。
H2の出し方
modiﬁed BB でフィット。τ850を出して、それを適当にA̲Kになおして(Lombardi 2009)、それをN[H2]/Av の値を古典的なものを仮定して(Bohlin 1978, Rachford 2002)出した。最終マップは38ʼʼ分解能、0.08pcに対応（天体までの距離は450pc）
結果、ハーシェルから導出した質量は5.52 10^4 M⊙ で、
12

Figure 3. Binned 2D histogram for W[ CO (2 - 1)] vs AV for theentire California cloud where W[12 CO] is detected at greater than
3 (& 1K). The gray dashed line is the standard relation for the
adopted MW X-factor (2 ⇥ 1020 cm-2 (K km/s)-1 ) which is 3.29
K km/s mag-1 . The blue histogram shows the number of pixels in abinned pixel of the plot on a log scale. The black dashed box shows
the region containing the bulk of data range covered by previous
studies (e.g., Pineda et al. 2008; Ripple et al. 2013).

On a plot of WCO vs AV , a single valued X-factor is represented by a straight line through the origin. The 2-D distribu-

クラウド全体の20%くらい。
CO観測について
The average RMS noise per 0.3 km/s channel weachieve is 0.11 K for C18O, 0.13 K for 13CO, and 0.11
K for 12CO.

T HE C OLD D EPTHS OF C ALIFORNIA

3

dent, with the 10 K km/s branch associated with the coldest
dust, the 35 K km/s branch associated with slightly warmer
dust, and the vertical branch with dust over ⇠ 25 K. The cold
branches extend to low extinctions beneaththe points cor/
responding to the hot dust in red; however 98% of the cold
CMC-West
pixels are below 40 K km/s. /
Cal-X
Further inspection of the data shows that these branches/N+Ÿ
actually correspond to different
regions of the cloud. The
1*&
nearly vertical branch corresponds to CO associated with hot(Tdust > 25 K) dust located in the vicinity of, and coincident
with, the H II region in L1482. This is the region studied by
K15 who found similar behavior in the W[CO]-Av relations
with CO excitation temperature, in particular, K15 showed
Figure 1. CO survey outlines over Herschel +Planck dust maps. The CO survey boundary is outlined in black and red, with the black outlinedthat theCOemissionin extent
the ofvicinityof initiallythe Hpresented
II regionis thealso
covering
L1482,
the southeastern
the cloud, beingin K15 andred outlines indicating the new data presented inthis paper (L1478, and CMC-West). CO and CO were observed in all 3 regions, while C O was not observed in CMC-West. The jagged
characterized
by thehightemperatures.
The
upper
black boarder marks
extentgasof theexcitation
Herschel dust data
(Lada et al. 2017), with
Planckdata used to fill in outside that. The Herschel data
continues further west about 1 deg. than shown on the map. The CO survey boarder was chosen to cover the majority of the A >3 mag dust.
horizontal
branch in the figure originates in the cooler, more
Two subregions of interest are shown. The dashed white circle shows the extent of the H II region (NGC 1579) associated with the star (shownin blue) LkH↵
In CMC-West
we show thefeature describedin Imarawithin
et al. (2017).
extendedstar101.
forming
regions
inCal-X
L1482
and fromthe
6 (SMT) using a
L EWISET AL . of
the Heinrich Hertz Submillimeter Telescope
We calibrated the intensity using periodic
observations
Cal-X
region in CMC-West. The lowest (⇠10 Kkm/s) horthe CO bright source W3(OH) (more details on the calibraprototype ALMA band 6 dual-polarization sideband separating receiverin combination
MHz – 256 chantion
method12canbefound in Bieging et al. (2010)). Since theizontalbranchin the with
Fig.the 0.25
4 consistsof the(T
Tablecoldest
1. Global
COdustX-factornel filterbank as the backend (0.25 MHz ⇠ 0.32 km/s at 230
beam and velocity resolution for each CO isotopologue are< 14K)
material
which
we
find
to
mostlyoriginateinthe the- 1)]GHz). We used two observing setups - (1) with 12 CO in theMethod all slightly different,
maps K15.
were convolved to the resoluX[12 CO (2131200
upper-sideband
and CO in thelower-sideband
andCMC.
(2) antion of the CO map (38 ) and regridded in velocity to 0.3
L1478
and 12CMC-Westregionsof18 thehN[H2 ]/W[12 CO]i
5.05 ± 1.80
4.4
other with CO in the upper-sideband and C O in the lowerkm/schannelstofacilitatecomparison.
The average RMS
1212hN[H2 ]i/hW[
CO]i
4.18
±channel
1.83 we 3.6
Thiscomplexstructureis quitedifferent
what
ofsideband.
This improves
the signal-to-noiseof theCO data. fromnoise
per 0.3is
km/s
achieve is 0.11 K for C18 O,
13 2.98 ± 1.74
Linearemission
fit throughorigin
0.13
K for
CO,
and 0.11 K for 12 CO. The average spectra
CMC wasmultiplewheredays during
the November
ten seen
in observed
other overclouds,
the
CO
at
low
ex2012 - April 2013 observing season. Those segmentsarefor each region4.1±0.8
are shown in Fig. 2. The reduced maps are
Average1
tinctions(i.e.,
1 overlayed
 AV onPlanck
3-4 +Herschelmagnitudes)or less
linshown as
outlines
dust map more
publicly
availableon the Harvard
Dataverse(Lewis
2020)
.
12
OTE
—Measurements of the global
CO (2in Fig. 1. CMC West was not observed in C18 O. N
The
surveyearlyboundaries
increases
with A , and then appears
to flatten
higher
3. DATAANALYSIS
for the atwholecloud
in unitsofwere designed toVcover regions with AV1)& 3X-factor
mag
-1
20
-2
10
cm
(K
km/s)
measurement
methods
3.1. CO Moment
Maps
in
the
Lada
et
al.
(2009)
extinction
maps.
(AV > 4 magnitudes) extinctions, though previous. The
studiesare discussedin §4.2. The K15 are scaled by 1.4
We observed in ’on-the-fly’ (OTF) mode (Mangum
et al.
We derivedettheal.
line parameters (integrated intensity, cenrarely
extend
much
into
the
high
extinctions(Frerking0
0
00-1
from the original XCO (1-0) to XCO
2007) - raster scanning 10 ⇥ 10 tiles at a rate ofto10converts , them tralvelocity, and velocity dispersion)using moment
analyPPP(2-1).
sampling
every
0.1 s. The
raw OTF data
were
put
T v v
T (v-M )v
1982;
Ladaspectra
et
al.
1994;
Lombardi
et
al.
2006;
Ripple
et
al.
sis:
M
=
T
v,
M
=
,
M
=
,
0
mb
1
2
M
M
00
onto a 10 grid using custom data reduction scripts in thewhere v is the channel width. M0 is the integrated inten2013;
Kong
et al.
2015;
Lee
al.and2018).
GILDAS
CLASS
softwarepackage(Petyet2018)
then had Indeed, none ofsity p(WCO ), M1 is the intensity weighted velocity centroid,a linear baseline
subtracted.
Those reduced
thenthe previous
studiescomeclose data
to wereachievingthe Mdynamic
andcorresponds to an intensity weighted velocity dis12

13

18

V

Figure 4. W[12 CO(2 - 1)] vs AV for the entire California cloud
where W[12 CO] is detected at greater than 3 (& 1 K). The points
are colored by their Herschel-derived dust temperature (Tdust ). The
points are plotted so that hotter points are plotted over cooler points.

vey bring the overall behavior of the W[CO]-AV relation into
sharper focus.
ハーシェルの温度で色分けしてみると、温かいところ
Confirming the earlier results of K15, it is clear from Fig.
3 that the CMC data are not well described by a single Xは線形な関係がありそう。
factor, i.e. there is no single, clear linear relationship present
25K以上のところはH
II
region
in
L1482
と一致。
12
between CO and AV . Typically, if there is a constant COabundance in a molecular cloud, a linear rise of W[12 CO]with AV is expected at low extinctions with a break at around

mb

mb

0

1

0

2

8

part of the cloud where CO is likely entirely in the gnamely in the region containing the hot dust. We
the total CO abundance by measuring the average ab
hNCO /N[H2 ]i, in the hot (Tdust >25 K) dust, using pi
13
18
signal-to-noise
>3
for
CO
and
>5
for
C
温かい領域 Tdust>25Kで13COとO and T
K, and find C18Oのアバンダンスを測った結果

L EWIS ET AL .

13

[ CO] = (1.31 ± .89) ⇥ 10
18

[C O] = (1.27 ± .59) ⇥ 10

-6-7

これはそれぞれstandard
These measurements
are consistent atomic
within the err
abundances abundances
derived from ratioですべてのcarbonがstandard atomic abundancassuming allCOになっていると仮定したときの値と
the carbon is in CO2 : [13 CO]MW = 1.4
18-7
consistentらしい[13CO]MW
[C
O]
=
1.80
⇥
10
. Looking at the=hot pixels
Figure 6. Plot of the (Left) CO and (Right) C O column density vs visual extinction colored by T . The colorbar shows colder pixels
MWas blue, transitioning through gray, and becoming redder as the temperature increases. Data is plotted so that higher temperature points are
1.45
10^­6,
[C
O]MW
=
1.80
10
that
a
straight
line
through
the
data
would
not
inte
plotted over cooler points to to emphasize the temperature dependent structure in the plot. The inset shows W vs A for the line with the samecoloring scheme.
^-7
origin. By performing
a linear regression to the dat
We see in Fig. 6 that again, as with W[12 CO]–AV , only
both emission (e.g.; towards the Pipe (Lombardi et al. 2006)
determinethelinearrelation
between
CO
and
H
.
2
and Cygnus (Schneider et al. 2016)) and absorption (e.g.; tothe hottest dust shows a clear linear relation. To measure the
13

18

dust

CO

V

13COとC18Oで同様の図を書いて比較。
true CO abundance, we must limit our measurements to thewards O & B stars, (Burgh et al. 2007)), so this offset is notpartcoldof the cloud
where CO is likely entirely in the gas phase,a threshold for CO formation.
brancheでは、見積もった光学的厚さは、13CO
namely in the region containing the hot dust. We measure
4.3.2. Abundance versus Tdust
1.5,
C18O
the が
total︎ COabundanceby は光学的に薄い(τ18measuring the average abundance,< 0.75)

1315N[ CO] = 3.19 ⇥ 10 (AV - 2.50)
N[C18 O] = 2.47 ⇥ 1014 (AV - 3.55)

hNCO /N[H2 ]i, in the hot (Tdust >25 K) dust, using pixels with
As with WCO vs AV , NCO vs AV shows a continuous trancold branchが３つくらいに分かれているのは場所後がsignal-to-noise >3 for 13 CO and >5 for C18 O and Tdust > 25
sition from the cold horizontal branches to the hot verticalThese
equations can be combined with equations
K, andfind
い。L1482とL1478
branch, indicating that the abundance may be tied
to the dust
temperature. Plotting the bin-averaged abundance as a func13-6(7) to derive a linear relation for WCO –AV . The A[ CO] = (1.31 ± .89) ⇥ 10
(8)tion of Tdust binned in 0.05 dex bins shows us that the abun-

T HE C OLD D EPTHS OF C ALIFORNIA

9

T HE C OLD D EPTHS OF C ALIF

Figure 8. Schematic drawing showing how the depletion factor isan average of the depletions along the line-of-sight and necessarily
underestimates the depletion in the center regions.

Figure 7. Binned log10 median abundance for (Top) 13 CO and (Bottom) C18 O versus dust temperature. The abundance is binned byTdust in logarithmic bins from 10 K - 63 K in 0.05 dex steps. The
error-bar is the median absolute deviation scaled to the standard deviation. The purple line with the purple shaded region marks theabundance measured in the hot dust and the corresponding 1 error.
The red dashed line is the cosmic abundance

et al. 2014) and was observed in the CMC by K15. Alternatively, the difference between our measured and the cosmic

ures 11,and C18
figures
contourcolumnthat the(see Apspatially
tors rantion fac
area coi
The mpletionobservestrongethan ⌘1
Both mpletiongeneralseen inIn their study of the L1482, K15 suggested desorption/depletion as the cause of abundance trends with Tex .
However, depletion depends on the dust temperature which
they were unable to make direct comparisons too at the time.
The temperatures we measure in the cold dust are significantly lower than the sublimation/depletion temperature andCO should freeze out onto grains in the cold, dense central
regions of the cloud. Similar variations in the nature of the18
10.Orion
Plot byof thetotalet(gasN[13 CO]–AV relation were Figure
found in
Rippleal. + ice) C O column density plotted againstthat(AVdepletion- 3.55) mag.
(2013) who came to the conclusion
was There- points are colored by Tdust - noteis different
than AinV Figs. 4 and 6. The red line showssponsible for the flatteningthatof the scale
relationshipbeyond
18
18
-7
the
N[C
O]
predicted
using
theabundance,
[C
O]
=
1.27
⇥
10
⇠ 10 mag.
18
ignoring
thematerial
below
the
3.55
mag
threshold,N[CO] /Our observations of the relationship between CO (WCO ,18[Cshaded region shows the 1 uncerV - 3.55).
NCO ) and the dust point to itO](Abeing
shaped The
by gray
desorp18
taintyintheC
O abundance.
The total column derived from addingtion/depletion processes and are consistent withprevious obcontribution
from iceis theconsistent with the total column denservations of depletion (e.g.;ininthe
Orion).
We conclude
thatsity derived
abundance
We pspatial variation in CO abundancein thefrom
CMCtheis duein largemeasured in the hot dust.
a dendrpart to changes in desorption and depletion caused by thedust temperature 冷たい領域については、CO氷とCO2氷の量を見積もっchanging plot
with indistancestar forming
and its u
red thefrom
line that
corresponds to the abundance we mearegions in the cloud.
In shorsured in the hot dust. The gray shaded region shows the errorてたしてやれば、全COの量をだいたいリカバーでき

est valuin the abundance. For this hot material we assume that all
5. DEPLETIONる。ここで、COとCO2iceの見積もりはAvから経験的
dent pethe CO is fully evaporated from the grains and thus that thisIn this section we measure the level of depletion across theand othmeasured abundance represents the true CO abundance in theに見積もられている先行研究を使っている。
CMC and create the first maps of depletion on giant molecustructurCMC.
lar cloud scales.

18. Star Formation Efficiency and Dispersal of Giant Molecular Clouds with UVRadiation Feedback: Dependence on Gravitational Boundedness and Magnetic
Fields, Jeong-Gyu Kim, Eve C. Ostriker,1 and Nina Filippova
Molecular clouds are supported by turbulence and magnetic fields, but
quantifying their influence on cloud lifecycle and star formation efficiency(SFE) remains an open question. We perform radiation magnetohydrodynamic
simulations of star-forming giant molecular clouds (GMCs) with UV radiation
feedback, in which the propagation of UV radiation via ray-tracing is coupled tohydrogen photo- chemistry. We consider 10 GMC models that vary in either
initial virial parameter (1 ≤ αvir,0 ≤ 5) or dimensionless mass-to-magnetic flux
ratio (0.5 ≤ μΦ,0 ≤ 8 and ∞); the initial mass 105M⊙ and radius 20pc are fixed.
Each model is run with five different initial turbulence realizations. In mostmodels, the duration of star formation and the timescale for molecular gas
removal (primarily by pho- toevaporation) are 4-8Myr. Both the final SFE (ε∗)
and time-averaged SFE per freefall time (εff) are reduced by strong turbulence
and magnetic fields. The median ε∗ ranges between 2.1% and 9.5%. The median
εff ranges between 1.0% and 8.0% and anticorrelates with αvir,0, in qualitative
agreement with previous analytic theory and simulations. However, the timedependent αvir(t) and εff,obs(t) based on instantaneous gas properties andcluster luminosity are positively correlated due to rapid evolution, making
observational validation of star formation theory difficult. Our median εff,obs(t)≈ 2% is simi- lar to observed values. We show that the traditional virial
parameter estimates the true gravitational boundedness within a factor of 2 on
average, but neglect of magnetic support and velocity anisotropy can sometimes
produce large departures. Magnetically subcritical clouds are unlikely torepresent sites of massive star formation given their unrealistic columnar
outflows, prolonged lifetime, and low escape fraction of radiation.

輻射磁気流体シミュレーションをやった。
UVはHydrogen photochemistryとカップルして解いた。
10 GMCモデル。ビリアル比を1-5、mass-to-ﬂux ratioを0.5-8, 無限大、でふった。
初期質量は10^5太陽質量、半径20pcは固定した
どれもだいたい4-8Myrで星形成が終わってガスが飛んでった。
median (時間平均した自由落下時間で割った星形成率) ε̲ﬀは、ビリア
ル比と反相関。
subcriticalクラウドは大質量星形成のサイトとしては不適切そう。

dance relative to the hydrogen nucleus. Each species
we follow is passively advected with the velocity field,with the right-hand side of Equation (4) representing
the net creation rate Cs due to various collisional reactions, cosmic ray ionization, and photodestruction processes. We explicitly follow the non-equilibrium abun数値計算について
dances of hydrogen in molecular (H2 ), atomic neutral
0
+
(H
),
and
ionized(H) phases but assume equilibrium- ベースはAthenaを利用。磁気流体パートは普abundances for carbon- and oxygen-bearing species (C0 ,通。ChemistryはH2,
C+ , CO, O0 ) (see SectionH0,
2.3).H+を非平衡ちゃんと扱
Adopting the ideal gas
law, assuming
0.1, and ignoring the abundance
He =Oは平衡状態を仮定って、C,C+, xCO,of trace species, the gas temperature can be written as- Tたぶんこの人達が頑張っているのは輻射パート= P/[(1.1 + xe xH2 ) kB nH ].
In addition to Equations (1)–(6), we solve radiativetransfer equations in the formn̂ · rI =I(7)

where I is the intensity,  is crossection per unit volume
in a given radiation component, and n̂ is the direction of
ray propagation. The scattering and emission by dust
grains are ignored. The radiation field is decomposed
into di↵use background and starlight: I = I bg + I ⇤ .
We use the six-ray approximation to solve for the diffuse background, meaning that only n̂ aligned with the- 輻射を２つに分け、diﬀuse backgroundとstarlight radiationを別々にちゃんと解いた。

age tage (e.g., Kim & Ostriker 2017). Using the stellar population synthesis model STARBURST99 (Leitherer et al. 2014), the time-dependent radiative output
per unit massis calculated in three frequency bins:(1) Photoelectric計算セットアップのパラメータのフリ方(PE; 110.8 nm < < 206.6 nm); (2)Lyman-Werner (LW; 91.2 nm < < 110.8 nm); (3) Ly- 輻射磁気流体シミュレーションをやった。乱流・磁man Continuum (LyC; < 91.2 nm). The PE and LW
clusterからのUV
feedbackあり
constitute FUV 場ありGMC、star(or non-ionizing) radiation
and are absorbed mainly by dust. Absorption of the FUV photons
初期：M0=
10^5Msun,radius
R0
=
20pc,
and
by small grains ejects electrons, which are important for
heating neutral gas.
The LW
photons
photodisfreefall
time
tﬀ ,0include= 4.68Myr.
sociation bands of H2 and CO molecules and also ion7
ize C0 . The LyC (or ionizing) photons are responsible
Table
for ionizing H0 (h⌫>1. Model
13.6parameterseV) and H2 (h⌫ > 15.2 eV).
B
M
M
We assumeModelthat↵ theµ mean energyof ionizing photons[km s ] [µG]is h⌫LyC = 18 eV. In Appendix
A, we show (tage ) inparameters(1)(2) Table(3) 1. Model
(4)
(5)
(6)
(7)
運動エネルギー高い
di↵erent frequency bins.
↵-series
Model ↵µB
M
M
We calculate
the
cross
dust absorption
A1B2
1.0
2.0
2.1 sections
13.5 13.6 for
1.3
A2B2
2.0
2.0
2.9
19.2
1.9
[km s 13.5
]over
[µG]
and photoionization
averagedthecluster’s UV specA3B2
3.0
2.0
3.6
13.5 23.5
2.3
(1)
(2)
(3)
(4)
(5)
(6)
(7)
trum and find
that
A4B2
4.0
2.0the4.1 dependence
13.5 27.2
2.7 on tage is weak.
A5B2
5.0
2.0
4.6
13.5 30.4
3.0↵-series
We thus takeconstant
cross
sections
PE = nH d,PE
-series
1.0
2.0
2.1
13.5 13.6
1.3
21
2
1
and LW =A1B2
n
,
where
=
10
cm
H
A2B05 H 2.0d,LW
0.5
2.9
54.0 19.2
d,PE0.5
A2B2
2.0
2.0
2.9
13.5 19.2
1.9 Figure 1. Parameters of simulated clouds in ↵- and 21
2
1
A2B1
2.0
1.0
2.9
27.0 19.2
0.9
series modelsvalues(see also Table 1). All clouds have the sameand d,LW A3B2
=
1.5
⇥
10
cm
H
.
These
3.0
2.0
3.6
13.5
23.5
2.3
mass M = 10 M and radius R = 20 pc. Thepbottom xA2B2
2.0
2.0
2.9
13.5 19.2
1.9
A4B2
4.0
2.7 axis is the
mass-to-magnetic flux ratio µ= 2⇡ G⌃ /B
are appropriate
for4.02.0
Weingartner&(2001a)’s
A2B4
2.0
2.9 4.1 6.7 13.5
19.2 27.2
3.8 Draine
=
A2B8
2.0
19.2 30.4
7.6
A5B2
5.0 8.02.0 2.9 4.6 3.4 13.5
3.0 and the left y-axis is the kinetic virial parameter ↵grain model
with
R
=
3.1.
For
LyC
radiation,
we
5R/(GM ). The top x-axis and right y-axis showV
A2Binf2.01
2.9
0.0 19.2
1
vir,0

,0

z,0

1d,0

0

A,0

z,0

0

1

vir,0

,0

1d,0

7

A,0

1

0

5

0

,0

0

z,0

vir,0

2
0 1d,0-series

磁場強い0

the corresponding magnetic field strength Bz,0 and onedimensional turbulent velocity dispersion 1d,0 , respectively.
Figure 1. Parameters of simulated clouds in ↵- and For each pair of (↵vir,0 , µ ,0 ), we run five simulations withseries seeds
models(see also
Table1). Alldi↵erent random
for initial
turbulentvelocity
field.clouds have the same

Note—Eachmodel is0.5
run with 2.9
5 di↵erent
randomseeds 0.5
A2B05
2.0
54.0
19.2
for turbulence. All models have the same initial mass
5 2.0
A2B1
1.0 R0 =2.920 pc. Columns
27.0 19.2
M
and radius
are as 0.9
0 = 10 M
follows:
model namemass M0 = 105 M and radius R0 = 20 pc. Thepbottom xA2B2 (1) 2.0
2.0 indicating
2.9 the initial
13.5 virial
19.2pa- 1.9
rameter (A) and magnetic flux-to-mass ratio (B). (2) iniaxis is the mass-to-magnetic flux ratio µ ,0 = 2⇡ G⌃0 /Bz,0
A2B4
2.0
4.0
2.9
6.7 19.2
3.8
2
tial virial parameter ↵vir,0 = 5 1d,0 R0 /(GMnote that andwe havethe e↵ects
the initial
tem-parameter ↵p 0 ). (4) inithe tested
left y-axis
is theof kinetic
virialvir,0 =
A2B8
2.0
8.0
2.9
3.4
19.2
7.6
tial magnetic mass-to-flux ratio µ ,0 = 2⇡ G⌃0 /|Bz,0 |.
perature choice, 2and found that they are unimportant as
5Rionization
Thearetopx-axisand right y-axis show
0 1d,0 /(GM0 ).
(3) initial velocity
dispersion.
(5) initial0.0
magnetic
field 1
thermal andequilibrium
rapidlyachieved.
A2Binf
2.0
1
2.9
19.2
thethecorresponding
magnetic
fielda passive
strength Bz,0 and oneWe assign
gas that is initiallymolecular
strength. (6) 3D sonic Mach number of initial turbuNote—Eachis runnumber
with of5 di↵erent
random seedsscalar fielddimensional
value scl = turbulent
1 and the velocity
ambient dispersion
atomic gas 1d,0 , respectively.
lence. (7) 3D model
Alfvén Machinitial turbulence.

- 簡単な結果- すべてのモデルにおいて、UVfeedbackは良く効いて、SFEを下げた。磁場や乱流が強いと更に下がる。

0. each
We pair
use sofdi↵erentiate
the run“cloudcl tofor turbulence. All models have the same initial masswith scl =For(↵vir,0, µ ,0 ), we
five simulations withmaterial”
fromtheambient
medium,but
itshouldbe
M0 = 105 M and radius R0 = 20 pc. Columns are asdi↵erent
randomseeds
forinitialturbulent
noted that the use of a passive scalar cannot serve as a velocity field.
follows: (1) model name indicating the initial virial pa-perfect surrogate for the Lagrangian tracer particle due
2.5.(A)
Model
and Parameters
rameterand Initialization
magnetic flux-to-mass
ratio (B). (2) ini-to numerical di↵usion. For analysis, we select gas that is2
neutral andinitiallyusing athefiltere↵ects
functionof the initial temWe virial
initialize
the simulationtial
parameter↵vir,0by=placing
5 1d,0aRuniform
0 /(GM0dennote
thatmolecular
we havebytestedp ). (4) ini-

SFE = 2.1-9.5% (α), 2.4-8.4% (β)-

sity cloud with mass M0 = 105 M

and radius R0 =(

36

- ビリアル比ふったやつ- 磁場強度ふったやつ

37- 全体の進化について
1. フィラメントができる。 2.sink particleができて、フィラメント内にHII 領域ができる。 3.
HII領域が壊れたり合体したりして、光蒸発を促したり、冷たい中性ガスを放出して、星形成を終わらせる。

- 乱流が強いと(αvir>2)
クラウドが広がり、
OUTFLOWで中性ガスを失う(Fig.17)- 強い磁場だと(μ<1)フィラメントは初期磁場の方向に垂直になる。
動きは非等方になる
18. Same as Figure 17 but for the -series modes with seed = 4, which have identical initial turbulent velocity field but
Figure 17. Snapshots of gas surface density projected along the y-axis from the ↵-series models with seed = 4 From topFiguretovarying initial magnetic field strength, decreasing from top to bottom. The corresponding initial mass-to-magnetic flux ratio
bottom, the initial virial parameter ↵increases from 1 to 5. The projected positions of star particles are indicated by circles.
(Fig.18)µ
varies from 0.5 to 8.
The size and color of circles indicate mass and age, respectively.
vir,0

,0

(2018). There, we focused on how evolution and final outcomes depend on the cloud surface density (for
6. Virial Ratios. We study the evolution of the viriala range ⌃0 ⇡ 10 103 M pc 2 ), and found that inratio ↵
˜ vir,tot that accounts for kinetic, thermal, andall cases star formation is quenched and clouds are dis- ビリアル比の時間進化
persed within a few freefall times or several Myrs aftermagnetic energies of the neutral cloud, and compare- 直接計算したビリアル比と、これまでよく使われていたビリアル比の見
star formation commences. Our new simulations, with
it with the traditional kinetic virial parameter ↵vir =2
more realistic physics (addition of realistic thermochem5R
on the cloud mass, size, and ve1d /(GMneu ) based
積もりαvirを比較した。
istry, magnetic fields, and time-dependent luminosity)
locity dispersion (Figures 5 and 13). The virial ratio- ビリアル比は、乱流が下がるため一旦下がって、それから上がる
confirm this evolutionary scenario. The short timescales
decreases due to decaying turbulence at early times,
we find are also consistent with recent empirical conbut it increases to large values (˜↵vir,tot > 10) at late
（feedback）。
straints on rapid cloud dispersal after the onset of startimes as feedback disperses the cloud. In spite of the- 初期条件ふったが、最初の星形成の頃に0.5‒1.7に落ち着く。磁場が強formation (e.g., Kruijssen et al. 2019; Chevance et al.
widely ranging initial conditions of the models, at the
いと別。 time of the first star formation ↵˜ vir,tot falls in a fairly
2020b,c).
narrow range 0.5–1.7. In strongly magnetized clouds,
Our findings of decreasing SFR and net SFE with in磁場があると、基本的に速度分散から計算したαvirを使うと、特に磁場
the kinetic energy may be significantly underestimatedcreasing ↵vir,0 are consistent with findings from previが強いときにビリアル比を過大評価してしまう。
or overestimated depending on the relative orientation
ous hydrodynamical simulations of global clouds (e.g.,
between the observer’s line of sight and the direction of
Raskutti et al. 2016; Bertram et al. 2015; Howard et al.
magnetic field. Averaging over directions (as would be
2016; Dale 2017; Kim et al. 2018), as well as hydrodyin fesc,Lyc than in fesc,FUV .

Figure 5. Time evolution of cloud size, velocity dispersion, energy, and virial ratios for three di↵erent models A2B2S4, A5B2S4,and A2B05S4 until t⇤,90% , when 90% of the final stellar mass has been formed. The grey dashed vertical lines indicate the time offirst star formation. (Top row) The e↵ective radius of neutral cloud R = 21/3 R50% , where R50% is the half-mass radius. (Secondrow) Velocity dispersion measured perpendicular ( 1d,? ) and parallel ( 1d,k ) to the background magnetic field, and averaged
over all directions ( 1d ). (Third row) Gravitational (|W|, blue), kinetic (Tkin , red), thermal (Tthm , green), and magnetic (B,yellow) energies (see Equations (21)–(24)). A simple estimate of gravitational energy is also shown (blue dashed). (Fourth row)
2
The kinetic virial parameter ↵vir = 5 1d
R/(GMneu ) for di↵erent viewing angles (gray lines) and various energy ratios.

the density-dependent optical depth isd nH Lshld , this corresponds to the shielding length of

Figure 13. Comparison of the virial ratio ↵˜ vir,tot = (2T + B)/|W| and the traditional kinetic virial parameter2↵vir = 5 1d
R/(5Mneu ) in the ↵-series (left) and -series (right) models. Each circle represents one snapshot anindicates the time of the snapshot normalized by the time of the first star formation. The grey circles mark values
The size of circles increases with increasing ↵vir,0 and Bz,0 (decreasing µ ,0 ). Rows show the traditional ↵vir calcu
velocity dispersion perpendicular (top) and parallel (middle) to the z-axis, and averaged along all three Cartesian axe

we use the tilde to emphasize the distinction.
The third row in Figure 5 plots the evolution of ki-

3GM02 /(5R0 ) at early times. At late times (tit begins to decrease as gas is photoevaporate

26

25

3.9.1. Comparison with Simulations and Theory
integral in some models to account for the freefall time27

of gasat di↵erent densities.
In Figure 14, we plot "↵,0 of all models against
thestructuresrap
We compare our results with the analytic models of
1/2
2
R0 / Oftio t↵,0 /tdyn,0 = ⇡ /40↵vir,0 , where tdyn,0 ⌘SFR.
1d,0
the six models presented in Table 1 of Federis the dynamical time (or turbulent crossingrathtime)
at& Klessen(2012), we consider the Krumholz & Mcmodel without the multi-freefall factor (KM), andt = 0. We perform a linear square fit of theKee↵-series
version of the Hennebelle & Chabrier星形成率のfreefall time依存性
models to "↵,0 = exp ( Ct↵,0 /tdyn,0 ) and find thethemulti-freefall
bestmodel (multi-↵ HC), both of which are extended to infit value C = 4.06, which is shown as a grey solid
line.
clude the e↵ect of magnetic field. The yellow and green
The empirical relation "↵,0 / exp( Ct↵,0 /tis
乱流弱いと高い、磁場弱いと高dyn,0 ) regions
shadedin Figure 14 show the predictions byqualitatively consistent with the star formationthe
law
proKM
and multi-↵ HC models, respectively. In both
posed by Padoan et al. (2012), who carried out
a set weof have assumed fixed values of b = 0.5 andmodels,い✏core / t = 1, and author-favored values of fudge facadaptive-mesh refinement simulations of star formationx =with driven turbulence in a periodic box. For ators
range
of 1.12 for the KM model and ycut = 0.1 forそれを観測と比較すると垂直に交
the multi-↵
3D sonic Mach number (10 and 20), Alfvén Machnum- HC model, respectively (see Federrath &2012). We also take = 0.02 and a range ofber (1.25–33), and the timescale ratio 0.23 < tKlessen↵ /tdyn <sonic Mach number 13.6 < M < 30.4 consistent withわってる。それを擬似観測で再現
1.78 (or 0.2 < ↵vir < 13), they measured "↵ thefromtheparametersof the ↵-series models.11
slope of a linear fit between the M⇤ (t)/M0 andThet/t↵,0.
qualitative
prediction by these analytic modelsできた
They found that "↵ has weak dependence
M andwith our simulation results, although the quanp onagrees
resultsdi↵er. Both models predict that ⇢crit
MA and is well fit by a relation exp( 1.6 3t↵titative/tdyn )10,
輻射でクラウドが大きくなる・
increases
linearly
with
↵
(for
fixed
other
dimensionvir
which is shown as a gray dashed line in Figure 14. Whileless parameters),so "↵ decreases with increasing ↵vir .
our best-fit value C = 4.06 is steeper than 2.77found
質量を失うので、R/Mneuが増
Higher overdensity is required for collapse at higher
by Padoan et al., a direct comparison is difficult
consid↵vir because the relative importance of turbulent stress,
ering the ambiguity in defining the dynamical whichtimescaletends to tear apart structures, versus gravity inえる。速度分散も増える。
and/or the gravitational energy in local simulations
withcreases.
In the Krumholz & McKee model, the depenperiodic boundary conditions. The main pointis thatdence
on Mach number is weak, because an increase
1/2
in Mholds
increases both the critical density (⇢crit / M2 )a similar exponential dependence of "↵,0 on ↵vir,0width of the density PDF, whose e↵ects tendtrue in global as well as local simulations, andand
fortherealto cancel each other. In the Hennebelle & Chabrieristic feedback as well as driving in k-space. model, however, ⇢ / M 2 and " increases monocrit↵
To explain the low observed SFR, a class oftocallytheoret(and sensitively) with increasing M.
2
1/2 1/2
ical 15.
modelshave
been
proposedt↵todescribe
FigureRelation
betweeninstantaneous/tdyn= (⇡ /40)how↵vir(magvs. "↵,obs in simulations (squares and diamonds; ↵- and3.9.2. Comparison
with Observations-seriesmodels)turbulenceand observations(circles). For
simulations,
we use the instantaneous
mass, size, and velocity dispersion of the
netized)regulates
SFR(e.g., Krumholz&
gasMcKeecloud and2005;
cluster’sbolometricluminosity.2011;
We showsnapshots
the compute
time interval< t < t⇤,90% . The“observed”
circles in
Next,in &wethet⇤,0time-dependent,
Padoan& UV
NordlundHennebelledi↵erent panels show data from di↵erent observational studies: SFE(a) low-mass,
nearbycloudscompiled
Evansobservations
et al. (2014); of
per freefall
time"↵,obsusingbymock
Chabrier
2011;
see
also
the
reviews
byFederrath&
1(b) molecular clouds in the Galactic Plane with massive star formation
rateclouds.
M⇤ > 5 M(Vutisalchavakul
et al.instanta2016).
simulated
For Myrthis exercise,
we use theKlessen2012
andPadoan
et
al.
2014)
In
brief,
these(c) GMCs with associated star-forming complexes in the Milky Way(Lee
et al.mass
2016).
with associatedstar-forming
neous
cloud
(M(d)
, assumed
to be recoverable
fromneuGMCsmodelsin theassumethat starCloudformation
place
in dataalinescomplexes
Large Magellanic(Ochsendorftakes
et al. 2017).
All
pointsare coloredby one-dimensional
velocitymolecular
or dustemission),size (R = 21/3
R50% ),
dispersion
sizes denote gas turbulent
mass. The solidand dashedgray
and the yellow
shadedthe
areas,are
1d and their(magnetized)
supersonically
cloud
where
thelines,and velocity
dispersion( 1d )and
to green
computedynambasedon the same functional
Figure 14,
but for instantaneous
panel (a), the grey bars indicate
ical
time measurements.
tdyn = R/ In
distributionof pmassforms
withas indensityfollows
a lognormal
1d and the free-fall time t↵ =⇥⇤p
the ranges of "↵,0 and t↵,0 /tdyn,0 plottedbars indicate the ranges of median values"↵,obs and23⇡/(32G¯⇢), where ⇢¯ = 3Mneu /(4⇡R3 ). ofTheinstanPDF: p(s) = ( 2⇡ s ) 1 exp in Figure(s 14,
s0 )while/(2thes2 )black, wheret↵ /tdyn for snapshots between t⇤,0 and t⇤,90% .
1/2
taneous
t
/t
/
(R/M
)
.
↵
dyn
1d
neu
s = ln(⇢/⇢0 ), s0 = s2 /2, ⇢0 is the mean density,and ratios
are derived from extinction maps, velocity dispersion(even
for cases
the MIR-based
may be unAssumingthatshown,an observer
is ableSFR
to recoverclusincreases
athigherMach
13number M, with some reducfrom molecular linewidth (mostly CO), and SFRs fromderestimatedby UV
a factor
2–3). Panel(c) massive
shows thestar
Milky
ter’s
bolometricluminosity
fromfortionforstrongermagnetic
fields.
counting young stellar objects (YSOs). Panel (b) shows
Way GMCs
associatedstar-forming
complexes
mationtracerswithsuch
as H↵ and/or
infrared
fluxes, inthe
The collapse
is assumed
in regions
the sample
of larger molecular
cloudstoin occur
the Galactic
Leewhere
et al. (2016) (their Fig. 7). Their samples
are1
instantaneous SFR
is calculated as Ṁ⇤ =LUV /tUV ,12
Plane
ongoing
starplus
formation,compiled
built
from the CO catalog of molecular clouds (Mivillethewith
supportbymassive
thermalturbulentpressurecannotwhere
LUV = LLyC + LFUV is the bolometric UV

-

-

-

Figure 14.
Time-averaged SFE per free fall time defined as h"↵ i"↵,0 = t↵,0 hṀ⇤ i/M0 = t↵,0 /tdep,0 as a function of the ratio between the initial freefall and dynamicaltimes in the ↵-series (red squares) and the -series (blue dip
amonds) models. The empirical law exp( 1.6 3t↵,0 /tdyn,0 )found by Padoan et al. (2012) is shown as a gray dashedline. The gray solid shows the fit to the ↵-series simulations
exp( 4.06t↵,0 /tdyn,0 ). The yellow (green) shaded region is
the prediction by the analytic model of Krumholz & McKee(2005) (Hennebelle & Chabrier (2011)) for the Mach number
range 13.6 < M < 30.4 with b = 0.5 and = 0.02 (see text
for more details).

t = 0 and the time-averaged SFR hṀ⇤ i ashṀ⇤ it↵,0"↵,0 = t↵,0=.
M0tdep,0(28)This is similar to the approach adopted by Padoan et al.
(2012); Raskutti et al. (2016). Since t↵,0 is the same forall of the present models, tdep,0 and "↵,0 are inversely
proportional to each other.

tUV = 6.7 Myr, to be consistent with the relationṀ⇤ = 1.5 ⇥ 10 4 M Myr 1 (LU V /L ) widely adopted
in observations (Kennicutt & Evans 2012; Liu et al.
2020). This assumes a more exact age for the stellar system is not known, although with spectral fitting cluster ages can be obtained (e.g. Adamo et al.
2017). The observed SFE per freefall time is then"↵,obs = (t↵ /tUV )LUV /( Mneu ) / M⇤ (R/Mneu )3/2 .12
Figure 15 plots as squares (↵-series) and diamonds( -series) "↵,obs for all models as a function oft↵ /tdyn . The top axis also shows instantaneous ↵vir =(40/⇡ 2 )(t↵ /tdyn )2 . We consider snapshots in the time
range t⇤, < t < t⇤,90% . The color of symbols indicatesthe 1D velocity dispersion. In contrast with "↵,0 in Figure 14, all of our models start from the lower left with↵vir ⇠ 0.5–1.5 and "↵,obs ⇠ 10 3 , and, as time progresses, they move upward slightly and then to the upper
right. This trajectory reflects the evolution of cloud’s
gas mass, size, and velocity dispersion, and the cluster’s mass and luminosity. The stellar mass M⇤ and UV
luminosity start very small and keep increasing during
the main star formation phase (note that tSF is comparable or shorter than tUV , except for a small number of late snapshots in strongly magnetized clouds).
Meanwhile, radiation feedback drives mass loss and radial expansion, leading to an increase in R/Mneu , together with an increase in velocity dispersion 1d . As
a result, evolution causes "↵,obs / M⇤ (R/Mneu )3/2 andt↵ /tdyn / 1d (R/Mneu )1/2 to increase in tandem. The
distribution of t↵ /tdyn extends to larger values (by a factor of a few) than in Figure 14 due to cloud evolution.
Even more extreme is the stretch in the distribution of"↵,obs by two orders of magnitude compared to "↵,0 in
Figure 14. We note, however, that in their later stages,
the simulated systems have very low gas mass (leading to"↵,obs & 0.1 and t↵ /tdyn & 1), and would not necessarily
be included in observed star-forming GMC samples.
Although there is a large variation during the evolution, we calculate the “characteristic” values of t↵ /tdyn
and "↵,obs as a median for snapshots equally spacedin time between t⇤,0 and t⇤,90% . We find 0.54 <med(t↵ /tdyn ) < 1.08 and 0.8% < med("↵,obs ) < 3.7% inthe ↵-series models, while 0.54 < med(t↵ /tdyn ) < 1.23
and 0.8% < med("↵,obs ) < 3.5% in the -series models.
The bars in Figure 15(a) indicate the ranges of the medians. These values are similar to the ranges of t↵,0 /tdyn,0
and "↵,0 shown in Figure 14, as indicated by grey bars
in Figure 15(a).
The circles in each panel show estimates of the ratio--

by Vutisalchavakul et al. (2016) (their Fig. 6). In their

Deschênes et al. 2017), cross-correlated with WMAP

21. Characterization of very wide companion candidates to young stars with planets and disks,F. Z. Majidi, et al.
Discovering wide companions of stellar systems allows us to constrain the dynamical
environment and age of the latter. We studied four probable wide companions of four different
stellar systems. The candidates were selected mainly based on their similar kine- matic
properties to the central star using Gaia DR2. The central stars are V4046 Sgr, HIP 74865,
HIP 65426, and HIP 73145, and their probable wide companions are 2MASSJ18152222-3249329, 2MASS J15174874-3028484, 2MASS J13242119-5129503, and2MASS J14571503-3543505 respectively. V4046 Sgr is a member of β-Pictoris Moving
Group while the rest of the stellar systems are acknowledged as members of the ScorpiusCentaurus association. The selected stellar systems are particularly interesting because all ofthem are already known to possess a low-mass companion and/or a spatially resolved disk.
Identifying wider companions of these stars can improve their eligibility as benchmarks forunderstanding the formation channels of various triple systems, and can help us to determine
the orbits of their possibly undiscovered inner, wider companions in case of higher
multiplicity. By analyzing the X-Shooter spectra of the wide companion candidates of thesestars, we obtained their stellar parameters and determined their ages. We find that 2MASSJ15174874-3028484 (0.11 M⊙, 7.4±0.5 Myr), an already recognized pre-main sequence(PMS) member of Scorpius- Centaurus association, is a highly probable wide companion of
HIP 74865. 2MASS J13242119-5129503 (0.04 M⊙, 16±2.2 Myr) is ruled out as a plausible
wide companion of HIP 65426, but confirmed to be a new sub-stellar member of the ScorpiusCentaurus as- sociation. 2MASS J14571503-3543505 (0.02 M⊙, 17.75±4.15 Myr) is a
probable sub-stellar member of the same association, but we cannot confirm whether or not it
is gravitationally bound to HIP 73145. 2MASS J18152222-3249329 (0.3 M⊙, older than 150
Myr) is determined to be a mildly active main sequence (MS) star, much older than membersof β-Pictoris Moving Group, and unbound to V4046 Sgr despite their similar kinematicfeatures. PMS wide companions such as 2MASS J15174874-3028484 might have formed
through cascade fragmentation of their natal molecular core, hinting at high multiplicity inshorter separations which can be confirmed with future observations.

Gaia DR2 のデータを使って、wide
companion の候補を4つ持ってきた(V4046 Sgr, HIP 74865, HIP
65426, and HIP 73145,)。
これらの天体はすでに伴星か円盤があることがわかっている。この研究は、そのさらに外側の星を探す。
X-Shooter を使って年齢を推定。4つについてそれぞれ議論した。

Majidi et al.: X-Shooter characterization of four new objects

Table 1: Physical parameters of the central stars.
Name
V4046SgrHIP 74865
HIP 65426
HIP 73145
GQ LupDistance a(pc)

Association

SpT

T e↵
(K)

Av
(mag)

Wide companion

Separation(00 )Separation
(au)

72.4
123.53109.21133.65151.82-Pictoris MG
Sco-Cen (UCL)
Sco-Cen (LCC)
Sco-Cen (UCL)
Sco-Cen (Lupus I)

K5/K7
F3V
A2V
A2IV
K7V

4370/4100
6720
8840
8840
4070

0b
0
0
0
0.7 c

2MASS J1815-3249
2MASS J1517-3028
2MASS J1324-5129
2MASS J1457-3543
2MASS J1549-3539

901
90
142
280
16

65232
11118
15508
37422
2429

a

Distances are calculated based on the parallaxes of the objects reported in Gaia DR2 catalog.
b(Stempels & Gahm 2004)c(Alcalá et al. 2017)
A&A proofs: manuscript no. 39031corr
The X-ray luminosity as derived from the observed X-ray
from XMM (source 3XMMflux
J181522.2-324932
from
Rosen
et
evaluationassuitable
indicators
arespreadindi↵erent
spectral(Rosen et al. 2016) coupled with our adopted stellar paramal. 2016), suggesting it is anetersactivestar.log(LInformationon the sep- regions. For details of the instrument description, the reader is
イントロ
yieldsX /Lbol ) = 4.14. This value is consistent witharation of the object from thecentralstar andmagnitudes
referredVernetthoseof Hyadesstars itsof similarcolors,isand wellbelowtothoseof et al. (2011). X-Shooter also enables us to deincluded in Tables 1 and 6, respectively.
the radial
members of younger clusters
or groups such termine
as the Pleiades
and velocity of objects within a precision range of・Gaiaの感度だと、褐色矮星の伴星<20Myr,-Pictoriswide
MG (Song
et al. 2008),
indicating
most (Frasca
probable et al. 2017) and is therefore useful for confirm2MASS J1517-3028 is thea possible
companionof HIP
2-5akm/sof severalassociation,
hundred million
74865. It is a member of theage
Sco-Cenhas years.
a brown ing or ruling out the physical association of the candidates with
<150pcくらいのものは見つけられる。
00
The age (0.13
inconsistency
with theMG, parent
which isstars.
sup- Four of our targets were observed through thedwarf companion at small separation= 16 au),and -Pictoris
was theirported by the position of 2MASS J1815-3249 on
the 00CMD, a low
00
00
discovered with Sparse Aperture
Masking
observations(Hinkley
1
.
0,
0
.
9,
and
0
. 9 slits for one or two cycles based on their G and・SPHERE-GTOターゲットの中で、Gaia
DR2で伴level of magnetic activity, and the absence of lithium absorption,
et al. 2015) and was recentlytogether
characterized
SPHERE(Canmagnitudes.
Among these four, for the three fainter obwith the withdiscrepant
RV withrespectJ tobandthat ofthe centalloube et al. in prep.). The tral
expectedmass
thetarget
new candidate
jectsmember
we performed
星っぽいものを探す→数十ターゲット。
star imply
thatofthis
is not a legitimateof the two cycles of ABBA nodding mode. Theseis about 0.17 M if coeval togroupthe central
star.
three targets are 2MASS J1324-5129, 2MASS J1457-3543, andand is significantly
older than expected.
2MASS J1324-5129 is a possible, very wide companion of 2MASS J1549-3539 with G magnitudes 18.71, 18.52, and 18.37
・分光観測して、若い天体に絞る。→表の通り
HIP 65426, and hosts a planetary-massline fluxes
with aJ1517-3028
precision(blue)
bettertogether with
4.2. 2MASS companion
J1517-3028 recently dis- respectively. As measuring
Fig. 2: Spectrumof 2MASS
covered with SPHERE at about 90 au (Chauvin et al. 2017). The than 20% –at least thein thevisible(VIS) andnear-infrared(NIR)
available
photometricdata(red dots) in theliterature and the
2MASS
J1517-3028
hasalreadybeen
identified
asaPMS
memsystem is part of the Sco-Cen association and its expected mass arms– was a necessity
achievingaims, a total on-source
fittedfor
BT-Settl
modelour(black).
ber of UCL (Damiani et al. 2019). The assessed age for the UCL
is about 25 M jup if the targetsub-associationis coeval to HIP
executionof 1.2 hr was requested for the aforementioned
was 65426.
reported to be 15±3 Myr
by Pecaut time& Ma2MASS J1457-3543 is a possible comoving object to HIP objects. For 2MASS J1517-3028, which is brighter than the

Majidi et al.: X-Shooter characterization of four new objects

Table 3: Physical stellar parameters of candidates. Except for 2MASS J1815-3249, mass, age, and log g are selected from the closest
matches o↵ered by Bara↵e et al. (2015) models. The same quantities are extracted from MIST isochrones for 2MASS J1815-3249.
Name
2MASS J1815-3249
2MASS J1517-3028
2MASS J1324-5129
2MASS J1457-3543
2MASS J1549-3539
a

SpT

T e↵
(K)

Av
(mag)

RV
(km/s)

L?
(L )

R?
(R )

M?
(M )

Age
(Myr)

log g

M3±0.5M4.5±0.5
M6.5±0.5
M8±0.5
M4±0.5

3410
3100
2710
2500
3190

0
0
0
0
1.0±0.517.7±2.4
1.2±6.915.3±5.8
4.0±8.9
1.3±4.0

0.03
0.018
0.00240.002
0.07 a

0.5
0.47
0.22
0.24
0.21

0.3
0.11
0.04
0.02
0.15

> 150
7.9
18.2
14.09
2

4.95
4.13
4.35
4.00
3.78

This object is sub-luminous. The value reported here is the object’s corrected luminosity obtained by Alcalá et al. (2020).

Table 4: ROTFIT results acquired for the targets. The estimated ages are concluded from Bara↵e et al. (2015) models.
Name
2MASS J1815-3249
2MASS J1517-3028
2MASS J1324-5129
2MASS J1457-3543
2MASS J1549-3539

T e↵
(K)

log g

RV
(km/s)

v sin i(km/s)

Age(Myr)

3562±303077±222646±502635±803230±1014.68±0.14
4.49±0.21
4.0±0.1
4.30±0.17
3.74±0.23

20.1±2.0
1.4±2.417.9±3.0
8.0±7.8
2.0±2.8

< 8.0
26±6
< 8.0
62.0±10.0
13.0±6.0> 150
7.4±0.5
16±2.217.75±4.15
2.75±0.75

Table 5: Kinematic properties of our targets and their associated stellar systems. For a discussion on the adopted kinematic properties
of 2MASS J1549-3539, the reader is referred to Alcalá et al. (2020); we have adopted the parallax of the central star for 2MASSJ1549-3539 because of its highly uncertain measured kinematic features as the only accreting object among the targets.
Nameparallax

µ↵

µ

RV

Majidi et al.: X-Shooter characterization of four new objects

Table 8: The wide-companion membership status checklist for our targets.
Name

Consistent kinematic properties
with the stellar system(yes/no)Age

ActiveContains Li i

(MS/PMS)(yes/no)(yes/no)2MASS J1815-3249
2MASS J1517-3028

no
yes

MS
PMS

no
yes

no
yes

2MASS J1324-5129
2MASS J1457-3543
2MASS J1549-3539

no
no
yes

PMS
PMS
PMS

yes
yes
yes

yes
ambiguous
yes

conclusion

field object
UCL member +HIP 74865 probable wide companion
LCC new member
ambiguous
Lupus I new member +GQ Lup probable wide companion

結論
V4046 Sgr - 候補生 2MASS J1815-3249は関係ない単独星だった。残念。
HIP 74865 and GQ Lupi:(a)Fig. 9: Log–log plot of total velocity di↵erence v (km/s) vs.
重たい星のすぐ近くを褐色矮星が周り、遠くに遠隔伴星がprojected separation s(au) for the five plausible wide companions studied in this work. v(km/s) (orange line) indicatesthe maximum total velocity di↵erence that bound binaries withいるシステム。
a total mass equal to 10 M in circular orbits can possess; see
Ramírez et al. (2019) for details.
HIP 65426 - 2MASS J1324-5129 とは関係なさそう。
max. Therefore, it might be true that the acquired age for 2MASSHIP 73145: 2MASS J14573543.
J1517-3028
based onとは関係なさそう
Bara↵e et al. (2015) models (that neglect

the e↵ect of magnetic fields) is younger than its real value by
a factor of two. Assuming the whole stellar system as a triple
system, HIP 74865 hosts a brown dwarf companion at a small
separation with an M-dwarf star at a larger separation. As an al-

面白そうな天体の簡易まとめ
HIP 74865について：F3V型, ScoCen association, 16auのところに褐
色矮星がいる、SPHEREで取られている、質量0.17Msun、2MASSJ1517-3028 のseparationは11118
au
GQ Lupiについて：K7V型、Sco-Cen(Lupus I) , 若い2Myr、GQLup Bが20-40 M jup、円盤ある、2MASSJ1549-3539 のseparationは2429au

A new prolonged decrease event in the brightness of theyoung stellar object V2492 Cygni
Sunay Ibryamov, Evgeni Semkov
Results from the BV(RI)c photometric observations of theyoung stellar object V2492 Cyg collected in the period from
April 2018 to September 2020 are presented. Theseobservations are a part of our monitoring of the star thatbegan in 2010 and continuing to date. V2492 Cyg is locatedin the Pelican Nebula, and its variability was explained byboth accretion and extinction events. The new photometricdata show that the star continues to exhibit rapid irregularvariability in all bands. In the period from March 2019 to
May 2020, we registered a prolonged decrease event in thelight curve of V2492 Cyg.

- V2492 Cygという星の2018-2020年のモニタリン
グ報告。

- プロジェクト自体は2010年から。
- 継続的に変動を見せ続けている。
- March 2019 からMay 2020にかけて、の減光イ
ベントを同定19. The CARMA-NRO Orion Survey: Filament Formation
via Collision-Induced Magnetic Reconnection – The Stickin Orion A, Shuo Kong, et al.
A unique filament is identified in the Herschel maps of theOrion A giant molecular cloud. The filament, which, wename the Stick, is ruler-straight and at an early evolutionary stage. Transverse position-velocity diagrams show
two velocity components closing in on the Stick. The
filament shows consecutive rings/forks in C18O(1-0)
channel maps, which is reminiscent of structures generated
by magnetic reconnection. We pro- pose that the Stickformed via collision-induced magnetic reconnection(CMR). We use the magnetohydrodynamics (MHD) code
Athena++ to simulate the collision between two diffuse
molecular clumps, each carrying an anti-parallel magnetic
field. The clump collision produces a narrow, straight, dense
filament with a factor of >200 increase in density. The
production of the dense gas is seven times faster than freefall collapse. The dense filament shows ring/fork-likestructures in radiative transfer maps. Cores in the filament
are confined by surface magnetic pressure. CMR can be an
important dense-gas-producing mechanism in the Galaxy
and beyond.

ハーシェルで見つかったOrionA分子
雲のフィラメントについて。Stickと名付けた。
PV図解析でみると、2つのコンポー
ネントがみえる。
Stickは、衝突励起の磁気リコネクシ
ョンで形成されたと提案する
Athna++で2つの分子クランプの衝突を計算、200倍くらい密度の高いフィラメントを再現。自由落下よりも7倍早くフィラメントを作れる
ring/fork-like構造も作れる。

The Cloud Factory II: Gravoturbulent Kinematics of Resolved Molecular Clouds in a
Galactic Potential
Andres F. Izquierdo, Rowan J. Smith, Simon C. O. Glover, Ralf S. Klessen, Robin G.
Tress, Mattia C. Sormani, Paul C. Clark, Ana Duarte-Cabral, Catherine Zucker
We present a statistical analysis of the gravoturbulent velocity fluctuations in molecular
cloud complexes extracted from our "Cloud Factory" galactic-scale ISM simulation
suite. For this purpose, we produce non-LTE 12CO J=1-0 synthetic observations andapply the Principal Component Analysis (PCA) reduction technique on a representative
sample of cloud complexes. The velocity fluctuations are self-consistently generated bydifferent physical mechanisms at play in our simulations, which include galactic-scale
forces, gas self-gravity, and supernova feedback. The statistical analysis suggests that,even though purely gravitational effects are necessary to reproduce standard
observational laws, they are not sufficient in most cases. We show that the extra injectionof energy from supernova explosions plays a key role in establishing the global turbulent
field and the local dynamics and morphology of molecular clouds. Additionally, wecharacterise structure function scaling parameters as a result of cloud environmental
conditions: some of the complexes are immersed in diffuse (inter-arm) or dense (spiralarm) environments, and others are influenced by embedded or external supernovae. In
quiescent regions, we obtain time-evolving trajectories of scaling parameters driven by
gravitational collapse and supersonic turbulent flows. Our findings suggests that a PCAbased statistical study is a robust method to diagnose the physical mechanisms that drive
the gravoturbulent properties of molecular clouds. Also, we present a new open source
module, the PCAFACTORY, which smartly performs PCA to extract velocity structure
functions from simulated or real data of the ISM in a user-friendly way. Software DOI:
10.5281/zenodo.3822718- Cloud Factory というシミュレーシ
ョン結果の解析の一つ- 分子雲におけるgravoturbulent
velocity ﬂuctuationの解析を主たる目的とした。
- non̲LTE 12COの疑似観測をやってPCA解析。
- 結果：乱流状態と分子雲のダイナミ
クスと形を再現するには、超新星爆発のようなエネルギー注入が重要であることがわかった。

Evolution of the density PDF in star forming clouds: the role of gravity
Etienne Jaupart, Gilles ChabrierWe derive an analytical theory of the PDF of density fluctuations insupersonic turbulence in the presence of gravity in star-forming clouds.
The theory is based on a rigorous derivation of a combination of theNavier- Stokes continuity equations for the fluid motions and the Poisson
equation for the gravity. It extends upon previous approaches first by
including gravity, second by considering the PDF as a dynamical system,
not a stationary one. We derive the transport equations of the density PDF,
characterize its evolution and determine the density threshold above
which gravity strongly affects and eventually dominates the dynamics ofturbulence. We demonstrate the occurence of two power law tails in the
PDF, with two characteristic exponents, corresponding to two different
stages in the balance between turbulence and gravity. Another importantresult of this study is to provide a procedure to relate the observed column
density PDFs to the corresponding volume density PDFs. This allows toinfer, from the observation of column-densities, various physicalparameters characterizing molecular clouds, notably the virial parameter.
Furthermore, the theory offers the possibility to date the clouds in units oftcoll, the time since a statistically significant fraction of the cloud started
to collapse. The theoretical results and diagnostics reproduce very wellnumerical simulations and observations of star-forming clouds. The theory
provides a sound theoretical foundation and quantitative diagnostics toanalyze observations or numerical simulations of star-forming regions andto characterize the evolution of the density PDF in various regions ofmolecular clouds.

- 分子雲における、超音速乱流下での密度関数を解析的に求めた。

- ナビエ・ストークスの連続の式と、重力についてはポアソン方程式を使った。

- PDF の移流方程式をだした。
- 結果、乱流と重力の釣り合いによ
り、二個のパワーローが出てくる。

- 観測される面密度PDFへのから空間密度PDFへの関係性も示した。