THE STAR FORMATIONNEWSLETTER
No. 320 (#36), No. 321 (#1-9)
塚越崇(NAOJ)
4 Oct. 2019

1. Northern Galactic Molecular Cloud Clumps in Hi-GAL: Clump andStar Formation within Clouds
2. HST spectra reveal accretion in MY Lupi
3. Gaia DR2 distances to Collinder 419 and NGC 2264 and new
astrometric orbits for HD 193 322 Aa,Ab and 15 Mon Aa,Ab
4. Inner rocky super-Earth formation: distinguishing the formation
pathways in viscously heated and passive discs
5. The First Detection of 13C17O in a Protoplanetary Disk: a Robust
Tracer of Disk Gas Mass
6. The fulcrum wavelength of young stellar objects – the case of LRLL 31
7. Radial Drift and Concurrent Ablation of Boulder-Sized Objects
8. The Structure of Dark Molecular Gas in the Galaxy – II. Physical Stateof “CO-Dark” Gas in the Perseus Arm
9. Stars and brown dwarfs in the sigma Orionis cluster. IV. IDS/INT andOSIRIS/GTC spectroscopy and Gaia DR2 astrometry
10.Dust accretion in binary systems: implications for planets andtransition discs

Northern Galactic Molecular Cloud Clumps in Hi-GAL: Clump and Star Formation within Clouds
Zetterlund, E., et al.

• Hi-GALデータとCOHRSデータを使った銀河面分子雲の調査• GMCとクランプ、SFRの関係性をglobalに調べる• GMCとclumpのmass fractionを用いて評価• 主にたくさんの結果を示している論文

• Mcloud - SFRの関係をdense substructureも念頭に入れて調べる• 2つの既存カタログをマッチングして使用• Recent star formation - 連続波(特に短波長)• HiGALのクランプカタログ [Ellsworth-Bowers+2013]+ distance determination tech. [Zetterlund+2017]• クランプの質量などを導出• Initial condition of star formation - CO
• CO High Resolution Survey (COHRS)のクラウドカタログ [Colombo+2019]• クラウドの質量やサイズ、ビリアル比等を導出• Clump formation fraction(efficiency)• 4704のHi-GAL clumpが1182のCOHRS cloudとマッチex)

• Clump mass fraction• Clump>3’のもので評価• 平均で~0.2

5.3. Relations with Physical Properties• GMC質量(面密度)と弱い相関• 1, 重いcloudほど多くのclumpを持つ• 2, 重いcloudほど重いclumpを持つ• Rgalにも非常に弱い相関がある=>環境効果

To determine what properties of a GMC inﬂuence its CFE
we plot clump mass fractions of clouds against various physica
properties (Figure 10). In examining cloud mass, surface massdensity, virial parameter, and line width we see no more thanslight correlations between clump mass fraction and thesecloud properties. In the case of cloud mass and cloud mass
surface density, we ﬁnd shallow power-law slopes of −0.14±0.03 and −0.13±0.05, respectively, with large dispersions. Acloud with more mass will, on average, have a slightly lower
clump mass fraction. A cloud’s stability against collapse as awhole (see Section 8) does not inﬂuence the fraction of acloud’s mass that will collapse into denser clumps.
In these plots, red points correspond to non-virialized clouds(αGMC<2), that is, those unstable to gravitational collapseThis comprises ∼25% of our sample. Panel (c) shows that thenon-virialized clouds sample from the full distribution of clump
mass fractions, as do virialized clouds. As seen in panels (a)and (b), these potentially collapsing clouds are more likely tohave higher masses and surface mass densities. However, weknow that lower-mass clouds (M∼103 Me) can collapse toform stars (Enoch et al. 2008). These clouds are not absenfrom the COHRS catalog, but we are less likely to match HiGAL clumps to them. This may be due to the small clouds’L (70)(erg s-1)
1SFR (M: yr ) =.
1.7 ´ 10 43(12)r about this relation is ∼0.2 dex, or about 60%.
this calibration was developed using individual
th typical scales of hundreds of parsecs and it will be
le for individual GMCs. At this wavelength, theraces light from stars of ages 0–100 Myr, but thession-contributing stellar age is 5 Myr.
ur clouds are deﬁned in PPV space, we are using atracer to measure star formation. We control for thiste our uncertainties due to this mismatch as follows.
the ﬂux densities required for conversion to SFRs,
ate the ﬂux density found within the velocityCOHRS mask. We account for background ﬂux
emission by creating a “rind” around the mask with a–1.5
resolution
elements,
processofclump
mass fractions
or mimicking
the SFR in theclumps
is smaller
in clouds
hotometry but for irregular ISM structures. We ﬁt awith greater clump mass fractions. Figure 15 (right)
GMCshows that
it is the former: clump masses tend to be larger in clouds withgreater clump mass fractions. Figure 15 also•showsthat clumpsLada+2010と異なる結果
with lower virial parameters tend to have higher SFRs (left) andthat more massive clumps tend to have smaller virial
Clumpparameters (right). While the scatter is verylarge, taken
together these weak correlations are consistent with anevolutionary trend in which clumps and clump mass fractions
grow within clouds over time, become increasingly gravitationally unstable, and have correspondingly increased SFRs.
Figure 16 shows that more massive clouds host more starformation. This behavior has been seen before in other studies,including the Gould Belt clouds in the solar neighborhood(Lada et al. 2010) and H II regions in the nearby galaxies NGC300 (Faesi et al. 2014) and NGC 628 (Kreckel et al. 2018),
among other work. We place these other data on Figure 16,alongside our own. The work of Lada et al. (2010) argued thatthe local molecular clouds followed a linear correlation• M

はSFRとnon-linear• M

はConstant SFE• K-S lawと比較• ややsteep

The First Detection of 13C17O in a Protoplanetary Disk: a Robust Tracer of Disk Gas Mass
Booth, A.S., et al.

• 13C17Oを原始惑星系円盤で初検出(HD163296)• 光学的に薄いのでガス質量を出せる• CO同位体を合わせるようにガス質量を推定=> 今まで考えられていた質量より2-6倍重い• ALMA cycle3 data (Notsu et al. 2019)• HD163296 (d=101.5pc,ただし論文では122pcを使用)• resolution: 0.87”x0.51”• S/N=7で13C17O(J=3-2)検出• そのほかのCO同位体データ• CO, 13CO, C18O (J=2-1)• CO, C17O (J=3-2)• optically thinを仮定して柱密度導出• 7.1x1015 cm-2 (<50au, 50K assumption)• C18Oから求めた柱密度との比~2.5=> C18Oはoptically thick• もしoptically thinなら比 ~100

• Qi+2011のモデルをベースに放射輸送計算+擬似観測• 13C17O以外のCOを再現する円盤モデル• CO depletion入り• 試したモデル• Model 1: モデルそのままで一定のX(13C17O) = 5.39x10-10• Model 2: ガス質量全体を増やす(目で合わせて3.5倍とした)• Model2の方が合う• Model1だと13C17Oを過小評価• total disk mass: 0.31M⦿
• g2d~260• これまでの見積もりより2-6倍重い• 光学的に厚いCOで求めていたせい• ISMの同位体比でフィットできている• selective photo-dissociationは起きていない？
• C18Oでも光学的に厚いので、rare COの高感度観測が必要• CO snow line?

• これまでのモデル• 1, 155au at 19K (Qi+2011); 2, 90au at 25K (Qi+2015)• 今回の観測は前者でフィットできるが、もっと高感度の観測が必要• 重力不安定?
thin, CO isotopologues, or by detecting molecules that,
due to chemistry, peak in abundance at a location re• 最小となるQでも~6(110au)くらい
lated to the snowline. Qi et al. (2011) use observations
CO and put the snowline at 155 au at 19 K: follow up=>”今は”安定•work suggested that the snowline was instead at 90 au
and 25 K (Qi et al. 2015). Out of these two scenarios
質量降着率から初期の円盤質量を概算すると
our 13C 17O observations fit best with the former option
given the data in hand. The relationship between the0.1Myrで1.33M
⦿cations, DCO+ (ring from 110 and 160 au) and N2 H+(inner edge of the ring at 90 au), and the location of the=>円盤初期では不安定
midplane CO snowline is not trivial, but both species
have been detected in this disk (Mathews et al. 2013;
Qi et al. 2015). Our analysis shows that the C 18O emission in both models tested is optically thick, and thus

Dust accretion in binary systems: implications for planets and transition discs
Chachan, Y. et al.

• 連星系惑星の観測• 連星系のdynamical stability limitのすぐ外側で見つかる傾向• そこで作られた？or 別の場所で作られて移動した？
• 前者はbinaryの影響が大きく難しい。
• 後者の方が優勢。truncated edgeでのdust trappingで作る• ただ現行の計算では観測のlocationはまだ再現出来ていないらしい
Dust accretion in binary systems

3• q~0.1くらいのCircumbinary diskでdust trappingが起き、惑星がでHere, dust and gas particles have masses md and mg ,ties ⇢d and ⇢g , and local velocities vd and vg respecy. P is gas pressure, (~r) is the gravitational potential
e binary (see appendix B) and avisc is acceleration due
tificial viscosity (see Springel 2005). In our simulations
ompute the drag coefficient assuming
constant Stokes
pber, S t = t s ⌦(r), where ⌦(r) = GM/r3 , the Keplerian
lar frequency due to the total mass of the binary, M.
• GADGET-2
The binary orbits in an anticlockwise direction and has
ss ratio q •= M
2 /M1 , total mass M = M1 + M2 , semi-major
ガス+ダストa, and eccentricity e. The simulations are run in the inframe centredon theregime
binary’s centreof mass. The stars(q<0.1)• stellar
reated as sinks and particles that come within 0.01a ofr star are• accreted
onto regime
it (Bate et(q<0.01)
al. 1995). The biplanetaryparameters are assumed to be constant over the duraof the simulationand the stellar gravitational
softeningdiskの観点• pre-transitional
h is set to 0.001a, chosen to be small enough to ensure
the circumstellar discs are not a↵ected by the softenWe do not include the e↵ects of disc self-gravity or the-reaction of the disc on the binary in our simulations. Aきるかどうか？を検証する• SPHによる計算

q=0.01のケース

The circumstellar disc masses are est
the mass within the Hill radius that

• ある程度の時間でCBdiskにダスト(St=0.1)がたまり、セカンダリへの降着はpure gasになる• 時間はeに依存include the results obtained
tablished (as in Young et alof dust and gas to the prim
the fractional accretion rat• ~15%のCB diskのガスがprimaryに入る• eに依存せず=• λdustは減少を続ける(e=0.2)

ṀprimaryṀprimary + Ṁsecondary.

We also use the ratio ṀD / Ṁof dust and gas accretion rratio.

ṀD Ṁprimary,dust + Ṁseconda=ṀGṀprimary,gas + Ṁseconda

Y. Chachan et al.

λwhere each Ṁ is normalized
ticles.
Y&& (2015)
0.5
Figure 3. The
gas and dust (for e = 0.2) accretedStfraction= 0. 1ofDust

onto the primary, , plotted as a function of time. After an initialtransient phaseGDs(⇠ 200tdyn , greyedfor gasACCRETION
settles to a value3 out),
DUSTof ⇠ 0.15. for dust can only be properly quantified for e = 0.2 as
0.4
dust accretion rate is too low for for lower eccentricites. There is
3.1 Planetary regime
a clear decline in for dust, implying that the amount of dust
accreted onto the binary (Figure 2) as well as onto the primary
0.3
We begin by considering ais going down with time.

0.20.10.00.00.2(i.e. a 10 MJ planet for a sknown that companions ondust in pressure maxima t
as is suspected to be thePinilla et al. 2012; Zhu et aFigure 10.4
0.60Dss 5DtiR (q)0.8• 近星点距離がバイナリ間隔の2.2-2.8倍になるとダスト降着が止まる• qが大きくなると降着を止めるには大きな近星点距離が必要• 実際のCB cavityは大きいのでCBで惑星を作る説が優勢かも• とはいえ観測されているplanet locationにはmigrationが必要• planetary mass regimeの結果について• eが大きければ降着が止まるまで時間がかかる(>300orbits)が、内側で惑星を作る時間(~Myr)より短い• そういう系にあるhot Jupiterはin situ formationではなさそう• eの大きいplanetary mass companionの存在は、pre-transitional diskの解釈になるかも• 複数の重い惑星がなくても良い

HST spectra reveal accretion in MY Lupi
Alcala, J.M., et al.

Gaia DR2 distances to Collinder 419 and NGC 2264 and new astrometric orbits for HD 193322
Aa,Ab and 15 Mon Aa,Ab
Apellaniz, J.M.

Inner rocky super-Earth formation: distinguishing the formation pathways in viscously heated andpassive discs
Bitsch, B.

The fulcrum wavelength of young stellar objects – the case of LRLL 31
Bryan, G.R., et al.

Radial Drift and Concurrent Ablation of Boulder-Sized Objects
Burn, R., et al.

The Structure of Dark Molecular Gas in the Galaxy – II. Physical State of “CO-Dark” Gas in thePerseus Arm
MBusch, M.P., et al.

Stars and brown dwarfs in the sigma Orionis cluster. IV. IDS/INT and OSIRIS/GTC spectroscopy andGaia DR2 astrometry
Caballero, J.A., et al.