星形成ゼミ
SFN #322 (66-75本目)
百瀬宗武（茨城大・理）詳細解説は以下の2本 + 1本(時間があれば)• #70: Magnetospheric Accretion as a Source of Hα Emission from Proto-planets around PDS 70(T. Thanathibodee et al. ApJ, 885, 94)• #74: X-shooter spectroscopy of young stars with disks: The TW Hydrae association as a probeof the final stages of disk accretion (L. Venuti et al., A&Ap in press, arxiv.1909.06699)• #66: Investigating the gas-to-dust ratio in the protoplanetary disk of HD 142527 (K. Soon et al.
PASJ, e-open: https://doi.org/10.1093/pasj/psz112)
その他はひと言メモを参照

LETTERS

70: Magnetospheric Accretion as a Source of Hα Emission
from Proto-planets around PDS 70

a

Kバンド観測A&A 617, L2 (2018)

T. Thanathibodee et al. ApJ, 885, 94
N_ALC_YJH_S (185 mas in diameter) apodized-Lyot coronhttps://arxiv.org/abs/1909.06450
agraph (Martinez et al. 2009; Carbillet et al. 2011). We usedtion and Handling software (v0.15.0, Pavlov et al. 2008) andadditional IDL routines for the IFS data reduction (Mesa et al.
2015). The modeling and subtraction of the stellar speckle pattern for both the IRDIS and IFS data set were performed witha smart Principal Component Analysis (sPCA) algorithm basedon Absil et al. (2013) using the same setup as described in
Keppler et al. (2018). Figure 1 shows the high-quality IRDIS
combined K1 K2 image of PDS 70. The outer disk and the planetary companion inside the gap are clearly visible. In addition, there are several disk related features present, which aredescribed in Appendix B. For this image the data were processedwith a classical ADI reduction technique (Marois et al. 2006) tominimize self-subtraction of the disk. The extraction of astrometric and contrast values was performed by injecting negative
point source signals into the raw data (using the unsaturated flux
measurements of PDS 70) which were varied in contrast andposition based on a predefined grid created from a first initial
estimate of the planet’s contrast and position. For every parameter combination of the inserted negative planet the data werereduced with the same sPCA setup (maximum of 20 modes, protection angle of 0.75 ⇥ FWHM) and a 2 value within a segmentof 2 ⇥ FWHM and 4 ⇥ FWHM around the planet’s position
was computed. Following Olofsson et al. (2016), the marginalized posterior probability distributions for each parameter was
A&A proofs: manuscript no. draftcomputed to derive final contrast and astrometric values andtheir corresponding uncertainties (the uncertainties correspondto the 68% confidence interval). For an independent confirmaH2H3 2015-05-03
H2H3 2015-05-31 tion of theK1K2
2016-05-14
L’ 2016-06-01
extracted
astrometry and photometry
we used SpeCal(Galicher et al. 2018) and find the values in good agreement witheach other within 1 uncertainty.

観測から得られた経験式を深く考えずに使うと，原始惑星の場合は質量降着率を桁で過小評価する可能性を見つけた。
【背景】• PDS 70の遷移円盤の穴に，原始惑星の兆候を発見• 赤外線の複数バンドでの点源• Hαの点源• サブミリ波点源

割と良い位置的な一致！（下・右の図）

赤外線のマルチバンド観測（2012-2016）Keppler et al. (2018)
L’ 2012-03-31

c
0.1"

b

2.3. Conversion of the planet contrasts to physical fluxes

The measured contrasts of PDS 70 b from all data sets (SPHERE,
NaCo, and NICI) were converted to physical fluxes following theapproach used in Vigan et al. (2016) and Samland et al. (2017),
who used a synthetic spectrum calibrated by the stellar SED toconvert the measured planet contrasts at specific wavelengths
to physical fluxes. In our case, instead of a synthetic spec-∆DEC (mas)

0
–100
–200

c

0.0

b

b

-0.5

see also
Wagner et al.
(2018)–300–400–400 –300 –200 –100

10
5
0
–5
0.650 0.655 0.660
Wavelength (µm)

0.5

0.0
∆RA (arcsec)

0

100

200

300

400

∆RA (mas)

c

10
5
0
–5
0.650 0.655 0.660
Wavelength (µm)

dSNR2.2. Data reduction

100

SNR

·
−1
IRDIS
dataM
were reduced
as describedin Keppler et al.
±
0.6
yr【結果】PDS 70bについては，log(M) = − 8.0The，70cに
Jup steps consisted(2018). The basic reductionof bad-pixel correction, flat fielding, sky subtraction, distortion correction(Maire et al. 2016), and frame registration.
·−1
The IFS data were reduced with the SPHERE Data Cenついてはlog(M) = − 8.1 ± 0.6 MJup yr でよく説明。また，Tタウリの
ter pipeline (Delorme et al. 2017), which uses the Data Reduc-

200

0.5
∆Dec (arcsec)

accretion モデルを流用することで説明を試みた。

Hα観測:
Haﬀert+(2019)【内容】PDS 70に付随する遷移円盤の穴中に発見された，降着を伴う原始惑星2個と思しきHα輝線の積分強度を，Tタウリ型星のmagnetospheric

Muller et al. (2018)

300

SNR

the IRDIS (Dohlen et al. 2008) dual-band imaging camera(Vigan et al. 2010) with the K1 K2 narrow-band filter pair
( K1 = 2.110 ± 0.102 µm, K2 = 2.251 ± 0.109 µm). A spectrum
covering the spectral range from Y to H band (0.96–1.64 µm,
R = 30) was acquired simultaneously with the IFS integral field
spectrograph (Claudi et al. 2008). We set the integration time
for both detectors to 96 s and acquired a total time on target ofalmost 2.5 h. The total field rotation is 95.7 . During the courseof observation the average coherence time was 7.7 ms and a
Strehl ratio of 73% was measured at 1.6 µm, providing excellent
observing conditions.

400

10
5
0
–5
0.650 0.655 0.660
Wavelength (µm)-0.5

Fig. 1. IRDIS combined K1 K2 image of PDS 70 using classical ADI
reduction technique showing the planet inside the gap of the disk around
PDS 70. The central part of the image is masked out for better display.
North is up, East is to the left.

Long et al. (2018). The spectrum Long et al. (2018). The spectrum covers a wavelength the entire IFS and IRDIS data set. To
obtain flux values for our data sets taken in L0 band at 3.8 µm, we
modeledthe stellarJournal
SED with
simpleblackbodies
to 10
account for
The AstrophysicalLetters,
879:L25(10pp), 2019 July
the observed infrared excess (Hashimoto et al. 2012; Dong et al.
2012). The final SED of the planet is shown in Fig. 2. The IFS
SED of the planet is shown in Fig. 2. The IFS spectrum has asteep slope and displays a few features only, mainly water values are listed in Table C.1.

Fig. 1 | Overview of the PDS 70 system. a, The Hα detection map with anoverlay of the contours of the orbital radii and a white dot in the centre
that marks the position of the star. The contours for PDS 70 c are theminimum and maximum orbital radii found for the different wavelength
observations. For both objects the square apertures that were used forthe photometry are shown, with the red aperture for PDS 70 b and the
purple aperture for PDS 70 c. b–d, The corresponding spectra divided bytheir standard deviation are centred around the Hα line position. The four
apertures in orange (panel a) indicate reference areas that are used tocompare with PDS 70 c. The orange reference spectrum (panel c) does notshow any spectral
while both PDS 70 b and PDS 70 c (panels d andIsella etfeature,al.
b, respectively) clearly show Hα in emission. DEC, declination; RA, rightascension; SNR, signal-to-noise ratio.

P Cygni profile with a blue-shifted emission and a red-shifted
absorption component, which is quite common and a sign of magnetospheric accretion21,22. The Hα line profile from PDS 70 b and3. Results
the other source do not exhibit these features, and show just a single
Gaussian-like profile.
3.1. Atmospheric modeling
The fact that the line shape, velocity offset and the line widths differ from what is expected from reflected starlight excludes reflected
We performed atmospheric simulations for PDS 70 b with the selflight as the origin and indicates that the line emission is locally generconsistent 1D radiative-convective equilibrium tool petitCODE
ated. The multi-epoch astrometry (Table 1) of the new point source(Mollière et al. 2015, 2017), which resulted in three di↵erent
reveals that it orbits in the direction of the Keplerian rotation of thegrids of self-luminous cloudy planetary model atmospheres (seecircumstellar disk and is co-moving with its host star, which makes
Table 1). These grids mainly di↵er in the treatment of clouds:
it improbable to be a background object. We therefore identify thispetitCODE(1) does not consider scattering and includes only
second source of Hα emission as the location of a second, forming
Mg2 SiO4 cloud opacities; petitCODE(2) adds scattering; petitplanet, PDS 70 c. The significant difference between the line-of-sight
CODE(3) contains four more cloud species including iron (Na2 S,
radial velocity of Hα emission from the accreting protoplanets comKCl, Mg2 SiO4 , Fe). Additionally, we also use the publicly availpared with their Keplerian velocities, 4.3 km s− 1 and 3.4 km s− 1 forable cloud-free petitCODE model grid (here called petitCODE(0);PDS 70 b and c, indicates that the dynamics of planetary accretionsee Samland et al. 2017 for a detailed description of this grid) andare different from the orbital dynamics of the planets. This is also thethe public PHOENIX BT-Settl grid (Allard 2014; Bara↵e et al.
case for accreting stars, where the line-of-sight velocity of the accret2015).
ing
material is different from the system velocity.
In order to compare the data to the petitCODE models
The K–L colour and non-detection in the H-band imply that
weFigureuse 1.
the
same
tools
Samlandet towardal. (2017),
ALMA
imagesof theas855describedμm continuuminemission
recordedthe inner regions of the PDS70 circumstellar disk. The image on the left was obtained
redderellipse70
on ctheisbottom
left. than PDS 70 b, which could be either due to aadoptingrobust parameterof 2 andhas anemcee
angular resolution(FWHM of the synthesizedusing
thea python
MCMC
code
(Foreman-Mackey
et al. beam) of 0 093×0 074 indicated by the blackPDS
The image in the middle panel was created adopting a robust parameter of 0.3 and has an angular resolution of 0 067×0 050. White contours are drown at 3 and 5
2013)
N-dimensionalmodel
linearlyinterpolated
attimes onthe rmsnoise level. See the text
for thegrids
details. The
labels PDS70b
smm and PDS70csmm indicate the compact sources of continuum emission discovered inside
NATUR
750 NIR image publishedthedustcavity.
The
image
on
the
right
shows
an
overlay
betweenthe
VLT/Spectro-Polarimetric
High-contrast Exoplanet REsearch (SPHERE)
each evaluation. We assume a Gaussian likelihood function
al. (2018),
and the the
ALMAspectral
low-intensity
red contours fromimage
with robust=2.0.
andin Müller
take etinto
account
correlationof thethe
IFS

サブミリ波連続波
Isella et al. (2019)【モデル】
Muzerolle+ (2001) のTタウリ型星 Magnetospheric Accretion modelを周惑星円盤の推定値(Batygin 2018)でT TAURI MAGNETOSPHERIC ACCRETION. II.

945

当てはめたもの

Basri 1995 ; Johns-Krull & Hatzes 1997), although proÐle
variability in many other cases appears to be more stochastic with a variety of timescales (Johns-Krull & Basri 1997 ;Smith et al. 1999).
There are some remaining problems with the magnetospheric infall interpretation of the line emission. Strong
emission lines have very broad wings, often extending out to^500 km s~1 from line center ; this is especially true at Ha,
the most optically thick optical line (Johns & Basri 1995 ;Johns-Krull & Basri 1997). The previous magnetosphericaccretion models could not account for this high-velocity
emission, since the line widths depended only on the infall
velocity of the gas, which for typical CTTS parameters isaround 200È300 km s~1 (projection e†ects cause the actualline width to be less than this in most cases). They also
predicted larger and more frequent line asymmetries,
because of occultation e†ects, than observations indicate.
Another problem is that the Balmer decrement line ratios ofHa and Hb given by the magnetospheric models were generally much smaller (D1.5È2 ; see MCH) than observed
values (D4È6).
With these challenges in mind, we now critically explore
the limit to which our axisymmetric radiative transfer
models can explain the observed behavior of CTTS line
emission. We have included a treatment of line broadening
in order to test whether this can produce extended wingemission, making line proÐles less asymmetric and raising
the Balmer decrement. We have also added calculations for
sodium as an additional diagnostic with di†ering properties
from the hydrogen lines. The grid of models we present here
fully explores parameter space, covering a wide range ofplausible temperatures (6000È10,000 K), inclinations, thejupiter
entire observed range of accretion rates (10~9È10~6M_yr~1), and four di†erent magnetospheric geometries(varying both width and extent). In some additional cases,
we have investigated the e†ects of rotating magnetospheres
and whether rotation produces signiÐcant e†ects on the lineproÐles. Comparing models with observations, we attempt
to glean more detailed insight into the physics of the accretion Ñow and provide some constraints to help guide moredetailed theoretical models of the interaction between
young stars and their accretion disks.
1

周惑星円盤での概念図: Figure taken from Batygin (2018)

The Astronomical Journal, 155:178 (7pp), 2018 April

B

磁場構造の概念図:

Figure taken fromMuzerolle+ (2001)

FIG. 1.ÈSchematic of the dipole magnetic Ðeld geometry used in themodels (from HHC).

this and an input parameter accretion rate value, we canderive the density distribution. We employ a semiempiricaltemperature distribution, with a shape determined by thedensity distribution (and kept relatively Ðxed regardless ofvariations in other parameters), and an absolute level determined by a parametrized heating rate that is a free parametertheeﬀ(see HHC for further details). Figure 2 shows
jupitervelocity, density, and temperature distributions for a Ðdu-先行研究による 70b, 70cに対する見積もりをカバーするように

on of Planets around PDS 70

Planet Massは 6 M,T3= 1200K, Rp=1.3R

で固定

Table 1. Range of Model Parameters
Parameters

Min.

Max.

Step

-11

-5

1

Tmax
(K)
2.1. Overview

8000

11000

200

R (R )28

2

log(Ṁ/MJup yr

2. MODEL CALCULATIONS)

We build on the models Ðrst presented in HHC andMCH ; see those papers for details of the physical framework and radiative itransfer.pThe following discussion is abrief review of the relevant points retained here. We employ
an axisymmetric, dipolarÐeld geometry for the
r magneticpaccretion Ñow (see Fig. 1). The inner and outer radii of theÑow at the disk are free parameters and set both the size andextent of the Ñow and the boundaries of the accretion shockat the surface of the star. The stellar photosphere and accretion shock radiation are each described by a characteristic
blackbody temperature diluted by the appropriate Ðlling
factors. For all of the results presented here, we use thecanonical T Tauri values of mass and radius (0.5 M , 2 R ).
_ we _The disk is considered opaque at all wavelengths, anddo
not include any source of radiation from it.
Given a speciÐc geometry, the velocity Ðeld of the gas in
the magnetosphere follows, assuming free fall from the disk

W (R )
cos(i)

0.3

1, 2, 4, 6
0.9

Truncation radius
flowの幅0.2

温度は密度によって決まる形式

FIG. 2.ÈVelocity, density, and temperature distributions for a Ðducial
case of mass accretion rate, M0 \ 10~7 M
yr~1, temperature_(characterized by the maximum value of the distribution),T \ 7500 K,
maxR \ 2.2È3
and magnetospheric inner and outer boundaries at the disk,
m with a
R . Dotted lines show the distributions for the same parameters*rotating magnetosphere, V \ 10 km s~1 ; the dashed line in the top panel*
shows the azimuthal componentof the velocity (zero everywhere in the

Figure 2. A newly formed giant planet embedded within the circumstellar disk. Meridional ﬂow of gas within the Hill sphere feeds a circumplanetary de-cretithat connects back onto the protoplanetary nebula. Thermal ionization of alkali metals ensues in the inner regions of the disk, coupling the planetary magnetic
the quasi-Keplerian motion of the gas. Consequently, spin angular momentum is expelled into the circumplanetary disk, resulting in a secular decay of the plrotation rate.

【結果：青がHαフラックス観測値と1σ以内の一致，オレンジが3σ以内の一致】
PDS 70bに対する計算

PDS 70cに対する計算
65

Magnetospheric Accretion of Planets around PDS 70

Thanathibodee et al.

6

Wr [Rp ]

4

0.8

0.8

cos(i)

2

0.6

0.4

11

11

9

の物理量はあまり制約されてい
ない）

0.6

0.4

10

·
M に対してはピークが明瞭（他42Tmax [103 K]

Tmax [103 K]cos(i)Wr [Rp ]

6

10

9

8

log(Ṁ/MJup yr 1 )8log(Ṁ/MJup yr 1 )一番上が，それぞれの物理量で観測値との一致条件を満たした
場合の数

6

8

10
2

4

6

Ri [Rp ]

8

2

4

6

Wr [Rp ]0.4

0.6

cos(i)

0.8

8

9

10
3

Tmax [10 K]

11

10

8

6
1

log(Ṁ/MJup yr )

6

8

10
2

4

6

Ri [Rp ]

8

2

4

6

Wr [Rp ]0.4

0.6

cos(i)

0.8

8

9

10

11

3

Tmax [10 K]

10

8

6
1

log(Ṁ/MJup yr )Figure 2. Same as Fig. 1 but for PDS 70c.
Figure 1. Corner plot showing regions in parameter space that are in agreement with the observed H↵ flux of PDS 70b. Blueand orange points show the models which H↵ line flux are within 1 and 3 of the measured line flux of the planet, respectively.
is calculated from observations of a large number of T
The right panel of Figure 4 shows the relationship beRed crosses show models outside of 3 . The histograms show the distribution of the parameters of the models that fit Tauri
the stars across a wide range of ages using contempotween Ṁ and W10 for all models in our grid. At low
observations, with the same color scheme as in scattered plots.
raneous UV and optical spectra, providing direct meaaccretion rates, Ṁ < 10 8 M
yr 1 , the line width is
Jup

Using simple linear regression, we fitted a powerlaw to the LH↵ model results between Ṁ = 10 10
71able, which yields
log(Ṁ) = (0.280 ± 0.002) log(LH↵ )surements of the accretion luminosity from the UV andof the line luminosity from the optical. Similarly, Rigli(6.14 ± 0.02). aco(2) et al. (2012) used simultaneous observations, but fornot sensitive to Ṁ even if it varies by orders of magnitude. In contrast, W10 increases steeply with Ṁ for Ṁ
& 10 8 M
yr 1 with a scatter significantly increased.

【Hα観測から質量降着率を求める経験式に関する
Tタウリ型星の場合との比較】

Thanathibodee et (2)al. 10%レベルの線幅(W10)(1) LHαとの関係

8

500

0

Model
TTS, Natta+04
PDS70b
PDS70c

今回のモデル計算で，質量降着率ごとでLHαとW10を導出

·
−10
−7
- M = 10 − 10 MJupiter /yr の範囲ではLHαはべき

2
400
4

みにすると大きく間違える可能性がある（同じ質量降着率でも，サイズスケールや密度・温度がTタウリと大きく異なるため？）。

- W10についてもTタウリの経験式(Natta+ 2004)と異なる。密度の違いにより光学的厚みが大きく異なる，大降着率の範囲では圧力幅が効くようになるため（continuum の状況も変わるので，質量降着率を決めるのにW10は役に立たなくなる）

今回のモデル計算で得た関係は，特定の惑星質
量・半径に対するものなので，包括的な議論のためにはさらにモデルの精緻化が必要

Model W10 [km/s]ていたものに比べて steep。Tタウリの経験則を鵜呑

log(LH /L )乗則に従う。過去のTタウリ型星に対して見積もられ
6
8
10
Model
TTS, Ingleby+13
TTS, Rigliaco+12
Power law fit
PDS70b, MUSE
PDS70b, MagAO
PDS70c

12
14
16
11

10

9

8

7

log(Ṁ/MJup yr 1 )

6

300

200
R=3000R=5000100
R=10000
R=20000

5

11

10

9

8

7

6

log(Ṁ/MJup yr 1 )Figure 4. Left: Predicted H↵ luminosity as a function of mass accretion rate for planets with parameters similar to those of
PDS 70b and c. The green and pink shaded regions show 1 and 3 , respectively, around the mean value (blue line) for a given
Ṁ. The line luminosity saturates once the accretion rate reaches the low end of mass accretion rate in pre-main sequence stars(T Tauri stars, TTS), Ṁ ⇠ 10 10 M yr 1 ⇠ 10 7 MJup yr 1 . The dashed purple line and the dash-dotted dark-blue line are
extrapolation of empirical relationships between Ṁ and LH↵ for TTS. The relationships underestimate the Ṁ of the planets by⇠2 orders of magnitude. The solid black line is the power law fit to the relationship for log(Ṁ) between -10 and -7; the dottedline is the extrapolation of the relationship outside the fitted range. Shaded horizontal lines show the line measurements of PDS
70 planets. Right: The W10 of the model H↵ line as a function of log(Ṁ). The purple lines are the relationship for T Tauristars from Natta et al. (2004), with extrapolation to W10 < 200 km s 1 in dotted line. The line width is a very weak functionof Ṁ. Horizontal dotted lines show W10 of Gaussian instrumental profiles for several spectral resolutions. The shaded crimson
dashed line and dash-dotted gray line show the measured W10 of PDS 70b and PDS 70c, respectively.

74: X-shooter spectroscopy of young stars with disks: The
TW Hydrae association as a probe of the final stages of disk
accretion(L. Venuti et al. A&Ap in press)
https://arxiv.org/abs/1909.06699【内容】年齢 8-10Myrの範囲にあると言われているTW Hya association(~50pc)に含まれる赤外超過=円盤を伴う星11個(TW Hyaの円盤付き天体の85%をカバー)をλ=300-2500μmで高分散分光，質量降着有無を調査。
【背景】円盤質量の散逸タイムスケールは研究が進んでいるが，質量降着のタイムスケールは研究が不足。特に中心星質量が小さい領域は未開·【星のパラメータ(M★, R★, L★, age: LaccからMacc出すのに必要)】• 進化モデルによって違いがあるが，本研究ではuniformity を確保するためBaraﬀe et al. (2015)をもとに決定，距離はGAIA2の最新値採用。

• 見積もったM★, R★と，観測で得られたスペクトルからL★を切り分け·たLaccから，Maccは，
LaccR★
·
Macc ∼ 1.25で求めるのが基本線。
GM★
·• 水素原子のライン強度からMaccを求める経験式とも比較する。

拓。TWAは近傍(~50pc)にあり，これらの克服には理想的な領域。
【サンプルの性質：HR図上他での理論モデルとの比較（黒: this study，【生データと大気吸収補正後のスペクトル】

A&A proofs: manuscript no. TWA_XS_accretion
グレー: diskless, 先行研究でターゲットとなっていた天体）】

等質量線

等年齢線

(Baraffe+
2015)
0.01M⦿
0.02M⦿
0.05M⦿
(詳細は
0.1M⦿
脚注参照)
0.2M⦿
0.4M⦿
Fig. 4. HR diagram for stars in Table 2 (black dots) compared with disk–free TWA objects investigated in Manara et al. (2013b) and Stelzer et al.
(2013) (gray diamonds), rescaled to Gaia DR2 distances and Herczeg & Hillenbrand’s (2014) SpT–based T . Filled symbols identify probable
0.6M⦿
TWA members, open symbols targets with disputed membership. Errors on T and log L are as in Fig. 3 (y-axis) and Fig. 2 (x-axis), respectively.

1Myr
2Myr
5Myr
20Myr
100Myre↵e↵‹

Left: Theoretical isochrones from di↵erent model grids (Bara↵e et al. 2015, dotted blue line; Siess et al. 2000, dash-dot red line; Tognelli et al.
2011, dashed green line; Spada et al. 2013, solid black line; Spada et al. 2017, solid yellow line) are shown for comparison on the diagram. The
illustrated isochrone ages are, from top to bottom, 1 Myr, 2 Myr, 5 Myr, 20 Myr, and 100 Myr (the isochrones at 1, 2, and 5 Myr are not available
from the models of Spada et al. 2013, and the 5 Myr isochrone is not available from the models of Spada et al. 2017). Right: Mass tracks from thesame model grids are illustrated for comparison. Displayed masses are, from right to left: 0.01 Md , 0.02 Md , 0.05 Md , 0.1 Md , 0.2 Md , 0.4 Md ,and 0.6 Md (Bara↵e et al. 2015); 0.1 Md , 0.2 Md , 0.4 Md , 0.6 Md (Siess et al. 2000; Spada et al. 2013); 0.2 Md , 0.4 Md , 0.6 Md (Tognelli
et al. 2011; Spada et al. 2017).

(他の詳細は脚注参照)【スペクトルのモデリング：次ページ】

緑：テンプレートの恒星光球からの放射【UV-optical連続波：バルマージャンプ@364.6nmあたりのスペクトル】(左の天体の別データ)

B·
黒：降着するスラブのモデル放射（←ここからMacc を求めるのが王道)
青：トータルフィット
赤：実際の観測されたスペクトルU

B

14天体観測したが，主に扱うのは11天体- 1天体(TWA3B) については過去の研究でdisklessと認定されていたので議論から除く- 2天体(TWA31, TWA22)はGAIAデータからTWSメン

P

P

バーとしての疑念発生，統計的議論からは除く“Bona fide accretors”
アク リ ー シ ョ ン が 明 確 に あ る も の （ バ ルマー
BB

continuumのレベルが光球より高いもの） B
4+(1)天体：TWA1, TWA27, TWA28, (TWA31), TWA32“Possible accretors”
アクリーションがmarginalにあるもの（バルマーP(左の天体の別データ)

continuumのレベルがmarginalに高いもの） P

3+(1)天体：TWA30A, (TWA22), TWA8A, TWA8B
P

B“Upper limit”
アク リ ー シ ョ ン が 明 確 で な い も の （ バ ルマー
continuumのレベルが光球と同等）

U

1天体：TWA3A
その他：“Underluminous”, “No slab model solution)

3天体：TWA30B, TWA40, J1247-3816

【Line放射量(H, He, Ca)との比較】
Continuum でPossible accretorやupper limitだった·【Macc vs. M★L.(1)】
Venuti et al.: X-shooter spectroscopy of young stellar objects: disk accretion in the TWA

Table 3. Mass–dependent values of M9 acc noisstars at an age of 10 Myr using the models of

y of youngstellar objects: disk accretion in the TWA
ものも含め，比較的良い相関が見られる→概ね降着起源

今回のサンプルはより低質量の星まで
M‹
log M9 acc,no

d: i)

rMd s
含んでいる。

model
‹

n the
oints
then

model
e↵

n the
olate
final
lated

le 2.
multiM9 acc ,
(1)

d not
meaects,
d luil insamined
te of
c on
hese
r tar-

y the
used

·
M

11(+2)天体中8(+2)天体で降着の兆候
TWA の円盤天体の~ 72%
円盤付き天体はTWA全体の24%程度なので，降着円盤はTWA 全体の13-17%

Fig. 5. Comparison between the log Lacc values measured from fitting
the excess continuum with a slab emission model (Sect. 3.1.1), andthose derived from the emission line luminosity (Lline ) measurements,using the empirical Lline ´ Lacc calibration relationships from Alcalá
et al. (2014). Error bars on the x-axis correspond to a typical uncertainty of 0.25 dex on the Lacc measurements from the multicomponent
spectral fit; error bars on the y-axis follow from error propagation including uncertainties on the measured line fluxes, and uncertainties on
the calibration relationships, as listed in Alcalá et al. (2014). Green lines
connect multiple measurements for the same object. Light dotted lines
delimit an area of ˘0.3 dex around the equality line (dashed in gray on
the diagram).

low-mass側で少し上にずれているのは※H , He I 5876, H↵, and Pa ), when detected, were converted

to Lacc using therelationships of Alcalá et al. (2014)6 . The fichromosphere活動期限の輝線のせい？
racc (lines) estimate for a given object was then calculatednal L

若い星団(Lupus等)に比べ低い割合。一方で，長期間降着が続く円盤があることも示唆

0.02-12.5380.03-12.2420.04-12.1890.05±-12.061
2.1
0.5
0.06
-11.980acc★0.07-11.886
0.08-11.798
0.09-11.725
0.1-11.693
一部の先行研究で可能性が
0.11-11.643
0.13-11.540示唆されていた
0.15-11.457
Double power-lawの0.17-11.3710.2-11.283
兆候は見られなかった。
0.3-11.0340.4-10.8140.5-10.625
0.6-10.461
0.7-10.291

Fig. 7. M9 acc vs. M‹ distribution for stars in our sample, where M9 acc is
calculated from the measured continuum excess luminosity for all objects (filled dots) except TWA 40 and J1247-3816 (empty dots). Smaller
red dots mark the locations of TWA 31 (upper point) and TWA 22 (lower
point). Green vertical lines connect multiple M9 acc measurements for thesame objects. The shaded gray area traces the estimated chromosphericnoise limit for M9 acc detection, as a function of stellar mass, for stars
aged between 3 and 10 Myr on Bara↵e et al.’s (2015) isochrone grid.
The dash-dot cyan line traces the bimodal M9 acc ´ M‹ relationship derived by Manara et al. (2017b) in the Chamaeleon I star-forming region,
with a break at M‹ „ 0.29 Md . The solid magenta line traces the bimodal relationship derived in Alcalá et al. (2017) in Lupus, assuming a
break at M‹ “ 0.2 Md . Double arrows mark the 10th–90th percentile
range in M9 acc covered by the Lupus (magenta) and Chamaeleon I (cyan)
populations as a function of mass. The dotted blue line traces a linear
fit to the datapoints, which accounts for objects dominated by chromospheric emission (labeled on the diagram) as upper limits.

the HR diagram covered by our targets in Fig. 5. The result of∝M

we consider the corresponding M9 acc as ana true detection for the reason explained
maining sources (8/12, excluding TWA
some accretion activity. The measured M
on the stellar mass. Following Venuti et aa Kendall’s ⌧ test for correlation (Feigels2012) to assess the presence of a correlat
the quantitative relationship taking into a(values consistent with the chromosphe
beled in Fig. 7) as censored data. Assumdescription, we obtained a best-fitting tre
a significance of 2.5 . This result is cons
ships reported in the literature for youn·
【Macc vs. M★ (2): X-shooterで得られ·【Macc の領域間での比較】

no. TWA_XS_accretion
L. Venuti et al.: X-shooter spectroscopy of young A&A
stellarproofs:
objects:manuscript
disk accretionin the TWAていた他領域での結果との比較】

¡7

Cha I
Lupus

¡8

¾ Ori
TWA
´ Cha

¡9

_ acc =M?2
log M

_ acc (M¯=yr)
log M

½ Oph

and a few more from ⇢ Oph and ⌘ Cha, with measured M9 acc
below the median accretion level detected in the corresponding
mass bin for the composite Lupus and Cha I sample (horizontal¡7:5
line in Fig. 11). At lower masses(M‹ „ 0.08 ´ 0.25 Md ), where
the samples overlap, stars in ⇢ Oph appear located systematicallyabove the M9 acc levels measured in the other populations. However, this e↵ect should be interpretedwith caution since the ac¡8:0
cretion rates for ⇢ Oph targets were measured from the emission
lines, contrary to the quasi-totality of other measurements shownon the diagram (which were derived from the Balmer continuum
emission), and are subject to larger¡8:5 uncertainties (Manara et al.
2015). If we exclude ⇢ Oph, on the other hand, it appears that the½ Oph
9 acc at the lowest
remaining populations cover similar ranges
in
M
Lupus
stellar masses, with striking evidence for some 5-10 Myr-old acCha I¡9:0
9creting stars with measured Macc higher than
the typical level of¾ Ori
accretion detected in the combined Lupus + Cha I sample. This
TWA
comparison might be a↵ected by di↵erent
completeness limits´ Chaof the various surveys over the mass
¡9:5 range of intersection among
5:8 surveys
6:0
6:2 not 6:4
6:8
7:0
the various samples (for instance, the Lupus
didin- 6:6
log MaAge [yr]clude many young star candidates with M‹ À 0.1 Md , andnara et al. 2017b quote a completeness limit of SpT „ M6 forFig. 10.theMedianage bias
and would
accretion
properties normalized to squared
the Cha I sample). Nevertheless,selection
impactmass for six young regions targeted with X-shooter: ⇢ Ophiuchi (Man9predominantly the low Macc
at gray
the corresponding‹,ara regimes
et al. 2015;
empty triangle),M
Lupus(Alcalá et al. 2014, 2017;and is unlikely to cause a dearthof diamond),stronger accretors
in themass et al. 2017b; cyan square),magenta
Chamaeleon
I (Manara
range of interest for the twoOrionis
younger
populations.
The
picture(Rigliaco
et al. 2012;
yellow
triangle), TWA (this study; blackin Fig. 11 therefore suggests
low-mass stars
stilldot),thatand very⌘ Chamaeleontis(Rugelthat
et al.
2018; green empty square). The
retain a disk at the age of horizontal
the TWA error
can sustain
a prolonged
bars markthe measuredacage spread among cluster memberspossiblyin the caseof ⇢“stable”
Oph, Lupus,
Cha I, and the uncertainty on thecretion activity with an M9 acc
more(i.e., andslowlyformassive
Ori (Bellet al.
2013)
evolving) in the long-term nominal
than on age
morestars.
This
mayand the associations TWA and⌘ Cha(Bell etphotoevaporation
al. 2015). The verticalerrorsmark the logarithmic disperperhaps be related to a more
efficiente↵ectin
9 acc {M‹2 measured across the various populations. For a clearer
sion
in
M
dissipating material in the inner disks around higher mass stars.
depiction, the datapoints representing Lupus and Cha I were shifted by
To ascertain whether thisconclusion might be a↵ected by-0.015 and 0.015, respectively, from the nominal median age (3 Myr).
the incompleteness of the X-shooter survey of disk–free PMS

¡10
¡11
¡12
¡13
¡1:5

¡1:0

¡0:5

0

log M? (M¯)
Fig. 11. Comparison between the M9 acc ´ M‹ distributions obtained fromX-shooter data for ⇢ Ophiuchi, Lupus, Chamaeleon I , Orionis, TWA,and ⌘ Chamaeleontis (symbols and sources of data as detailed in thecaption of Fig. 10). Upper limits in M9 acc are marked by downward arrows. The horizontal lines mark the median M9 acc measured in bins of
0.25 dex in log M‹ for the combined Cha I and Lupus datasets.

年齢が古いので，降着を示す星の割合は低いものの
average accretion level as a function of ·cluster age as una↵ected
降着を示すものは同程度の"Maccを示すようだ
by the specificmass coverage as possible.
Fig. 10 seems to suggest a marginal decrease in M9 acc withincreasing age, the extent of which is, however, appreciably
smaller than the 1 scatter in value measured in any region.

若干，減少傾向が見えなくはないが，有意とは言えない·【Macc vs. Mdisk】

manuscript no. TWA_XS_accretiontion

M9 acc onto
Mdisk ) is asical pro(see, e.g.,that, in ate of viser the age
esses thathotoevapct) wouldr disk rer disk re† Mdisk {t.
which canmassivean be achis case,(Manaraoter meak to infer
ars in Luompatible(e.g., 古典的な粘性降着)_ acc = 0:1 Myr
Mdisk=M¡7_ acc = 1:0 Myr
Mdisk=M

Inside-outでdisk clearingが進むのか_ acc = 10: Myr
Mdisk=M(e.g., 中心星放射のphoto-evaporation)¡8_ acc [M¯=yr]log Mnd stellar
f M9 acc insive stars
dy of ther correla0.4 Md )pointing
idual age

降着が，quasi-steady なのか

TWA 1

Outside-in でdisk clearingが進むのか¡9¡10(e.g., 外界放射のphoto-evaporation)

TWA 30A
TWA 31の手がかりを与える

TWA 32

¡11
¡12

Lupus
Cha I

J1247 ¡ 3816

¡13
¡6

TWA

¡5

¡4

¡3

¡2

log Mdisk [M¯]Fig. 12. Logarithmic M9 acc vs. disk mass distribution for stars in the
TWA, compared to the populations of the Lupus and Chamaeleon I starforming regions. The straight lines mark the M9 acc ´ Mdisk relationshippredicted for di↵erent stellar ages in a scenario of viscous quasi-steadydisk evolution (Rosotti et al. 2017). Upper limits are indicated by correspondingly oriented arrows.

nal age of the association). TWA 30A and TWA 31 appear even
above the assumed steady-state locus with a viscous timescaleof 0.1 Myr, which crosses the upper envelope of the datapoints
distributions pertaining to Lupus and Cha I. TWA 30A is, however, a nearly edge-on star-disk system, which may have introduced additional uncertainties on the measured parameters (see·Macc ∼ Mdisk /t·
Macc < Mdisk /t·
Macc > Mdisk /t若い星団に対する過去の結果は，·" acc ∼ Mdisk /t を示唆？
M·
TWA31については，"Macc /Mdisk が大きい方にずれているかもTWA 31については，最近のGAIA観測の結果から，TW Associationのメンバーでないかもしれないとの指摘があるが，その場合，より年齢が古くなるセンスなので，quasi-steady
的なpictureとは一層ずれることになる※TWA 30Aについては，edge-onであることが分かっているので，·" acc, Mdiskいずれの見積もりにも注意が必要
M66. Investigating the gas-to-dust ratio in the protoplanetary
disk of HD 142527
K.-L. Soon et al. (PASJ, e-open)
https://doi.org/10.1093/pasj/psz112

used ALMA Data

The Hebig Fe star HD 142527
d = 157±1pc (Gaia collaboration 2016)
Mstar = 2.4M⦿ + 0.13M⦿ @ 0.1” (15.7au)
separation (Arun+ 2019; Biller+ 2012)age = 2.96 Myr (Arun+ 2019)
The lopsided disk
ring-like structure with the inner hole ofr ~ 100au (Fujiwara+ 2006; Fukagawa+2006)
incl. = 27°, Kepler rotation (clockwise)strong dust concentration in northern
regions was revealed based on Band 7
data (Muto+ 2015; Boehler+ 2017)

100

r[AU]
200

300

variation of G/D, dust accumulates in high-p regions
disk is optically thick (τ≳1) in the emissions at λ=0.9mm,
requiring detailed RT calculation→ radial profiles were obtained in just two directions

400

North,gasNorth,dust
surface density [g/cm2]

surface density [g/cm2]

10

0.4
0
0.5
South, gasSouth,dust

0.4
0.3

10-1

10-2

0.1
0
100

200
r[AU]

300

400

This study

Dust continuum &13CO/C18O emissions
at λ=3mm

10-3

0.2

Band 7 (in Cycle 0; Fukagawa+2013, Muto+2015)13CO & C18O (J=3-2); δv=0.12km/s
continuum (336 GHz) with BW=2GHz
beam(FWHM): same as Band 3• smoothed by “imsmooth” in CASANorth, gas
North, dust
South, gasSouth, dust

00.8Band 3 (July 2015, in Cycle 2)
13CO & C18O (J=1-0); δv=0.3km/s &0.12km/s (for deriving peak-TB)
continuum (98.5GHz) with BW=2GHz
beam(FWHM): 0.54” × 0.44” (PA=78°)• Briggs weighting (robust =0.5) + selfcalibration

101

1.6
1.2

ガスとダストの面密度分布，ダストのβを直接導出，G/Dやβの空間分布を議論

Introduction🙆🙅【内容】HD142527円盤に対するALMA B3+B7データを同時に解析し，

10-4

100

200

300

400

r[AU]

Fig. 26. Surface density distribution of gas (solid line) and dust (dashed line) in northern (red) andsouthern (blue) part of the disk. The left panels show surface density in the linear scale (top panel for the

Red: North (PA=21°) ;Blue: South (PA=221°)
Muto et al. (2015)← lower optical depths
allow us to map Σdust &Σgas in the whole disk

Discussions

Results & Analyses(1) Grain size is larger in the
Opacityindex βnorthern
regions
0.6

τdust(ν) = κdust(ν)Σdust,
β
ν
where
κdust(ν) = 10cm2g−1 .
( 1THz )
0.05

0.10

0.15

0.20

0.25

0.30

β

2

13CO (J=3-2)

3

2

1

0
−2
Σ∆R.A.
)g (g cm
(′′ )

0.0

0.2

0.4

−1

0.6

−2
0.8

Σgas

1.2

1.20.8
1.0

6 1.4
1. 1
.2

1..46
1

1.6

~1.7

0.2

0.4

0.8

0.2

0.6

−3

0

−1

2

1

0

−1

−2

∆R.A. (′′ )X(C18O) = 1.79 × 10−8

3

2

1

−1

0

−2

−3

∆R.A. ( )(2) Dust is accumulated
in north(higher-pressure)G/D ≡ Σ /Σ regions
′′

Σgas

−2

3

1.6

−2

1

∆Dec. (′′ )

1.4

1 .0

2

−3

1.
4

−3

(b)

Temperature

0

−1

3

assumption
peak TB=Tdust = Tgas

0.6

1.
2

∆Dec. (′′ )

0.1
0

−3

≲1

0.8

1.6

Σdust

−2

1 .8

1.0

5
0.0

∆Dec. (′′ )250.

β & Σdust

−1

0
1.

0.20

0

1.0

2
1.

1

(a)

0.05

0.6

0.8

2

15
0.

1 .6

0.35

3

1

1 .4

1.2

0.00

1 .2

(c)

Σd (g cm−2 )336GHzpeak-TB + Cont.
(optically thick)

1 .0

3

{

98.5GH

0.8

1.
2

Dust

Iν(dust) = [Bν(ν; T ) − Bν(ν; 2.7K)]
× [1 − exp(−τdust(ν))]

1.2

peak-TB(98GHz) ~8K < peak-TB(336GHz) ~ 24K,
indicating that τ(98.5GHz) < 1

−3

g

10

0

10

d

1

10 2

C18O (J=1-0) Mom. 0

3

NC18O
assumption
Δv to be thermal

G/D

30

3

10

0
10

∆Dec. (′′ )10suggesting that it traces Σgas330

Comparisons with

1

10

3

3

2

3

(cont. subtracted intensity)Iν(C18O) = [Bν(T ) − Bν(2.7K)]
× [1 − exp(−τC18O)] exp(−τdust)
where τC18O(T, NC18O, Δv) .
C18O(J=1-0) well delineates dust continuum,

~3

(c)

3

Gas (C18O)(ISM value; Wilson+ 1999)

30

10

−1
10

J=1-0 is used, except
180°≤P.A.≤240° where J=1-0 is

−2

−3

~20
3

2

1

0
′′

−1

−2

−3

Discussions (continued)(3) Σg ∝ Σd0.47, or G/D ∝ Σd –0.53
along the dust ridge!

"0.47±0.02

=

10

10
0

/D

=

G

/D

10 0

G

/D

=

1

G

Σg (g cm−2 )Σd(a) Σg = (1.27 ± 0.07)
0.10 g cm−2(4) Σdust peak, shifted ahead of Σgas peak ?
→ particle size ~ 5mm (from St≈1 & Σg)

P.A. sector
10 −1
10 −3

(000◦ , 060◦ )(180◦ , 240◦ )(060◦ , 120◦ )(240◦ , 300◦ )(120◦ , 180◦ )(300◦ , 360◦ )

10 −2

10 −1
−2

Σd (g cm)

10 0

direction of rotation (clockwise)
consistent with the theoretical prediction(Baruteau & Zhu 2016) when St ≳ 1

67

The DSHARP Rings: Evidence of Ongoing Planetesimal Formation?
Sebastian M. Stammler1 , Joanna Dra̧żkowska1 , Til Birnstiel1 , Hubert Klahr2 , Cornelis P. Dullemond3,
and Sean M. Andrews4
1

University Observatory, Faculty of Physics, Ludwig-Maximilians-Universität München, Scheinerstraße 1, D-81679
Munich, Germany; 2 Max-Planck-Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany; 3 Zentrum für
Astronomie, Heidelberg University, Albert-Ueberle-Straße 2, D-69120 Heidelberg, Germany; 4 Center for Astrophysics— Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138, USA
E-mail contact: stammler at usm.lmu.de
Recent high-resolution interferometric observations of protoplanetary disks at (sub-)millimeter wavelengths reveal omnipresent substructures, such as rings, spirals, and asymmetries. A detailed investigation of eight rings detected infive disks by the DSHARP survey came to the conclusion that all rings are just marginally optically thick with opticaldepths between 0.2 and 0.5 at a wavelength of 1.25 mm. This surprising result could either be coincidental or indicatethat the optical depth in all of the rings is regulated by the same process.
We investigated if ongoing planetesimal formation could explain the “fine-tuned” optical depths in the DSHARP rings
by removing dust and transforming it into “invisible” planetesimals. We performed a one-dimensional simulation ofdust evolution in the second dust ring of the protoplanetary disk around HD 163296, including radial transport of gas
and dust, dust growth and fragmentation, and planetesimal formation via gravitational collapse of suﬃciently dense
pebble concentrations.
We show that planetesimal formation can naturally explain the observed optical depths if streaming instability regulates the midplane dust-to-gas ratio to unity. Furthermore, our simple monodisperse analytical model supports the
hypothesis that planetesimal formation in dust rings should universally limit their optical depth to the observed range.
Accepted by ApJL

https://arxiv.org/pdf/1909.04674.pdf

http://arxiv.org/pdf/1909.04674• DSHARPで得られた円盤ダスト連続波のλ=1.3mmでの光学的厚みが天体ごとで「揃っている」ように見える点を，streaming instabilityによって赤道面付近でg/d≈1が保たれつつ微惑星形成が起こっている状況下で再現できるとする議論
46

68

The Infall Motion in the Low-Mass Protostellar Binary NGC 1333 IRAS 4A1/4A2
Yu-Nung Su1 , Sheng-Yuan Liu1 , Zhi-Yun Li2 , Chin-Fei Lee1 , Naomi Hirano1 , Shigehisa Takakuwa1,3
and I-Ta Hsieh1
1Institute of Astronomy and Astrophysics, Academia Sinica, 11F of ASMAB, AS/NTU No.1, Sec. 4, Roosevelt
Road, Taipei 10617, Taiwan; 2 Department of Astronomy, University of Virginia, Charlottesville, VA 22904, USA; 3
Department of Physics and Astronomy, Graduate School of Science and Engineering, Kagoshima University, 1-21-35
Korimoto, Kagoshima 890-0065, Japan
E-mail contact: ynsu at asiaa.sinica.edu.tw
We report ALMA observations of NGC 1333 IRAS 4A, a young low-mass protostellar binary, referred as 4A1 and
4A2. With multiple H2 CO transitions and HNC (4–3) observed at a resolution of 0.′′ 25 (∼70 au), we investigate thegas kinematics of 4A1 and 4A2. Our results show that, on the large angular scale (∼10′′ ), 4A1 and 4A2 each display
a well-collimated outflow along the N-S direction, and an S-shaped morphology is discerned in the outflow poweredby 4A2. On the small scale (∼0.′′ 3), 4A1 and 4A2 exhibit distinct spectral features toward the continuum centroid,with 4A1 showing simple symmetric profiles predominantly in absorption and 4A2 demonstrating rather complicatedprofiles in emission as well as in absorption. Based on radiative transfer modeling exercises, we find that the physicalparameters inferred from earlier low-resolution observations cannot be directly extrapolated down to the 4A1 inner
region. Possible reasons for the discrepancies between the observed and modelled profiles are discussed. We constrain
the mass infall rate in 4A1 to be at most around 3 × 10−5 M⊙ yr−1 at the layer of 75 au. For the kinematics of the 4A2
inner envelope, the absorbing dips in the H2 CO spectra are skewed toward the redshifted side and likely signaturesof inward motion. These absorbing dips are relatively narrow. This is, like the case for 4A1, significantly slower than
the anticipated inflow speed. We estimate a mass infall rate of 3.1–6.2 × 10−5 M⊙ yr−1 at the layer of 100 au in 4A2.
Accepted by ApJ• Low-mass protobinary である NGC1333 IRS4A 周囲のガス詳細観測，速度構

http://arxiv.org/pdf/1909.12443
造からインフォールレートなどを議論

https://arxiv.org/pdf/1909.12443.pdf
An ALMA Study of the FU-Ori Type Object
V900 Mon: Implications for the Progenitor
Michihiro Takami1 , Tsu-Sheng Chen1 , Hauyu Baobab Liu1 , Naomi Hirano1 , Ágnes Kóspál2,3 , Péter

http://arxiv.org/pdf/1909.12443

69

An ALMA Study of the FU-Ori Type Object V900 Mon: Implications for the Progenitor
Michihiro Takami1 , Tsu-Sheng Chen1 , Hauyu Baobab Liu1 , Naomi Hirano1 , Ágnes Kóspál2,3 , Péter
Ábrahám2 , Eduard I. Vorobyov4,5 , Fernando Cruz-Sáenz de Miera2 , Timea Csengeri6 , Joel Green7 ,Michiel Hogerheijde8, Tien-Hao Hsieh1 , Jennifer L. Karr1 , Ruobing Dong9,1 , Alfonso Trejo1 and Lei
Chen2
1

Institute of Astronomy and Astrophysics, Academia Sinica, 11F of Astronomy-Mathematics Building, AS/NTU
No.1, Sec. 4, Roosevelt Rd, Taipei 10617, Taiwan, R.O.C.; 2 Konkoly Observatory, Research Centre for Astronomy
and Earth Sciences, Hungarian Academy of Sciences, Konkoly-Thege Miklós út 15-17, 1121 Budapest, Hungary; 3
Max Planck Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany; 4 Research Institute of Physics,
Southern Federal University, Roston-on-Don 344090, Russia; 5 Department of Astrophysics, University of Vienna,
Vienna 1080, Austria; 6 Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, D-53121 Bonn, Germany; 7Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD 21218, USA; 8 Leiden Observatory, Leiden
University & Anton Pannekoek Institute for Astronomy, University of Amsterdam, PO Box 9513, 2300 RA, Leiden,
The Netherlands; 9 Department of Physics & Astronomy, University of Victoria, Victoria, BC, V8P 1A1, Canada
E-mail contact: hiro at asiaa.sinica.edu.tw
We present ALMA observations of 12 CO, 13 CO, and C18 O J=2–1 lines and the 230 GHz continuum for the FU Oritype object (FUor) V900 Mon (d∼1.5 kpc), for which the accretion burst was triggered between 1953 and 2009. We
identified CO emission associated with a molecular bipolar outflow extending up to a ∼104 au scale and a rotating
molecular envelope extending over >104 au. The interaction with the hot energetic FUor wind, which was observed
using optical spectroscopy, appears limited to a region within ∼400 au of the star. The envelope mass and thecollimation of the extended CO outflow suggest that the progenitor of this FUor is a low-mass Class I young stellar
object (YSO). These parameters for V900 Mon, another FUor, and a few FUor-like stars are consistent with theidea that FUor outbursts are associated with normal YSOs. The continuum emission is marginally resolved in ourobservations with a 0”.2×0”.15 (∼300×225 au) beam, and a Gaussian model provides a deconvolved FWHM of ∼90
au. The emission is presumably associated with a dusty circumstellar disk, plus a possible contribution from a windor a wind cavity close to the star. The warm compact nature of the disk continuum emission could be explained withviscous heating of the disk, while gravitational fragmentation in the outer disk and/or a combination of grain growth
and their inward drift may also contribute to its compact nature.
47
Accepted by Astrophysical Journal
https://arxiv.org/pdf/1909.05520.pdf
http://arxiv.org/pdf/1909.05520

Magnetospheric Accretion as a Source of Hα Emission from Proto-planets around PDS
V900 Mon に付随する星周物質に対するALMA観測から周囲のガスの構造・運動を議論し，• FU Ori型のoutburstを示した
70
1
1
2
3
Thanawuth
Thanathibodee,
Nuria
Calvet,JaehanBae,James
Muzerolle, and Ramiro Franco
このprogenitor が Class 4I 的な天体らしいことを議論
Hernándezand the width at the 10% height of the line is insensitive to the accretion rate at M < 10

MJup yr.

Accepted by ApJ

71

http://arxiv.org/pdf/1909.06450

The VLA/ALMA Nascent Disk and Multiplicity (VANDAM) Survey of Orion Protostars I. Identifying and Characterizing the Protostellar Content of the OMC2-FIR4 andOMC2-FIR3 Regions
John Tobin1 , S. Thomas Megeath2 , Merel van ’t Hoﬀ3 , Ana Karla Diaz-Rodriguez4 , Nickalas Reynolds20 ,Mayra Osorio4, Guillem Anglada4 , Elise Furlan5 , Nicole Karnath2 , Stella S. R. Oﬀner6 , Patrick Sheehan1 ,Sarah I. Sadavoy7 , Amelia M. Stutz8 , William J. Fischer9 , Mihkel Kama10 , Magnus Persson11 , James Di
Francesco12 , Leslie W. Looney13 , Dan M. Watson14 , Zhi-Yun Li15 , Ian Stephens7 , Claire J. Chandler1 ,Erin Cox16 , Michanel M. Dunham17 , Kaitlin Kratter18 , Marina Kounkel19 , Brian Mazur2 , Nadia M.
Murillo3, Lisa Patel20 , Laura Perez21 , Dominique Segura-Cox22 , Rajeeb Sharma20 , Lukasz Tychoniec3
and Friedrich Wyrowski23
1

National Radio Astronomy Observatory; 2 University of Toledo; 3 Leiden Observatory; 4 Instituto de Astrofisica
de Andalucia; 5 Infrared Processing and Analysis Center; 6 University of Texas; 7 Harvard-Smithsonian Center for
Astrophysics; 8 University of Concepcion; 9 Space Telescope Science Insititute; 10 Cambridge University; 11 Chalmers
University; 12 NRC-Herzberg; 13 University of Illinois; 14 University of Rochester; 15 University of Virginia; 16 Northwestern University; 17 SUNY-Fredonia; 18 University of Arizona; 19 University of Western Washington; 20 Universityof Oklahoma; 21 University of Chile; 22 Max Planck Institute for Extraterrestrial Physics; 23 Max Planck Institute forRadio Astronomy
E-mail contact: jtobin at nrao.edu
We present ALMA (0.87 mm) and VLA (9 mm) observations toward OMC2-FIR4 and OMC2-FIR3 within the Orion
integral-shaped filament that are thought to be the nearest regions of intermediate mass star formation. We characterize the continuum sources within these regions on ∼40 AU (0.1 arcsec) scales and associated molecular line emission
at a factor of ∼30 better resolution than previous observations
at similar wavelengths. We identify six compact con48
tinuum sources within OMC2-FIR4, four in OMC2-FIR3, and one additional source just outside OMC2-FIR4. This
continuum emission is tracing the inner envelope and/or disk emission on less than 100 AU scales. HOPS-108 is the
only protostar in OMC2-FIR4 that exhibits emission from high-excitation transitions of complex organic molecules(e.g., methanol and other lines) coincident with the continuum emission. HOPS-370 in OMC2-FIR3 with L ∼ 360 L⊙ ,also exhibits emission from high-excitation methanol and other lines. The methanol emission toward these two protostars is indicative of temperatures high enough to thermally evaporate methanol from icy dust grains; overall these
protostars have characteristics similar to hot corinos. We do not identify a clear outflow from HOPS-108 in 12 CO,
but find evidence of interaction between the outflow/jet from HOPS-370 and the OMC2-FIR4 region. The multitude
of observational constraints indicate that HOPS-108 is likely a low to intermediate-mass protostar in its main massaccretion phase and it is the most luminous protostar in OMC2-FIR4. The high resolution data presented here are
essential for disentangling the embedded protostars from their surrounding dusty environments and characterizingthem.
Accepted by ApJ
http://arxiv.org/pdf/1910.00605• 原始星のミリ波・サブミリ波サーベイ（VANDAM）の
結果報告第一弾。個別天体のダスト連続波の構造やア
ウトフロー検出などを議論

https://arxiv.org/pdf/1910.00605.pdf

72

Binary stars in Upper Scorpius
Andrei Tokovinin1 , Cesar Briceño1
1

Cerro Tololo Inter-American Observatory, Casilla 603, La Serena, Chile

E-mail contact: atokovinin at ctio.noao.edu
To address the statistics of binary stars in the 8-Myr old Upper Scorpius star formation region, we conducted speckleinterferometric survey of 614 association members more massive than 0.4 M⊙ (spectral types earlier than M3V) basedon the list of Luhman et al. (2018). We resolved 187 pairs, 55 of those are new discoveries. Using also the publisheddata and the Gaia DR2, a catalog of 250 physical binaries is produced. We carefully estimated detection limits foreach target and studied binary statistics in the separation range from 0.′′ 06 to 20′′ (9 to 2800 au), as well as clustering
at larger separations. The frequency of companions with mass ratios q > 0.3 closer than 50 au reaches 0.2 per decadeof separation (1.4× more than for the field solar-type stars) and does not depend on the mass. At wider separations
from 100 to 104 au, the separation distribution and companion fraction resemble those of solar-type stars in the field.
However, unlike in the field, we found a deficit of binaries with q > 0.7 at separations of ∼500 au, compared to smaller
and larger separations. The distribution of q clearly depends on the separation, with a preference of larger q and alarger fraction of twins with q > 0.95 at smaller separations. The binary population of Upper Scorpius diﬀers frombinaries in the field in several ways, proving that binary statistics is not universal.
Accepted by AJ
http://arxiv.org/pdf/1909.12700• Upper Scoのバイナリー統計探査。Gaiaデータをもとにメンバーを決め，バイナリー探査を行った。
https://arxiv.org/pdf/1909.12700.pdf
An ALMA view of molecularfilaments in the Large Magellanic Cloud II: An early stageof high-mass star formation embedded at colliding clouds in N159W-South

Kazuki Tokuda1,2 , Yasuo Fukui3,4 , Ryohei Harada1 , Kazuya Saigo2 , Kengo Tachihara3 , Kisetsu Tsuge3 ,Tsuyoshi Inoue3 , Kazufumi Torii5 , Atsushi Nishimura1 , Sarolta Zahorecz1,2 , Omnarayani Nayak6 , Margaret Meixner6,7 , Tetsuhiro Minamidani2 , Akiko Kawamura2 , Norikazu Mizuno2,8 , Remy Indebetouw9,10 ,73

An ALMA view of molecular filaments in the Large Magellanic Cloud II: An early stageof high-mass star formation embedded at colliding clouds in N159W-South
Kazuki Tokuda1,2 , Yasuo Fukui3,4 , Ryohei Harada1 , Kazuya Saigo2 , Kengo Tachihara3 , Kisetsu Tsuge3 ,Tsuyoshi Inoue3 , Kazufumi Torii5 , Atsushi Nishimura1 , Sarolta Zahorecz1,2 , Omnarayani Nayak6 , Margaret Meixner6,7 , Tetsuhiro Minamidani2 , Akiko Kawamura2 , Norikazu Mizuno2,8 , Remy Indebetouw9,10 ,Marta Sewi!lo11,12 , Suzanne Madden13 , Maud Galametz13 , Vianney Lebouteiller13 , C.-H. Rosie Chen14
and Toshikazu Onishi1
1

Department of Physical Science, Graduate School of Science, Osaka Prefecture University, 1-1 Gakuen-cho, Naka-ku,
Sakai, Osaka 599-8531, Japan; 2 National Astronomical Observatory of Japan, National Institutes of Natural Science,
2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan; 3 Department of Physics, Nagoya University, Chikusa-ku, Nagoya 4648602, Japan; 4 Institute for Advanced Research, Nagoya University, Furo-cho, Chikusa-ku, Nagoya 464-8601, Japan;
5
Nobeyama Radio Observatory, 462-2 Nobeyama Minamimaki-mura, Minamisaku-gun, Nagano 384-1305, Japan; 6Department of Physics & Astronomy, Johns Hopkins University, 3400 N. Charles Street, Baltimore, MD 21218, USA;
7

Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA; 8 Department of Astronomy, School of Science, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 133-0033, Japan; 9 Department of49
Astronomy, University of Virginia, P.O. Box 400325, Charlottesville,
VA 22904, USA; 10 National Radio Astronomy
Observatory, 520 Edgemont Road, Charlottesville, VA 22903, USA; 11 CRESST II and Exoplanets and Stellar Astrophysics Laboratory, NASA, Goddard Space Flight Center, Greenbelt, MD 20771, USA; 12 Department of Astronomy,
University of Maryland, College Park, MD 20742, USA; 13 AIM, CEA, CNRS, Université Paris-Saclay, Université
Paris Diderot, Sorbonne Paris Cité, F-91191 Gif-sur-Yvette, France; 14 Max Planck Institute for Radio Astronomy,
Auf dem Huegel 69, D-53121 Bonn, Germany
E-mail contact: tokuda at p.s.osakafu-u.ac.jp
We have conducted ALMA CO isotopes and 1.3 mm continuum observations toward filamentary molecular clouds of
the N159W-South region in the Large Magellanic Cloud with an angular resolution of ∼0.′′ 25 (∼0.07 pc). Although
the previous lower resolution (∼1′′ ) ALMA observations revealed that there is a high-mass protostellar object at
an intersection of two linear-shaped filaments in 13 CO with the length scale of ∼10 pc (Fukui et al. 2015), thespatially resolved observations, in particular, toward the highest column density part traced by the 1.3 mm continuum
emission, the N159W-South clump, show complicated hub-filamentary structures. We also discovered that there are
multiple protostellar sources with bipolar outflows along the massive filament. The redshifted/blueshifted componentsof the 13 CO emission around the massive filaments/protostars show complementary distribution to each other, whichis considered to be a possible piece of evidence for a cloud-cloud collision. We propose a new scenario that thesupersonically colliding gas flow triggered the formation of both the massive filament and protostars. This is a
modification of the earlier scenario of cloud-cloud collision which postulated the two filamentary clouds prior to thehigh-mass star formation by Fukui et al. (2015). A recent theoretical study of the shock compression in colliding
molecular flows by Inoue et al. (2018) demonstrates that the formation of filaments with hub-structure is a usualoutcome of the collision, lending support for the present scenario. In the theory the filaments are formed as denseparts in a shock compressed sheet-like layer, which resembles “an umbrella with pokes”.

• ALMAによる，LMCの分子雲フィラメントのAccepted by ApJ
https://arxiv.org/pdf/1811.04400

https://arxiv.org/pdf/1811.04400

X-shooter spectroscopy of young stars with disks. The TW Hydrae association as a

詳細観測

75

A solution to the proplyd lifetime problem
Andrew J. Winter1,2 , Cathie J. Clarke2 , Giovanni P. Rosotti3, Alvaro Hacar3 , Richard Alexander1
1

Department of Physics and Astronomy, University of Leicester, Leicester, LE1 7RH, UK; 2 Institute of Astronomy,
University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK; 3 Leiden Observatory, Leiden University, P.O.
Box 9513, 2300-RA Leiden, The Netherlands
E-mail contact: a.j.winter at le.ac.uk
Protoplanetary discs (PPDs) in the Orion Nebula Cluster (ONC) are irradiated by UV fields from the massive star
θ1 C. This drives thermal winds, inducing mass loss rates of up to Ṁwind ∼ 10−7 M⊙ yr−1 in the ‘proplyds’ (ionised
PPDs) close to the centre. For the mean age of the ONC and reasonable initial PPD masses, such mass loss rates imply
that discs should have been dispersed. However, ∼80% of stars still exhibit a NIR excess, suggesting that significant
circumstellar mass remains. This ‘proplyd lifetime problem’ has persisted since the discovery of photoevaporating
discs in the core of the ONC by O’Dell & Wen (1994). In this work, we demonstrate how an extended period ofstar formation can solve this problem. Coupling N-body calculations and a viscous disc evolution model, we obtainhigh disc fractions at the present day. This is partly due to the migration of older stars outwards, and younger stars
inwards such that the most strongly irradiated PPDs are also the youngest. We show how the disc mass distribution
can be used to test the recent claims in the literature for multiple stellar populations in the ONC. Our model also
explains the recent finding that host mass and PPD mass are only weakly correlated, in contrast with other regions ofsimilar age. We conclude that the status of the ONC as the archetype for understanding the influence of environment
on planet formation is undeserved; the complex star formation history (involving star formation episodes within ∼0.8
Myr of the present day) results in confusing signatures in the PPD population.
Accepted by MNRAS
http://arxiv.org/pdf/1909.04093• OrionでUVフィールドにより比較的大きな質量放出を持つ光蒸発の兆候が見えている原始惑星系円盤のタイムスケール（寿命）が長いように見える点を，cluster中の長い期間でのstar formation historyおよび個々の天体のMolecular cloud distances based on the MWISP CO survey and Gaia DR2

migration を考慮して解決しようとする議論
Qing-Zeng Yan1 , Ji Yang1 , Yan Sun1 , Yang Su1 and Ye Xu1
https://arxiv.org/pdf/1909.04093
1
Purple Mountain Observatory and Key Laboratory of Radio Astronomy, Chinese Academy of Sciences, Nanjing
210034, Peoples Republic of China