星形成ゼミ
SFN #318 (28-37本目)
百瀬宗武（茨城大・理）詳細解説は以下の3本• #30: Streaming Instability for Particle-Size Distributions (L. Krapp et al. ApJL, 878, L30)
• #35: The anomalously low (sub)millimeter spectral indices of some protoplanetary disks may
be explained by dust self-scattering (H. “Baobab” Liu, ApJL, 877, L22)
• #36: Large-scale periodic velocity oscillation in the filamentary cloud G350.54+0.69 (M. HongLi Liu et al. MNRAS, 487, 1259)
その他はひと言メモを参照

30. Streaming Instability for Particle-Size Distributions
L. Krapp et al. (ApJL, 878, L30)
https://arxiv.org/abs/1905.13139【内容】Linear Streaming Instability の最大成長率Ｓ
Ｓ

颸辶㫳サイズの分布

巸て考慮して解析．幅広いダストガス質量比(ε)に対してダスト総質量Ｓ遵圡奭奭 颸辶㫳サイズ分布（離散的サイズ分布で，最大・最小限
界や粒柳種の数を変化させる）を振り，成長率Ｓ 卋槗【本研究で考える状況と手法】• Epstein則領域（ストークス数 Ts ∝ a）を考え，[Ts, min, Ts, max]の範囲を
Log等間隔でN種類のダストサイズを考える．

• ガスダスト質量比をεとし，εを振る．ダストにN種類のサイズがある
とき，一サイズあたりのガスダスト質量比εi = ε/N になる(i=1,..., N)．

【結果】多くの場合で，シングルサイズの場合より成長率がかなり遅くな• 本論文では，dn/da ∝ a^{-3.5}に限るが，将来的な一般化は容易．Nはる．Streaming Instability が有効なのは，ダスト濃集が起こる(ε~1)，特定の最大2048だが，Nが大きくなるにつれ成長率が収束する場合は，連続

サイズのダストが集中する(シングルサイズに近い)ときのみ？

分布の良い見積もりが得られたと考えられる．

• シアリングボックス内でN種のダストが混じった流体系の線形解析．
【背景】• 円挽 養翆

方法は別論文に出ている(Benitez-Llambay et al. 2019)が，ポイントは
翆

榆藴

悗

琰瓜輩槗

• Streaming Instabilityは固体微惑星の形成プロセスとして有望視．
• ガスとダストの２流体中でダスト密度の擾乱が発生すると，ダスト摩擦に
よりガス運動が変化，ダスト高密度領域での落下速度が低下（両者がなじむ）→ ダスト濃集領域が拡大，やがて自己重力で凝集できるほどに…
合ってますか？＞理論の方 (see Youdin & Goodman 2005）• 時間スケール：dynamical timeよりは長いが落下時間よりは短い．

• 計算結果はたくさんあるが，線形領域での系統的な解析が欠けていた．
• より現実的（かつ多様）な状況下での評価の基礎• 数値計算結果のベンチマーク

N種のダスト＋ガスに対して，固有モードの時間発展が̂ , k )e i(kx x+kzz)−ωt]に従うと考えて，4(N+1)次元の固有ベクトルRe[δ f(kx ŷ , k )とそれに対応したω(k , k )
δ f(k
x z
x z
つ不安定(Re(ω) < 0の中から)Ｓ亻槗【線形解析の妥当性を数値計算で検証】

線形解析でのN vs. σwhen the number of dust species doubles from N=2 toN=2048. For the case DTsI , the maximum growth rate

【より広いパラメータ空間での最大成長率σの線形解析結果】【単一ダストサイズ vs. N=2048 との比較】

単一サイズの場合（左）と比べると，
サイズ分布を考えた方が成長率がずっと長くなり，成長率が大きい領域はεが大きい領域に限定される．
maximum varies, i.e., DTs, min = [10-4 , Ts,max] with Ts,max Î{10-3, 10-2 , 10-1, 1}, or the maximum is ﬁxed and theminimum varies, i.e., DTs,max = [Ts,min, 1] with Ts,min ={10-1, 10-2 , 10-3, 10-4}. For each of these intervals in Stokes
numbers, we consider an increasing number of dust species by
doubling N from 2 to 2048 while keeping the dust-to-gas mass
ratio characterizing the distribution constant. This procedure

rate obtained in (Kx, Kz)-space, following the method describedin Section 2.
The most relevant outcomes are (i) the growth rate of themost unstable modes corresponding to the majority of the
distributions with low dust-to-gas mass ratios ò0.1 have notconverged and decreases below 10−3Ω0, independent of Ts,max.
In particular, when Ts,min=10−4 is ﬁxed, (Figure 4, upperponds to σ=10−3Ω0 in all panels.

【成長率
σ vs. １サイズ種あたりのダストガス質量比ε1 】
3. Outcome
and Implications

e have provided the ﬁrst systematic study of the linearεを固定してNを変化させた時の振る舞い
wth of the multi-species streaming instability. We foundσ ∝oneϵ1hand,
ガスが十分多い極限では，，単一ダストサイ
different- typesof behaviors. On thethere areibutions for which convergence of the growth rates is
ズの場合とほぼ一致：単純にN種類の２流体不安定の
hed by considering between a handful and a couple of
重ね合わせ
red dust species.
In the majority of cases we considered,ever, we-wereonly able to ﬁnd upper limits to the growthσ ∝ ϵ1，
それ以外，ダスト質量分布が効く領域では，, which are,in many cases, well below the values obtained
単一サイズの場合から大幅にずれる．
n only one dust
speciesisinvolved.
This
resultisbetter- 理由その１：各サイズのダストドリフト速度はeciated in Figure 5, where we show the maximum growth
トータルのダスト質量で決まるため，単一ダスト
for the classical (gas and one dust species) streaming
bility (leftmostサイズの時とは異なる．
panel) together with those obtained for the- 理由その２：ガスに対する摂動が増えるに連れ，
ibutions with
2048 species studied in Section 2.5 (center
異なるサイズのダスト間の結合が強まりお互いのrightmost panels).
elow, we brieﬂy
discuss some of the most relevant aspects
干渉が無視できなくなる．
consequences of our ﬁndings.
rowth rate decay and connection with resonant dragbilities. For aダストのサイズ分布を考慮すると，
large fraction of the particle distributions we
ed, the maximum growth rate σ decreases as the number ofStreaming
Instabilityが効くのは
ies N increases (and the mass per species decreases). A
rete example is これまでの予想よりは狭く，illustrated in Figure 2, where the growthII-4 , 1] decreases
D
T
=
[
10
he distribution
with
ò=1
and
εが大きく，ダストサイズが揃っている領域にs/N for sufﬁciently large N. It is natural to compare this to
限られるかもしれない
growth obtained when considering a single dust species
a dust-to-gas mass ratio ò1∼ò/N corresponding to the

Figure 6. Maximum growth rate as a function of the dust-to-gas mass ratio
ò1∼ò/N of the species with Stokes number that leads to the fastest growthwhen considered in isolation, Ts=1 in this case (solid black curve). The
dashed colored curves correspond to distributions with different dust-to-gas
mass ratios (each ﬁlled circle is obtained, from right to left, by doubling thenumber of species N from 2 to 256).

35. The anomalously low (sub)millimeter spectral indices ofsome protoplanetary disks may be explained by dust selfscattering
H. “Baobab” Liu (ApJL, 877, L22)
https://arxiv.org/abs/1904.00333【本研究でやったこと】• 2015.A.00005.S (PI: Tsukagoshi)と2013.1.00114.S (PI: Oberg)を再解
析，確かにα < 2の領域があることを確認• Birnstiel+ (2018)に倣ったダストモデルと光学的に厚いslabを想定し，【内容】最近のALMA観測で捕らえられた，光学的に厚いプランク関数

amaxを振りながら近似的な輻射輸送計算（エディントン近似的なも（レイリージーンズ領域）よりさらにシャローなスペクトル指数を，ダストの大きな散乱アルベドで説明しようという提案．これが事実ならス
ペクトル指数を広い波長範囲で計るだけでダストサイズ分布がわかる？

の）により各所のスペクトル指数再現を試みる．

• 近似的な輻射輸送計算の妥当性を定性的にチェックすべく，RADMC3Dを用いた計算も行う．

with independent measurements), and from HD 163296 (Dentet al. 2019).
【背景】TW HyaやHD163296などで，円盤内域で表面輝度 Iν∝ να として
Based on simpliﬁed radiative transfer models, the work
presented here argues that if we take scattering opacity into
α < 2の領域が捉えられている． テクニカルな分野の権威による仕事
consideration, it is possible to reproduce the anomalously low
values at (sub)millimeter bands from an isothermal, high(e.g., Dent+2019)もあり，較正エラーやartifact とは考えにくい． αopticaldepth dust emission source with amax∼0.1 mm. The【表面輝度およびαの再現】
analysis will be comparedspeciﬁcally to Atacama Large
Millimeter Array (ALMA) observations of the Class II
protoplanetary disk, TW Hya (d ∼ 60 pc; Gaia Collaboration
et al. 2016, 2018). Thanks to its approximately face-on
projection (for more information of this target source, see Qi
et al. 2004; Andrews et al. 2016, and references therein), it may
be sufﬁcient to consider the analytic solution of radiativetransfer equation for a thin slab, without requiring full threedimensional Monte Carlo radiative transfer modeling. Therefore, the analysis can be based on fewer free parameters and theresults would remain robust and comprehensive. In addition,
there is less concern about the confusion of free–free emission
thanks to the low bolometric luminosity and low protostellar
mass of TW Hya.
The observational data used in this work are brieﬂy
30au
introduced in Section 2.1 while more details are given in the
Appendix. The analysis of the SEDs is provided in Section 2.2.
Section 3 discusses the general implication of this work to other
Tsukagoshi et al. observational(2017） case studies, while our conclusion is nearly
identical to the Abstract.

ALMA偏光観測でも散乱の重要性が指摘（e.g., Kataoka+ 2015)
2. Data Analysis

【用いたダスト光学特性】

has α=2.0 in the Rayleigh–Jeans limit. Here we can clearly
see that when dust scattering is taken into account, and whenamax∼0.1 mm, the anomalously low values of α<2.0 are
reproduced at millimeter wavelengths, manifesting as a ﬂatter
SED than that of the ordinary blackbody. The feature ofanomalously low α shifts to centimeter wavelengths whenamax∼1 mm.
Markov chain Monte Carlo (MCMC) ﬁttings to the Tbr
proﬁles of TW Hya (Figure 1) were carried out to examine

【RADMC-3Dによる近似輻射輸送計算の妥当性の定性的チェック】

rnal Letters, 877:L22 (7pp), 2019 June 1⎛ r ⎞-0.4z) = 50[K] ´ ⎜⎟ ,⎝ [10 au] ⎠

Liu

(4 )

muthal angle. Our simulations assumed
gradients to what was actually observed
higher absolute temperature scale. This wasurpose is to test whether or not including
散乱なし
indeed lead to theanomalously low α
r dust temperatures can avoid producing
non-Rayleigh–Jeans effects. In addition,
ng dust temperature based on radiative
the assumed radial temperature proﬁle.
s because on the spatial scales of ouran be heated due to viscous dissipation isaddition, to simulate anisotropic dust
cally very thick limit, we need to use fullrids with rather small grid sizes, which
temperature evaluation computationallyasible for us. On the other hand, we dotions with and without dust scattering tot temperature proﬁles, which will in turnn about the effects of dust scattering on α.
f our simulations were deﬁned in sphericaliform intervals of polar angle, azimuthal
adius. Using the RADMC-3D code, we
n view of the 散乱あり
disk at 232.990 andsimulated images have some numerical∼3 au radius due to the very rapid changes
ity with the radius and polar angle, which
2 errors of the derived spectral indices.
ed the inner 3.2 au radius in the simulated
ows the derived α distributions from thesecases with and without switching on
the case without switching on scattering,n the innermost, high optical depth region;
中心部分は計算のswitched on, α can converge to values
hich supports our arguments about the
pplying the信頼性が落ちるの
analytic thin slab solution.
3. Discussion でマスク

(2016) reported that α has a value ∼3.7 at

Σd(r)とT(r)を手で与えた静水圧平衡モデル円盤で輻射輸送計算を実行

散乱を効かせた状態（下）では，確かにα < 2が再現できる．

一般に，散乱入りの厳密な輻射輸送計算は難しいが，温度勾配があまり急でない光学的に厚い円盤なら，slabモデルに基づく近似的な輻射輸送計算でもFigure 5. Spectral index (α) in between 144.988 and 232.990 GHz derived
from the RADMC-3D simulations for a face-on disk. Top and bottom panels
show the cases with no scattering and with anisotropic scattering evaluatedbased on Henyey–Greenstein approximation, respectively. These simulations
were based on the identical assumption of density and temperature distributions(see Section 2).

良いだろう．

36. Large-scale periodic velocity oscillation in the filamentary
cloud G350.54+0.69【全体像・観測諸元】• 距離 1.38±0.34 kpc(1’ ≈ 0.4 pc)，2本のフィラメント N(長さ5.9pc)，

M. Hong-Li Liu et al. (MNRAS, 487, 1259)
https://arxiv.org/abs/1905.08292

S(長さ2.3pc)で構成．ダスト質量は810M⦿(N)，110M⦿(S)．

【内容】Herschelダスト連続波で見つかったフィラメントG350.54+0.69について，APEXで13CO&C18O(J=2-1)でマッピング，速度構造から，力学平衡や速度構造を議論．特に系の中心速度の周期変動を発見し，その起源を議論．

• ダスト連続波のデータから，フィラメント中には低質量星原始星に付随する９つのコアが同定→ 低質量星形成領域．

• 速度分解能 = 0.15 km/s, ビームサイズ 28” (FWHM), or ~ 0.2pc.
• 速度的に２成分に分けられる．C18Oの方がより細い(高密度領域)をトレース．
Large-scale periodic velocity oscillation in G350.54+0.69

2

H.L. Liu et al.
2
CO
C18O

1

I (K km s 1)
2 3 4 5 6

7(b)

I (K km s 1)
0.5 1.0 1.5 2.0 2.5 3.0 3.5
(c)

24

1

Dec (J2000)

Tmb (K)10

36 20 (a)13

2

I (K km s 1)
4
6
8

3

0
15

10

5
Vlsr (km s 1)

0

5

28
32
36

Figure 1. Average spectra of 13 CO, and C18 O (2-1) over the G350.5 filamentary cloud. The black color stands for the observed spectrum. The redline is the Gaussian fitting to the spectrum of 13 CO, which is a sum of thetwo Gaussian components (blue lines), the weak-emission one from 9.0
to 5.7 km s 1 , and the major one from 5.7 to 2 km s 1 . The green lineis the single-component Gaussian fitting to the spectrum of C18 O, whichmatches the major component of the 13 CO spectrum. The vertical dashedblack line represents a systematic velocity of 3.9 km s 1 .

40s

20s

18m00s17h17m40s

RA (J2000)

40s

20s

18m00s17h17m40s

RA (J2000)

40s

20s

18m00s17h17m40s

RA (J2000)

Figure 2. (a–b:) Velocity-integrated intensity maps of the two velocity components of 13 CO (2-1) toward G350.5, the major one from 5.7 to 2 km s 1 ,and the weak-emission one from 9 to 5.7 km s 1 . (c:) Velocity-integrated intensity of C18 O (2-1), integrated over 7 to 1 km s 1 . The pixel size of thethree maps is reduced to 700 , a quarter of the beam size, for a better visualization. The cyan curves in all panels represent the spines of the two discontinuous
filaments as identified in Paper I and the black contours stand for the NH2 column density (see Paper I), starting from 1.8 ⇥ 1022 cm 2 with a step of
0.6 ⇥ 1022 cm 2 ( = 0.04 ⇥ 1022 cm 2 ).

3.3

Velocity dispersion along the filament measured fromcosmicbackgroundradiation;
⌧
is
the
o
Tackenberg J., et al., 2014, A&A, 565, A101
1318
both CO and C O
速度幅の分布（C18Oでも十分なS/Nで検出されている領域について，
of the telescope beam filled by emission
Tafalla M., Hacar A., 2015, A&A, 574, A104
4
H.L.
Liu
et
al.
The G.
observed
velocity( obs ) is constantmeasuredandfromVázquez-Semadeni E., Palau A.,
Ballesteros-Paredes J., Gómez
C.,and µdispersion
are the rotational
thebop
複数速度成分を切り分けたガウスフィットで）13
18
in the mainrespectively.
structure of thefilament
whe
Zamora-Avilés M., 2019, arXiv e-prints, arXiv:1903.11247 CO and C
M3DO (2-1)
of molecules
Given
the opti
13
0.0
1.9 species
3.0are well-detected.
4.0
4.9 In principle, obs can be estimat
both
CO
Yuan J., et al., 2016, ApJ, 820, 37
10
Tmb , weobtainedtheexcitationtempe
10τの補正をすると，
C O(b) from the second-order moment
C O
of the PPV data cube. Howev
Yuan J., et al., (a)2018, ApJ, 852, 12
CO a result, the average excitation
CO
Eq.
B1.
As
this G.
methoddue to the simp
この領域では概ねWilliams G.8 M., Peretto N., Avison A., Duarte-Cabral 8A., FullerA., will cause additional uncertainties
τの補正あり
1)
in
the
main
structureofthefilamentis1
τの補正なしtreatmentofmultiple-velocity
components
as
one.
To
betteres
2018, A&A, 613, A11
非熱的成分は
6
6
13the corresponding average dust
less
than
CO 2-1, we define 47 positions alo
obs especially from
Wu G., Qiu K., Esimbek J., Zheng X., Henkel C., Li D., HanmateX., 2018,
熱的成分の2倍程度
fromthe dust
temperature
map6 (PaperI).
the ridgeline of thefilament,
41 in
G350.5-N andin G350.5
4
4
A&A, 616,
A111
13
Neighbour positions
areto
separated
one pixel, halfofof thebea(
lated
non fullbythermalization
CO
2
13
18
2
size of both COof
andG350.5.
C O maps. N(Histhenobtainedfromtobs )からBを見積もり，
2
13
18
Gaussian fitting to the spectra of both CO, and C O (2-1) at t
0
0
1.0
1.5
2.0
2.5
3.0
1
2
3
4
ビリアル解析(次頁)
APPENDIX A: OPTICAL
DEPTH
selected
47
positions.
The
fitting
plots
are
shown
in
Fig.
C1.
/
/
tot,intthtot,obth
すると， from bo
The velocity dispersion combines the contributions
13 CO (red) and C18 O (2-1) (blue) for G350.5. In panel (a) is the observed velocity
Figure 4.theHistogramof the total velocity
dispersionequations
measured from bothwith
Followingradiative
transferanassumptionof
APPENDIXC:GAUSSIAN
FITTINGthermalandnon-thermalgasmotionsinturbulentclouds.
Howev
重力エネルギーが卓越．
13
18
dispersion. The meanvaluesintheobserved
velocity
dispersions
measured
from
CO
and
C
O
are
2.9
±
0.4
(red
dashedline),
and
2.0
±
0.3
th
th
18
13
optically
thinline),
C respectively.
O (2-1)(e.g.,(intrinsic)
Gardenal. 1991),
weare 2.1 ± 0.2 can(blue dashed
In panelemission(b) is the opacity-correctedvelocityet
dispersion.
The meancontributions
values(red dashedth frombeCO
theiroverestimateddue to the line broadeni
磁場と乱流の寄与は
1813line), and 1.9 ± 0.3 th from C O (blue dashed line). Theof the velocity dispersions of CO before and after the opacity correction
13different distributions 18
calculated
thethatoptical
depth) of ofCO,andC dispersion.
O (2-1)forthe depth.
opticalSuch overestimate
can be expressed analytica
demonstrate
the line opacity
can cause (⌧the broadening
the intrinsic
velocity
Theby
intrinsic
velocity
dispersion distributions
indicate that the同じくらいoverall turbulent motions of the filament G350.5 are supersonic (i.e., the 3D Mach number follows M3D > 2).
A
This
paper
has
been
typeset
from
a
T
X/L
TEX
as
a
functionoftheopticaldepth(Phillips
etal.
1979):
selected positions in the main structure of G350.5 as below:
E
"
!#
τの補正
1/2
13
hydrogen, and µp = CO
2.33 is the mean molecular⌧13 weight per free ⌧13 (M ), and gravitational potential (W ). Tnet accounts for the difTmb ratio N (H)/N
1 (He)
e = 10 and a1negligi-e ference between the internal (Tint ) and external
1 (Text ) turbulent ⌧particle for an abundance
τの導出
R=7.7
⇡
⇡
,
(A1)
p
=
ln
,
(
int
obs
18
ble admixture of metals(e.g.,
Kauffmannetal.
2008).
Thegas
kienergy,
assumingthatmolecular
cloudsareconfinedbytheexter⌧
⌧
/R
C O
2
18
13
1
e
1
e
T
ln
netic temperaturemb
here was assumed to be equal to the average dust
nal pressure rather than completely isolated. ln2
Note that all energies e ⌧ +1
18

18

13

temperature (Td ⇠ 18 K) of the filament instead of the excitation

NumberNumber13

here are measured per unit length.

given the 1D measurement NT,int , and assuming isotropic tu
フィラメント単位長さあたりでのビリアル解析（W:
重力，2T
運動＋外圧，M:
net:inlence in three dimensions(e.g.,
Kainulainen& Federrath
2017
Large-scale periodic
velocity
oscillation
G350.54+0.69
5 磁場)

4

3.4

W
2Tnet

45
10
3
(b) G350.5-S
2Virial analysis along the filament
M
2Tnet + M

Engery (erg pc 1)

Engery (erg pc 1)

45106(a) G350.5-N

W
2Tnet
M
2Tnet + M

1

The complex interactions between turbulence, gravity and m
0
neticfieldshappeneverywhereinmolecular
clouds
and
reg
0
1
C1
star
formation
We can describe these interactions thro
C8 C7B
C7A
C6 C5
C4 C3 therein.
C2
2
2
0
1the 2virial
3 theorem.
4
5
0.0
0.5
1.0
1.5
2.0
According
to
Fiege
&
Pudritz(2000a),Distance along filament (pc)Distance along filament (pc)virialequationofself-gravitating,magnetizedturbulent
filamenFigure 5. Virial analysis through the comparisons between turbulent (T ), magnetic energies (M ) and gravitational potential (W ) for G350.5-N (in panel a),
and for G350.5-S (in panel b). Several peaks are related to dense cores (also see Fig. 3) as indicated by vertical lines. The comparison shows that the gravity
cloudscanbeis awritten
inglobalthe
form:
dominates over both turbulence
and magnetic
fields, andmajor driver to thefragmentationof G350.5 from clouds to dense cores.
22T+ W + M = 0.

where C 力学平衡フィラメントの場合は，is a constant, and NH is the column density of atomic hylation
netin G350.5-N. Since G350.5-S is disconnected to G350.5-N
drogen gas. This constant was given to be 3.8 by Crutcher (2012)
in Fig. 3 and G350.5-S itself has insufficient detection in the diffuse
under the assumption of magnetic critical condition and a sphericalregion, the nature of the velocity field of G350.5-S requires detailed
本フィラメントの場合は，概ね，|W|> 2Tnet+ M（特にコア領域），重力が卓越
shapeof clouds, and improved to be 1.9 by Li et al. (2014) asinvestigation
with sensitive and higher-resolution observations. To
suming a sheet-like shape. The mean column density of the cloud
better visualize the velocity oscillation along G350.5-N, we make

Therefore, the virial equilibrium of the clouds depends on the cFigure 2. (a–b:) Velocity-integrated intensity maps of the two velocity components of 13 CO (2-1) toward G350.5, the major one from 5.7 to 2 km s 1 ,and the weak-emission one from 9 to 5.7 km s 1 . (c:) Velocity-integrated intensity of C18 O (2-1), integrated over 7 to 1 km s 1 . The pixel size of thethree maps is reduced to 700 , a quarter of the beam size, for a better visualization. The cyan curves in all panels represent the spines of the two discontinuous
filaments as identified in Paper I and the black contours stand for the NH2 column density (see Paper I), starting from 1.8 ⇥ 1022 cm 2 with a step of
0.6 ⇥ 1022 cm 2 ( = 0.04 ⇥ 1022 cm 2 ).
6 H.L. Liu et al.

中心速度の分布：周期的な変動が見られる

4.0

3.8

3.6

3.4

V (km s 1 )

平均速度図

36 21
C8

24 C7B

G350.5-N

Dec (J2000)

C7A

C6

27

C5
C4
C3

G350.5-S

30
C2

1 pc

33
45s

C1

30s

15s

18m00s 17h17m45s

RA (J2000)

Figure 3. Velocity centroid map of C18 O (2-1) for the filamentary cloudG350.5. Only the main filament is shown because of sufficient detectionof C18 O (2-1) therein. The pixel size of the map is 1400 , half of the beamsize, as recommended by the APEX observatory. All pixels correspond to
detection greater than 3 (i.e., = 0.4 K km s 1 ). The red and blue color
scales can be regarded as the red and blue-shifted velocities with respect tothe systematic velocity 3.9 km s 1 . A large-scale periodic velocity fluctuation appears almost along the main structure. The black contours represent
the NH2 column density, starting from 1.8 ⇥ 1022 cm 2 with a step of
0.6 ⇥ 1022 cm 2 . The dense cores identified in Paper I are also labelled as
C1–C8.

3.3

2.5

3.0

Velocity dispersion along the filament measured fromTmb (K)18
both 13 CO andCO
2The observed velocity(a)dispersion( obs ) is measured from bothPV
diagram
13
CO and C18 O (2-1) in the main structure of the filament where
3
both species are well-detected.
In principle, obs can be estimated1.32 pc
from the second-order moment of the PPV data
cube. However,this method will cause additional uncertainties due to the simple
treatment of multiple-velocity
components as one. To better esti4
mate obs especially from 13 CO 2-1, we define 47 positions along
the ridgeline of the filament, 41 in G350.5-N and 6 in G350.5-S.
Neighbour positions
are separated by one pixel, half of the beam
5
size of both 13 CO and C18 O maps. obs is then obtained from theGaussian fitting to the spectra of both 13 CO, and C18 O (2-1) at the
selected 47 positions. The
plots arein Fig. C1.
(b) fitting
NH2 alongtheshown
filament
C5bothThe velocity dispersion combines the contributions fromthermal and non-thermal gas motions in turbulent clouds. However,C7a
C6 broadening
their contributions can be overestimated due to the line
4 overestimate can be expressed analyticallyby optical depth. SuchC7b
as a function of the optical depth (Phillips et al. 1979):

NH2 (1022 cm 2)

4.2

2.0

3.5

4.0
Fit curve
Blue-shifted fit
Observed
Intensity-weighted velocity

PV図V (km s 1)

4.4

1.5

C8
"

1
ln
obs = p
ln2

2

⌧
ln

2
e ⌧ +1

!#1/2

int ,

(1)

Cubic interp.
Nearest-neighbour interp.

VLSRの変化点と

C4
C3

C2

柱密度（コア）のピークが一致しているように見えるwhere int is the intrinsic velocity dispersion unaffected by optical
0 A). For comparison,
1245
depth ⌧ (see Appendix
we calculate
the gas 3
dispersions (i.e., non-thermal dispersion NT , and the total ddisper(pc)sion tot ) based on the observed, and intrinsic velocity dispersions
Figure 6. from(a-b:) Velocity
columnextracted
the spine of the filament G350.5-N. (a): Position-velocity diagram of C18 O (2-1).
measuredboth 13andCOand density
C18 O,distributions
respectively,
as along
follows:

The detections greater than 3 ( q= 0.42 K) are shown in grey-scale within the velocity channels from 5.0 to 2.5 km s 1 . The cyan crosses are the1weighted average velocities by brightness
kTkand fitted with a sine function with a wavelength of ⇠ 1.3 pc and an amplitude of ⇠ 0.12 km s .
2 temperatures,,
NT =Literally, the blue color curves
stand forobs/int
the blue-shifted
mvelocities with respect to the systematic velocity of 3.9 km s 1 , and the red ones for the red-shiftedqvelocities. (b:) Density distribution. Density peaks relatedi.e., C2 to C8, which are extracted from Fig. 3. The velocity and column
kTk to dense cores are indicated,(2)=,
thone-to-one
density distributions do not show a good
between one another.
µpcorrespondence
mH
tot =

p

2
NT +

2
th ,

vanishing velocity. Assuming that both density and velocity pertur-

since (random) small-scale kinks would not maintain a large-scalewhere
can be replaced either withor with(seebationsobs/int(oscillations) are sinusoidal (periodic), a /4obs
shift between int coherent, periodic oscillation. On the other hand, a regular but osEq.them1), k
constant,thegasdetails).
kinetic
k isforcanisbeBoltzmanns
expected (Hacar& Tafalla T
2011
more
Thetemperacillating geometry driven by some physical mechanism could beLarge-scale periodic velocity oscillation in G350.54+0.697

中心速度の周期的な変動の起源は？
(a) Model 1 (core formation)(b) Model 2 (oscilating filament)(Model 1) フィラメント状クラウド内の重力不安定による分裂
Observer◎フィラメント内のコアの間隔が理論と良い一致
× ρのピークと速度の変曲点の「位相ずれ」が見えない(Model 2) 何らかの振動，Large-Scale
0
V

0
V

0

0

-V

-V
L (distance)

MHD-transverse wave ?

両者の組み合わせかもしれないL (distance)

Figure 7. (a:) schematic representation of GI-driven core formation along a filament, which is adapted from Fig. 12 of Hacar & Tafalla (2011). Due to GI,
the motions of core-forming gas converge towards the core centers where the density enhances but the relative motions become static with respect to thesystematic velocity. In this scenario, assuming that both density ( ⇢) and velocity ( V ) follow sinusoidal distributions, a /4 phase shift between them would
be expected. (b): schematic representation of a physical oscillating filament (e.g., induced by an MHD wave). The filament moves toward and away from an
observer. In this scenario, the correspondence between the velocity extremes (i.e., red/blue-shifted velocities) and the density enhancements is expected if thefilament (density) oscillates sinusoidally. The red and blue colors indicate the red, and blue-shifted velocities (motions) with respect to an observer. Note thatthe physical intensities of ⇢ and V are ignored in the schematic demonstration.

28

ALMA Observations of the rho Ophiuchus B2 Region: I. Molecular Outflows and Their
Driving Sources
Takeshi Kamazaki1 , Fumitaka Nakamura1,2,3 , Ryohei Kawabe1,2,3 , Chihomi Hara3 , Shigehisa Takakuwa4 ,Naomi Hirano5 , James Di Francesco6 , Rachel Friesen7 and Motohide Tamura3
1

National Astronomical Observatory, Mitaka, Tokyo 181-8588, Japan; 2 The Graduate University for Advanced
Studies (SOKENDAI), 2-21-1 Osawa, Mitaka, Tokyo 181-0015, Japan; 3 Department of Astronomy, Graduate Schoolof Science, The University of Tokyo, Tokyo 113-0033, Japan; 4 Department of Physics and Astronomy, Graduate
School of Science and Engineering, Kagoshima University, 1-21-35 Korimoto, Kagoshima, Kagoshima 890-0065, Japan;
5
Academia Sinica Institute of Astronomy and Astrophysics, P.O. Box 23-141, Taipei 106, Taiwan; 6 NRC Herzberg
Inst of Astrophysics, 5071 W Saanich Rd, Victoria BC V9E 2E7, British Columbia, Canada; 7 Dunlap Institute for
Astronomy and Astrophysics, the University of Toronto, 50 St George St, Toronto ON M5S 3H4, Canada
E-mail contact: kamazaki.takeshi at nao.ac.jp
We present the results of ALMA mosaic observations of dust continuum and 12 CO (J = 2 − 1), 13 CO (J = 2 − 1) and
C18 O (J = 2 − 1) molecular lines toward the rho Ophiuchus B2 region. The 1.3-mm dust-continuum image made fromthe combined 12-m and 7-m Array data reveals not only the dense cores identified by past single-dish observations, butalso their detailed internal substructures. The 12 CO (J = 2 − 1) images show very complex structures of protostellar
outflows. They suggest that the gigantic outflow lobes in Oph B2 are presumably driven at least by two protostars,
EL 32 and EL 33. We do not detect clear high velocity components associated with the Class I protostar SST c2d
J162730.9-242733 and the Flat Spectrum object SSTc2d J162721.8-242727. In addition, we find interesting striations
with 15arcsec (approx 2000 au) separations in both the 12 CO and 13 CO channel maps. The CO striations appear to
be roughly parallel to the magnetic field direction, and we speculate that the directions of the striations may follow
the magnetic field in the envelope of Oph B2.
Accepted by ApJ

https://iopscience.iop.org/article/10.3847/1538-4357/aaf857• ALMA-ACAのρOph領域のモザイクマップ
Extremely
Dense Cores Associated with Chandra Sources in Ophiuchus A: Forming
Brown
Dwarfs Unveiled?
• コアやアウトフローの同定
1

2

1

1

1

Extremely Dense Cores Associated with Chandra Sources in Ophiuchus A: Forming
Brown Dwarfs Unveiled?

29 Ryohei Kawabe , Chihomi Hara , Fumitaka Nakamura , Kazuya Saigo , Takeshi Kamazaki , Yoshito
1

2

1

1

1

Shimajiri3 , Kengo Tomida4 , Shigehisa Takayuwa5 , Yohko Tsuboi6 , Masahiro N Machida7 , James Di
Francesco8 , Rachel Friesen9 , Naomi Hirano10 , Yumiko Oasa11 , Motohide Tamura2 , Yoichi Tamura12 ,Takashi Tsukagoshi1 and David Wilner13
1

National Astronomical Observatory, Mitaka, Tokyo 181-8588, Japan; 2 Department of Astronomy, School of Science,
University of Tokyo, Bunkyo, Tokyo, 113-0033, Japan; 3 Laboratoire AIM, CEA/DSM-CNRS-Universit Paris Diderot,
IRFU/Service d’Astrophysique, CEA, Saclay, F-91191 Gif-sur-Yvette, France; 4 Department of Earth and Space
Science, Osaka University, Toyonaka, Osaka 560-0043, Japan; 5 Department of Physics and Astronomy, Graduate
School of Science and Engineering, Kagoshima University, 1-21-35 Korimoto, Kagoshima, Kagoshima 890-0065, Japan;
6
Department of Physics, Faculty of Science and Engineering, Chuo University, 1-13-27 Kasuga, Bunkyo-ku, Tokyo 1128551, Japan; 7 Department of Earth and Planetary Science, Faculty of Science, Kyushu University, Hakozaki 6-10-1,
8Higashi-ku, Fukuoka, 812-8581, Japan; NRC Herzberg Inst of Astrophysics, 5071 W Saanich Road, Victoria, BC V9E
2E7, British Columbia, Canada; 9 North American ALMA
24 Science Center, National Radio Astronomy Observatory,
520 Edgemont Road, Charlottesville, VA 22903; 10 Academia Sinica Institute of Astronomy and Astrophysics, P.O.
Box 23-141, Tapei 10617, Taiwan; 11 Faculty of Education, Saitama University, 255 Shimo-Okubo, Sakura-ku, Saitama
338-8570, Japan; 12 Division of Particle and Astrophysical Science, Graduate School of Science, Nagoya University,
Aichi 464-8602, Japan; 13 Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138,
USA• 他波長データ解析により，ρOph領域のProtostar candidates について詳細な議論

E-mail contact: ryo.kawabe at nao.ac.jp
On the basis of various data such as ALMA, JVLA, Chandra, it Herschel, and it Spitzer, we confirmed that two
protostellar candidates in Oph-A are bona fide protostars or proto-brown dwarfs (proto-BDs) in extremely early
evolutionary stages. Both objects are barely visible across infrared (IR, i.e., near-IR to far-IR) bands. The physicalnature of the cores is very similar to that expected in first hydrostatic cores (FHSCs), objects theoretically predictedin the evolutionary phase prior to stellar core formation with gas densities of sim 1011−12 cm−3 . This suggests thatthe evolutionary stage is close to the FHSC formation phase. The two objects are associated with faint X-ray sources,
suggesting that they are in very early phase of stellar core formation with magnetic activity. In addition, we found the
CO outflow components around both sources which may originate from the young outflows driven by these sources.
The masses of these objects are calculated to be sim0.01 − 0.03 Modot from the dust continuum emission. Their
physical properties are consistent with that expected from the numerical model of forming brown dwarfs. These facts(the X-ray detection, CO outflow association, and FHSC-like spectral energy distributions) strongly indicate that thetwo objects are proto-BDs or will be in the very early phase of protostars which will evolve more massive protostars
if they gain enough mass from the surroundings. The ages of these two objects are likely to be within sim103 years
after the protostellar core (or second core) formation, taking into account the outflow dynamical times (lesssim 500
yrs).
Accepted by ApJ

https://iopscience.iop.org/article/10.3847/1538-4357/aae153

31

Magnetic activity of the young solar analogue V1358 Ori
L. Kriskovics1 , Zs. Kövári1 , K. Vida1 , K. Oláh1 , T.A. Carroll2 , and T. Granzer2
1

Konkoly Observatory, Research Center for Astronomy and Earth Sciences, Budapest, Hungary; 2 Leibniz Institutefor Astrophysics (AIP), Potsdam, Germany
E-mail contact: kriskovics.levente at csfk.mta.hu
25
Young, fast rotating single stars can show dramatically diﬀerent magnetic signatures and levels of magnetic activity
as compared with the Sun. While losing angular momentum due to magnetic breaking and mass loss through stellar
winds, the stars gradually spin down resulting in decreasing levels of activity. Studying magnetic activity on such solar
analogues plays a key role in understanding the evolution of solar-like stars and allows a glimpse into the past of the Sun
as well. In order to widen our knowledge of the magnetic evolution of the Sun and solar-like stars, magnetic activity of
the young solar analogue V1358 Ori is investigated. Fourier analysis of long-term photometric data is used to derive
rotational period and activity cycle length, while spectral synthesis is applied on high resolution spectroscopic data inorder to derive precise astrophysical parameters. Doppler imaging is performed to recover surface temperature maps
for two subsequent intervals. Cross-correlation of the consecutive Doppler maps is used to derive surface diﬀerential
rotation. The rotational modulation of the chromospheric activity indicators is also investigated. An activity cycleof ≈1600 days is detected for V1358 Ori. Doppler imaging revealed a surface temperature distribution dominatedby a large polar cap with a few weaker features around the equator. This spot configuration is similar to othermaps of young solar analogues from the literature, and supports recent model predictions. We detected solar-likesurface diﬀerential rotation with a surface shear parameter of α = 0.016 ± 0.010 which fits pretty well to our recently
proposed empirical relation between rotation and diﬀerential rotation. The chromospheric activity indicators showeda rotational modulation.
https://www.aanda.org/articles/aa/abs/2019/07/aa34927-18/aa34927-18.html
Accepted by A&A
http://arxiv.org/pdf/1905.08259• 太陽的な若い星 V1358 Ori に対する18年間にわたる分光測光観測に基づき，変光の周期解析やドップラーイメージングを行い，磁気活動を議論

Magnetic field structure in Serpens South32

Magnetic field structure in Serpens South
Takayoshi Kusune1 , Fumitaka Nakamura1,2,3 , Koji Sugitani4 , Shuji Sato5 , Motohide Tamura3 , Jungmi
Kwon6 , Kazuhito Dobashi7 , Tomomi Shimoikura7 and Benjamin Wu1
1

National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan; 2 The Graduate
University for Advanced Studies (SOKENDAI), 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan; 3 Department of
Astronomy, University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan; 4 Graduate School of Natural
Sciences, Nagoya City University, Mizuho-ku, Nagoya, Aichi 467-8501, Japan; 5 Department of Astrophysics, Nagoya
University, Furo-cho, Chikusa-ku, Nagoya, Aichi 464-8602, Japan; 6 Institute of Space and Astronautical Science, Japan
Aerospace Exploration Agency, 3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa 252-5210, Japan; 7 Departmentof Astronomy and Earth Science, Tokyo Gakugei University, Koganei, Tokyo 184-8501, Japan
E-mail contact: takayoshi.kusune at nao.ac.jp
We made near-infrared polarimetric observations toward Serpens South. This region contains three dense filamentsthat are roughly parallel to one another. Using the histogram of relative orientations, the three filaments are foundto be roughly perpendicular to the global magnetic field. The morphology of the plane-of-sky (POS) magnetic field
and molecular gas suggests that the magnetic field plays an important role in the filament formation and evolution.
Applying the Davis-Chandrasekhar-Fermi method, the POS magnetic field strengths are estimated to be 30–140 >muG. The estimated magnetic field strengths indicate that the filaments are close to magnetically critical. Since
pretellar and protostellar cores already formed along the filaments, the magnetic diﬀusion such as a ambipolar diﬀusion
and turbulence reconnection is likely to play an important role in the formation of dense cores. We speculate that thefilaments are formed by fragmentation of a sheet-like cloud that was created through the gravitational contraction ofa magnetized, turbulent cloud.
Accepted by PASJ

• Serpens South をSIRPOL(NIR POL)とNRO45m/FOREST(COs) で観測，分子雲構造と磁場構造を議論

Ahttps://academic.oup.com/pasj/advance-article-abstract/doi/10.1093/pasj/psz040/5480475?
SiO J = 5→4 Survey Toward Massive Star Formation Regions
1,2,3

1,4

redirectedFrom=fulltext
52
1,3

Shanghuo Li, Junzhi Wang , Min Fang , Qizhou Zhang , Fei Li
Qingfeng Zhu8 , and Shaoshan Zeng91, Zhi-Yu Zhang6,7 , Juan Li1 ,2

33

A SiO J = 5→4 Survey Toward Massive Star Formation Regions
Shanghuo Li1,2,3 , Junzhi Wang1,4 , Min Fang5 , Qizhou Zhang2 , Fei Li1,3 , Zhi-Yu Zhang6,7 , Juan Li1 ,Qingfeng Zhu8 , and Shaoshan Zeng91

Shanghai Astronomical Observatory, Chinese Academy of Sciences, 80 Nandan Road, Shanghai 200030, China; 2Center for Astrophysics — Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138, USA; 3 University ofChinese Academy of Sciences, 19A Yuquanlu, Beijing 100049, China; 4 Key Laboratory of Radio Astronomy, Chinese
Academy of Sciences, 210008, Nanjing, China; 5 Department of Astronomy and Steward Observatory, University of
Arizona, 933 N Cherry Ave., Tucson, AZ 85721, USA; 6 Institute for Astronomy, University of Edinburgh, Blackford
26
Hill, Edinburgh, EH9 3HJ, UK; 7 European Southern Observatory,Karl-Schwarzschild-Strasse 2, D-85748 Garching,
Germany; 8 Astronomy Department, University of Science and Technology, Chinese Academy of Sciences, Hefei 210008,
China; 9 School of Physics and Astronomy, Queen Mary University of London, Mile End Road, E1 4NS London, UK
E-mail contact: shanghuo.li at gmail.com
We performed a survey in the SiO J = 5→4 line toward a sample of 199 Galactic massive star-forming regions at
diﬀerent evolutionary stages with the SMT 10 m and CSO 10.4 m telescopes. The sample consists of 44 infrared darkclouds (IRDCs), 86 protostellar candidates, and 69 young Hii regions. We detected SiO J = 5→4 line emission in
102 sources, with a detection rate of 57%, 37%, and 65% for IRDCs, protostellar candidates, and young Hii regions,
respectively. We find both broad line with Full Widths at Zero Power (FWZP) > 20 km s−1 and narrow line emissonsof SiO in objects at various evolutionary stages, likely associated with high-velocity shocks and low-velocity shocks,
respectively. The SiO luminosities do not show apparent diﬀerences among various evolutionary stages in our sample.
We find no correlation between the SiO abundance and the luminosity-to-mass ratio, indicating that the SiO abundance
does not vary significantly in regions at diﬀerent evolutionary stages of star formation.
Accepted by ApJ
http://arxiv.org/pdf/1905.08359• 大質量星形成領域199個(IRDC, Protostars, HII領域からなる)をSiO 5-4でサーベイ（CSO, SMT 10m) ，

統計的な性質を議論

CO Outflow Candidates Toward the W3/4/5 Complex I: The Sample and its Spatial
Distributionhttps://iopscience.iop.org/article/10.3847/1538-4357/ab1e4c
1,2

1

1

1

1,3

1

http://arxiv.org/pdf/1905.08359

34
CO Outflow Candidates Toward the W3/4/5 Complex I: The Sample and its Spatial
Distribution
Yingjie Li1,2 , Ye Xu1 , Yan Sun1 , Qing-Zeng Yan1 , Yuehui Ma1,3 and Ji Yang1
1

Purple Mountain Observatory, Chinese Academy of Sciences, Nanjing 210008, China; 2 University of Science andTechnology of China, Chinese Academy of Sciences, Hefei, Anhui 230026, China; 3 University of the Chinese Academyof Sciences, 19A Yuquan Road, Shijingshan District, Beijing 100049, Peoples Republic of China
E-mail contact: xuye at pmo.ac.cn
Using the Purple Mountain Observatory Delingha 13.7 m telescope, we conducted a large-scale 12 CO (J = 1–0)
outflow survey (over ∼110 deg2 ) toward the W3/4/5 complex and its surroundings. In total, 459 outflow candidates
were identified. Approximately 62% (284) were located in the Perseus arm, including W3/4/5 complex and its
surroundings, while ∼35% (162) were located in the Local arm, ∼1% (5) in the Outer arm, and ∼2% (8) in two
interarm regions. This result indicated that star formation was concentrated in the Galactic spiral arms. The detailed
spatial distribution of the outflow candidates showed that the Perseus arm presented the most active star formation
among the study regions. The W3/4/5 complex is a great region to research massive star formation in a triggered
environment. A key region, which has been well-studied by other researches, is in the eastern high-density W3 complex
that neighbors the W4 complex. Two shell-like structures in the Local arm contain candidates that can be used to
study the impact on star formation imposed by massive or intermediate-mass stars in relatively isolated systems. The
majority of outflow candidates in the two interarm regions and the Outer arm are located in filamentary structures.
Accepted by ApJS
http://arxiv.org/pdf/1905.02018• 紫金山天文台（デリンハ）の14m電波望遠鏡を用いた12COサーベイデータに基づき，当該領域のアウトフロー候補を抽出
https://iopscience.iop.org/article/10.3847/1538-4365/ab1e55
The anomalously
low (sub)millimeter spectral indices of some protoplanetary disks may
be explained by dust self-scatteringhttps://arxiv.org/pdf/1905.08292.pdf

37

A Large-field J=1-0 Survey of CO and Its Isotopologues Toward the Cassiopeia A
Supernova Remnant
Yuehui Ma1,2 , Hongchi Wang1 , Miaomiao Zhang1,3 , Chong Li1,2 and Ji Yang1
1

Purple Mountain Observatory and Key Laboratory of Radio Astronomy, Chinese Academy of Sciences, 2 West
Beijing Road, Nanjing 210008, China; 2 University of Chinese Academy of Sciences, 19A Yuquan Road, Shijingshan
District, Beijing 100049, China; 3 Max-Planck-Institut für Astronomie, Königstuhl 17, 69117 Heidelberg, Germany
E-mail contact: mayh at pmo.ac.cn
We have conducted a large-field simultaneous survey of 12 CO, 13 CO, and C18 O J = 1−0 emission toward the Cassiopeia
A (Cas A) supernova remnant (SNR), which covers a sky area of 3.5◦ × 3.1◦ . The Cas giant molecular cloud (GMC)
mainly consists of three individual clouds with masses on the order of 104 − 105 M⊙ . The total mass derived from the13
CO emission of the GMC is 2.1×105 M⊙ and is 9.5×105 M⊙ from the 12 CO emission. Two regions with broadened(6−7 km s−1 ) or asymmetric 12 CO line profiles are found in the vicinity (within a 10′ × 10′ region) of the Cas A SNR,
indicating possible interactions between the SNR and the GMC. Using the GAUSSCLUMPS algorithm, 547 13 COclumps are identified in the GMC, 54% of which are supercritical (i.e. αvir < 2). The mass spectrum of the molecular
clumps follows a power-law distribution with an exponent of −2.20. The pixel-by-pixel column density of the GMC
can be fitted with a log-normal probability distribution function (N-PDF). The median column density of molecularhydrogen in the GMC is 1.6 × 1021 cm−2 and half the mass of the GMC is contained in regions with H2 column density
lower than 3 × 1021 cm−2 , which is well below the threshold of star formation. The distribution of the YSO candidates
in the region shows no agglomeration.

• 紫金山天文台（デリンハ）の14m電波望遠鏡を用い，CasAを12CO/13CO/C18O(1-0) でマッピング，GMCの全貌を明らかにし，547 Clumpsを同定．

28

https://iopscience.iop.org/article/10.3847/1538-4357/ab1ea7