SFN #327 (19-24本目)
百瀬宗武(茨城大・理)詳細解説は以下の2本• #21: Episodic accretion constrained by a rich cluster of outflows(T. Nony et al. A&A in press; arXiv:2002.05720)• #22: Episodic accretion constrained GG Tauri A: gas propertiesand dynamics from the cavity to the outer disk (N.T. Phuong et al.
A&A, 635, A12)

rent
assould
07).
with
018)
oto-

Observations were carried out in Cycle 2 between July 2014 andJune 2015 (project #2013.1.01365.S), with ALMA 12 m and 7 m(ACA) arrays and baselines ranging from 7.6 m to 1045 m. W43MM1 was imaged with a 7800 ⇥ 5300 (2.1 pc ⇥ 1.4 pc) mosaic
composed
33 al.
fieldsA&A
with them array and11 fields withT.
Nonyofet
in 12press;
arXiv:2002.05720
ACA. In the 12 m configuration, the primary beam is of 26.700(45.800 with ACA) and the maximum detectable scale is of W43-MM1⇠ 1200【内容】大質量星を多く含むprotocluster.
(⇠ 2100 with ACA). Data were reduced in CASA 4.7, applying
promanual and selfcalibration
scripts. We
ran
the CLEAN
algorithm(d=5.5kpc)
のALMA
maps
in
CO(2-1)&SiO(5-4),
原始星
endwith CASA 5.4 using a robust weight of 0.5 and the multiscalerees(with 0, 1, 3, 9 and 17 beams) option in order to minimize inコア27個が駆動するアウトフローローブ計46個を検出，
lecterferometric artefacts due to missing short spacings. Line datave-アウトフロー諸量や降着時間変動を統計的に議論。
were imaged using the merged (12 m + 7 m) configurations
ent,while the continuum was imaged with the 12 m and 7 m arrayso asseparately. The 12 m-only data provide the best sensitivity image
omchosen to identify compact cores (Motte et al. 2018b) while thes of
7 m only map is the best to probe the column density backgroundaria-【用いたデータ】右図を参照
of these cores. Parameters of the four line spectral windows used
city.
in 視野
this study
are presented
in Table
1. The 1.3
mm12m
continuum•
78”x53”(2.1pc×
1.4pc)
withArray + ACA
rred
00
00
images themselves have 0.44 and 5.3 angular resolutions and
phe1
1
1
rms
noise
levels
of
⇠0.13
mJy
beam
and
⇠9.1
mJy
beam
of, • CO(2-1) & SiO(5-4)はcombine dataによるimage(0.015 K and 0.007 K).
s oflled • Continuum@1.3mmは，コア同定は12m Array only,
PV Table 1. Parameters of the merged data for the spectral windows
backgroundの柱密度導出は
ACA
onlyで
mass used
here.
t al.
iller
Spectral⌫obs
Bandwidth
Resolution
rms
a
falla
window[GHz][MHz][00 ] [km s 1 ]t al.
SiO(5-4)
217.033
234
0.48
0.5
3.3
O13 CS(18-17) 218.125
234
0.46
0.25
3.2
5).
CO(2-1)230.462
469
0.46
1.3
2.3
n are
13
CS(5-4)231.144
469
0.46
0.35
2.9
Maud
(a)
1
1
1
rms
in
unit
of
mJy
beam
.
1
mJy
beam
corresponds to 0.12 K
dful19). at 230 GHz.

21. Episodic accretion constrained bya rich cluster of outflows【Blueshiftとredshift成分のマップ】
T. Nony et al.: Episodic accretion constrained by a rich cluster of outflows

Vsys : 98km/s in LSR
CO(2-1)
Blue: 42-85km/s (LSR)
Red:111-154km/s
SiO(5-4)
Blue: 43-93km/s (LSR)
Red:103-153km/s※ CO(2-1)は自己吸収の影響ある範囲を避けた

Fig. 1. Cluster of molecular outflows in the W43-MM1 massive protocluster. Contours of the integrated blue- and red-shifted emission of the CO(2-1) (in a) and SiO(5-4) (in b) line wings are overlaid on the 1.3 mm continuum emission of the 7 m array (grayscale, with a 5.300 beam and 10 mJy/beam rms). (a) The CO line is integrated over 42 85 km s 1 for blue lobes and 111 154 km s 1
for red lobes. Contours are 15, 35 to 235 by steps of 50 in unit of blue = 33 mJy beam 1 km s 1 and red = 42 mJy beam 1 km s 1 .
(b) The SiO line is integrated over 43 93 km s 1 for blue lobes and 103 153 km s 1 for red lobes. Contours are 8, 18, 28 to 203
in unit of = 23 mJy beam 1 km s 1 . Ellipses in the lower-left corner represent the angular resolutions of each image and a scalebar is shown on the top panel to indicate the size in physical units. Boxes outline the zooms-ins of Figs. 3-4.
2

アウトフロー(lobe)の同定(a)が大スケールのjets(b-d)は拡大図【個々のアウトフローは2成分からなる】• jet-like: 高速(|Δv| > 30km/s)でコリメート- 右図の青と赤のコントア- knots が数珠状，スペクトルでも確認- 途中で折れ曲がったような構造も• outflow: より低速，広がった分布を示す- 右図の薄い緑のコントア• 濃い緑のコントアがダストコア，グレーは12m Arrayで捉えたダスト成分

Fig. 3. Molecular outflows driven by cores in the central region of the W43-MM1 ridge. Outflows detected at high velocity (in a) are
shown in detail with their low- and high-velocity components in zooms b, c, and d. Contours of the integrated blue- and red-shifted
emission of the CO(2-1) line wings are overlaid on the 1.3 mm continuum emission of the 12 m array (gray scale). For blue lobes,
emission is integrated over 42 64 km s 1 (HV, blue contours) and 82 88 km s 1 (LV, cyan contours). For red lobes, emission is
integrated over 108 119 km s 1 (LV, orange contours) and 128 158 km s 1 (HV, red contours). For blue lobes, contours are 7, 15,
30 to 230 by steps of 40 (HV) and 7, 15 to 120 by steps of 15 (LV), in unit of = 20 mJy beam 1 km s 1 . For red lobes, contours
are 10, 20 to 160 by steps of 20, in unit of LV, R = 37 mJy beam 1 km s 1 (LV) and 7, 15, 30 to 280 by steps of 50, in unit of(HV). Green ellipses locate the W43-MM1 cores, arrows (sometimes broken) indicate the direction of their outflows, dashed line
outline the best developed outflow cavities. Knots of the blue lobe of core #8 are labelled B1 to B7 in zoom b (the first knot B0 is
not visible on this integrated map). Ellipses in the lower corners represent the angular resolutions of the images and a scale bar isshown in a and b to indicate the size in physical units (same physical units in c-d and Fig. 4 than in b). Some isolated features, likethose outlined by the blue contours below the Z2 box in a, correspond to di↵use gas not associated with the W43-MM1 outflows.

beam, by:H ,f2
more general trends is howeverdifficultbecauseobs⌦beamµ mourH 1.3m
limitations (e↵ect of self-absorption and missing short s
23 developi
2preclude the detection
cavities
I7mof the outflow
=
9.4
10
cm
NH2 ,f to=the systemic velocity (| V| < 10 km s 1 ).

⌦beam µ mH 1.3mm B1.3mm (T dust )!
(1)
I7mis the23
2 ⌦最高速度
where
beam solid
beam
3.4.
Cloud
environment
crossedby, outflow
lobes an
= 9.4 10 cm
Associated coreBlue lobe
Red lobe
Jy weight,
beam 1 mH is the mass
molecular
Name
RAa
Dec.a
Mcore aVLSR b
V max cLmax cPAd
AV e
V max cLmax cPAd
AV e
Comment f
We characterized the cloud density within which each
[J2000][J2000][M ][km s 1 ] [km s 1 ] [⇥10 2 pc][deg][⇥103 mag] [km s 1 ] [⇥10 2 pc][deg][⇥103 mag]is bytheestimating
dust opacity
percolumn
unit masscolo
コアからの最大距離
lobe
develops
themean
density,#1
18:47:47.02-1:54:26.86
102 ± 598
34
5.3
121.5 ± 0.5
2.5
where
⌦beam is the
beam
solid
angle,=
2.8
theB(T dust ) µ
isofthefunction
1.3mm#218:47:46.84-1:54:29.30
55 ± 599
57
5.0
74.5 ± 0.8
1.7
48
6.1
114.6 ± 0.9
2.4
alently
the visualextinction,
the Planck
cloudisoverthemeancom⇤#3
18:47:46.37-1:54:33.41
59 ± 498
83 -9111115 ± 0.5
0.78 (0.53)
60
13
105.5 ± 0.9
0.52 (0.28)weight,
molecularmHof isthemass lobe.
ofatomichydrogen,1.3mmBecause
W43-MM1
is acolumn
warm
and
tent
each
outflow
The (filtered)
densi
⇤
⇤
放出方向(inPA)#4
18:47:46.98-1:54:26.42
36 ± 3102
19 -25
6.9 -9.6
50 ± 1
2.4
Conf
anyof sight
mayderivedfrom, thecm
1.3g
is
the
dust
opacitypergiven
unitline
mass
columndensity
at
1.3
mm,
and
7m
T dust=23
Kbeand1.3mm
= I0.01⇤⇤#7
18:47:47.26-1:54:29.70
23 ± 299
50
8.3
25.1 ± 0.5
1.2
43
18
154 ± 10.39 (0.25)
densitymeasuredinthe
7 am-array
continuum
image
in.
#8
18:47:46.54-1:54:23.15
14 ± 2
97
82⇤ -904418⇤ ± 30.22 (0.13)10113171⇤ ± 2B1.3mm
1.4(T
(2.5)
)
is
the
Planckfunctionfor
dust
temperatureTetal.
2018b).
Since,
duetointerferodust
dustbeam, by: 埋れ度合い#918:47:46.48-1:54:32.54
17.8 ± 0.9
97
61
9.6
56⇤ ± 10.83 (0.68)
64
7.2
127.0⇤ ± 0, 5
0.99
00dense cloud,(inwe
Because
W43-MM1
is
a
warm
and
weAv)
assumedthan⇠21
are missing,
added
aA
#10
18:47:46.91-1:54:29.99
16 ± 1101
29
2.7
66.2 ± 0.5
1.7
Conf
#11
18:47:46.52-1:54:24.26
2.1 ± 0.3
95
20
4.8
25.9 ± 0.5
0.5552⇤ -65
1.9⇤ -6.1
173 ± 4 T dust =0.63
CO, Conf(R)
23 K and1.3mm = 0.01
g 1 (seearguments
in Motte
at thecm
outskirtof theW43-MM1
clou
I7m
#12
18:47:46.57-1:54:32.04
31 ± 498
42
2.9
157.5 ± 0.5
1.2
N00 larger
H2 ,f =
etal.
2018b).
Since,
duetointerferometer
filtering,
scalesumn
density
image,
with25
resolut
⌦
µ
m

B
#13
18:47:46.92-1:54:28.62
11 ± 39927
2.9
131 ± 42.6
18
2.952 ± 3
2.0
Conf
beam
H 1.3mm 1.3mm (T dust )00!
than ⇠21
are missing,
we addedmag levelmeasured#14
18:47:46.97-1:54:29.66
19 ± 4100
45
4.2
165.8 ± 0.5
1.8
Conf(2013).a Av=130
I7m
23
2
#16
18:47:47.02-1:54:30.78
36 ± 610060
3.7
110.2 ± 0.5
1.8
62
4.8
33 ± 1
1.5= 9.4 10cm
, folloattheoutskirt
oftheW43-MM1
cloud
ridge
on
the
Herschelcol1
Equation
1
then
leads
to
the
Jy
beam
#18
18:47:46.25-1:54:33.41
28 ± 496
25
2.9
24.7 ± 0.5
0.69
00
image, with 25extinction,
resolution,#19
18:47:46.88-1:54:25.74
9±3
100
27
2.1
107 ± 41.530
2.968 ± 1 umn density
2.0
AVby: Nguyen Luong et al.
#22-a 18:47:47.05-1:54:32.15
8±2
100
45
6.1
35.4 ± 0.5
0.91
46
7.2
159 ± 2 (2013).
0.83
where ⌦beam is the beam solid !angle, µ = 2.8 is th
#22-b
58
4.0
127.5 ± 0.50.70
molecular
weight,mH isNthe
mass of atomic
hydrogen#23
18:47:46.90-1:54:24.30
11 ± 3101
39
5.1
22 ± 11.1
COEquation
1
then
leads
to
the
followingforthevisualH2 equation
AV = per unit
⇡ 130density
mag +is the dust opacity
mass2 column
at 9401.3#31
18:47:46.73-1:54:17.53
4.3 ± 0.7
96⇤11⇤ -147.7
17.2⇤ ± 1
0.31
16⇤ -205.9
169.3⇤ ± 1
0.36
CO(R)
21
extinction,A:
V Conf(R)
10 cm#36
18:47:46.64-1:54:19.51
3.1 ± 0.5
96
11⇤ -148.5
165 ± 20.45
27
6.3
14.6 ± 0.6
0.32
B
(T
)
is
the
Planckfunction for a dust temperatu
1.3mm
dust
⇤
⇤
#49
18:47:46.59-1:54:20.48
3.7 ± 0.8
9649
6.1
120.1 ± 0.5
0.85
50
5.9
44.8 ± 0.5
0.33
!
! we
Because
W43-MM1 is a warm and dense cloud,#15
18:47:44.77-1:54:45.22
3.9 ± 0.3957813152 ± 20.31
68
18
24.6 ± 0.5
0.17 (0.14)NH
I17m
2
A
values
span
more
than
one(2)
or
T
=
23
K
and

=
0.01
cm
g
(see
arguments
⇤
dust
1.3mm
V
#26
18:47:44.94-1:54:42.84
5.2 ± 0.7
954518
131.1 ± 0.5
0.24
71
12
48.5 ± 0.5AV =0.31 (0.23)
⇡
130
mag
+
940
mag
⇥
.
1Sincescal
Since,
due
to interferometer
filtering,#29
18:47:45.22-1:54:38.81
3.1 ± 0.59539
9.3
78 ± 1
0.24-1021 cm 2 et al. 2018b).
to
2600
mag
(see
Table
2).
the
Jy
beam
00
⇤
#39
18:47:44.61-1:54:42.16
2.3 ± 0.4
95
23
19
30.1 ± 0.9
0.21
24
19
133 ± 4
0.28
than ⇠21 are
missing,
we addeda Av=130mag leveltemmshouldhaveratherhomogeneous#44
18:47:44.76-1:54:46.76
2.5 ± 0.3
9750
9.3
36.2 ± 0.5
0.19
at the outskirt of the W43-MM1 cloud ridge on the Hers⇤A
values
span
more
than
one
order
of
from 140
weestimatetherelative by
uncertaintyo#59
18:47:44.77-1:54:44.17
1.4 ± 0.5
95
62
7.5
24.9 ± 0.5
0.32
64
6.7
150.8 ± 0.5
V0.31
00 magnitude,
umn
density
image,
with25
resolution,
NguyenLuo⇤#67
18:47:44.09-1:54:48.89
2.2 ± 0.3
9574
2283.8 ± 0.5
0.14
67
19
101.5 ± 1.5to 2600
0.16mag (see Table 2). Since the median core environments
30% or less. Similarly,
the large-sc
vs.
(2013).
(a)
Values taken from Motte et al. (2018b).
(b)
shouldhaveratherhomogeneoustemperatureand dust
uncertain
byopacity,
upfor
to t
Equation
1 thenAleads
to the following
equation
V values
Velocity at rest estimated with the 13 CS(5-4) line, except for cores #1, #2, and #4, for which the O13 CS(18-17) line was used instead. A star marker indicatesthe velocityof theneighboringcore havethat is used for cores neither detected in 13 CS(5-4) nor in O13 CS(18-17).
extinction,AAs
we estimate the relative
uncertaintyof ourAV values
to of
be theabout
absolute
valuesv
V : for the(c)
100
Maximum velocity o↵set and length, uncertain by 2 km s and 0.2 (0.005 pc), respectively. A star marker indicates lower limits used in our calculations when measurements are confused by!
30% or less. Similarly, the belarge-scalefilteredcloud
structuresuncertain
by
up
to
a
factoroftwoline absorption features or neighboring outflows.
NH2コアのパラメータ
I7m(d)
Position angle of the molecular jet, measured from West to North. A star marker indicates molecular jets that are o↵set by more than 0.200 from their core havecenter. AV values uncertain
by
up
to
30%,
A
=
130
±
40
mag.1
AV =
⇡ 130badly
mag +constrained.
generally
V940 mag ⇥21
2(e)
Jy beam
Visual extinction measured over the complete extent of the outflow lobe, uncertain by 30%. A second measure at the tip of the molecular jet is given in parenthesis for lobes covered by more 10
than cmAsfortheabsolute
valuesofthevisualextinction,
they
could
2
3 pixels in the 7 m array continuum map (i.e., with Lmax 9.6 ⇥ 10 pc) and when it di↵ers by more than 20%.
(f)
Less reliable outflow lobes are pinpointed by mentioning their confused environments (denoted with ‘Conf’) or their non-detection in SiO(5-4) but in CO(2-1)(denoted withby‘CO’).
be uncertainup‘(R)’
toAVameans
factorspanof twoopacities
which aref
valuesmoreduethantoone
order of magnitude,that the note applies to the red lobe only.
jets the median core envir
generally badly constrained.
to 2600 mag3.5.
(see Knots
Table 2).inSinceshould have rather homogeneous temperature and dustof themolecular
jetsAobserved
we estimate Most
the relative
uncertaintyof our
V values tointegrated
maps consistof clouda conts
the large-scalefiltered
3.5. Knots in jets30% or less.4)Similarly,
Table 2. Main characteristics of each outflow lobe driven by W43-MM1 protostars located in its center (top rows) and south-western (bottom rows) part.

T. Nony et al.: Episodic accretion constrained by a rich cluster of outflows

T. Nony et al.: Episodic accretion constrained by a rich cluster of outflows

PV図上でのKnotsの同定,
速度コヒーレンスの確認• 46個のlobesうち，38のlobes(22個のコア)で86個のknotsを発見- 複数のknotsがあるものはepisodicな降着率の上昇を反映？
【目立つ特徴】• Hubble law的な構造（見かけの加速）- 右図の黒で示されているもの• ゼロ加速，減速しているような構造も- (a)のR2やB1, (b)のB2など

Fig. 5. PV diagrams along the molecular outflows of cores #8 (in a), #3 (in b), and #7 (in c). Position o↵sets, r on the X axis, aregiven with respect to the core center or the connecting point between the red and blue lobes. Negative and positive position o↵sets
correspond to the red and blue lobes, respectively. Each figure has the same physical scales on the X axis and on the Y axis. Contours
are 5, 10, 22 to 67 by 15 steps in unit of CO = 2.5 mJy beam 1 . Additional 16 CO contours are drawn in a and c. Vertical andhorizontal segments in red locate the detected knots (at rknot in Table 3) and outline their associated velocity range from which Vknot
was measured. Segments with arrows in a and b indicate the knots whose velocity range is limited by absorption features. Fingers,
with Hubble-law gas distributions connecting the knots with lower-velocity structures, are indicated with continuous black lines,
gas layers possibly associated with lateral forward shocks are indicated with dotted red lines. The PV area confused by anotheroutflow is hatched in b.
9

議論1. Outflows characteristics in a dynamical environment(1) コアからの最大距離 vs. (a)コア質量，(b)埋れ度合い
T. Nony et al.: Episodic accretion constrained by a rich cluster of outflows(2) CO(2-1)とSiO(5-4)との比較• SiOの高検出率(42/46=91%); 大質量星形成領域に特徴的• この領域の特徴：morphologyも類似。気相中のSiが例外的に高い環境にあるため？
T. Nony et al.: Episodic accretion constrained by a rich cluster of outflows

Fig. 8. Maximal outflow length observed projected on the plane of the sky, Lmax , vs. the mass of the launching core, Mcore , (in a)and the visual extinction of the cloud background crossed by the outflow, AV , (in b). For the longest outflows, measurements attheir tip (in parenthesis in Table 2) are used instead of that averaged over their complete extent. Filled and empty squares pinpoint
outflow lobes developing in the central and south-western parts of W43-MM1, respectively. Blue and red colors show the blue- andred-shifted lobes, respectively. The anti-correlation found in b (best fit line Lmax / A↵V with ↵ = 0.5 ± 0.1) suggests that molecular
outflows are shorter in denser environments.

(b)はAvと逆相関＝深く埋れたアウトフローほど短い• 周囲の分子雲成分との相互作用で邪魔されるwith the outflow momentum flux (e.g., Bontemps et al. 1996;
Duarte-Cabral et al. 2013; Maud et al. 2015).

down or even disperse the gas contained within protostellar outflows. This global infall is probably neither spherical nor continuous in time but it would rather consist of sporadic accretion flows along filaments thus following privileged directions(e.g,. Schneider et al. 2010; Peretto et al. 2013; Motte et al.
2018a). The most massive protostars #1 and #4 are the best sites
where this scenario could apply since they are located at the center of gravity of the infalling protocluster (Herpin et al. 2012;
Cortes et al. 2010) and they only developed small, low-velocity
monopolar (red) outflows (see Fig. 3c). Interestingly, both their
missing lobes and the largest gas inflow observed toward theclump hosting cores #1 and #4 are from the blue-shifted part ofthe spectrum (Pouteau et al. in prep.). Therefore one can expect
this inflow of material to partly quench and disperse the outflowing gas, thus preventing the development of well-organized• 個別のアウトフローのB/R非対称もよく説明（前ページの#7, 8はcluster領域から飛び出る側のLobeが長Regardless the correlations or non-correlations discussed
above, they are all obviously limited by our ability to detect protostellar outflows in the central part of W43-MM1. This ridge
contains extremely-dense, 104 107 cm 3 (see Sect. 3.4, Nguyen
Luong et al. 2013; Louvet et al. 2014, 2016), material that shouldslow down the propagation of molecular outflows. While noglobal slowing down is observed in Fig. A3b, the inspection of
the PV diagrams of the asymmetric outflows of cores #8 and #7
possibly give some insights (see Figs. 5a and 5c). For core #8,
the velocity of knots, when they remain observable, are not much
impacted, with the velocity o↵sets, Vknot , of R3 and R4 similar to that of R1 (see Fig. 5a) and to the mean velocity of theい；3ページ前のマップも参照）✴ Wiggling pattern: binaryによるjet precession ?

Fig. 9. Molecular outflows observed in SiO(5-4) compared with their CO(2-1) emission in the central region (in a, Z1 field o1ferred to as ‘Hubble laws’. We here interpret them as the sum of
1/ molecular gas ejected by the protostar and ballistically propagating plus 2/ ambient gas dragged and accelerated by the jet(see also Arce et al. 2007; Rohde et al. 2019). In our case, fingers
with Hubble-law gas distributions represent traces of enhanced
ejection events that add to a more continuous average ejectionlevel. At the tip of these fingers observed in PV diagrams (seeparticularly Figs. 5a-c and Figs. 6b-c, 6f), vertical segments appear that correspond to knots. These knots, also called internalmark enhanced velocity and density spots of the forward
T. Nony et al.: Episodic accretion constrained by a richshocks,
cluster of outflows
shock where the ejecta and/or its accompanying accelerated gas
catches up with slower material from previous ejections (e.g.,
Jhan & Lee 2016). In some PV diagrams, a gas layer/structure
presents a vertical or strongly declining slope that connects the
knot to the VLSR , recalling the features seen in simulations (e.g.,
Lee et al. 2001). This gas layer could correspond to the lateralforward shocks.
In our general case, fingers with increasing velocities, knots,
and layers with decreasing velocities create triangular-shapestructures in the PV diagrams (see Figs. 5a, 5c, Figs. 6a-b and6e-f). Several of these triangular shape structures add up forming
nested shocks that appear like a continuous emission plateau in
the PV diagrams (see, e.g., Figs. 5a, 5c). At large o↵sets, most of
the low- to intermediate-velocity (<30 km s 1 ) emission disappears from PV diagrams, leaving the high-velocity vertical segFig. 10. Outflow dynamical timescale corrected for a mean projection angle of 57.3 , tdyn, corr , vs. the mass of the launching core,ments Arelatively
isolated (case of B5 to B7 in Fig. 5a). For a
Mcore , (in a) and the visual extinction of the cloud background crossed by the outflow,
V , (in b). The dynamical timescale of agiven core is evaluated in two ways: as the oldest ejection event (triangles) and ascouple
the Lmaxof/ Vknots,(squares).
Both dynamical
highestvelocity appears upstream, drawmax ratio theirtimescales are plotted in a, showing that they generally are consistent within errorbars.
Asingle estimatefast-tail’is used in b,that correinga‘slow-head,profile(see B2 and B3 of core #67
sponds to the oldest ejection event (triangles) when it can be measured and the Lmax / V max ratio otherwise (squares). No significant
6f).
Santiago-Garcia
al. and(2009)and more recently Jhan
correlation is observed in a. The anti-correlation between tdyn and AV in b, tdyn in/ Fig.
A↵V with↵ =0.7 ± 0.1 (magentaetline)adyn
Pearson correlation coefficient of -0.67, recalls that between Lmax and AV in Fig. 8b.
& Lee (2016) attributed these ‘inverse-Hubble laws’ to projection e↵ects of internal shocks, in agreement with the pulsed jetof the beam (⇠0.200 or ⇠0.005 pc). We consider that this sec- have started insimulations
the south-westernpart of W43-MM1
first.(1993).
Thisof Stone& Normanond method provides the most reliable estimates of the ejection fits in the frameworkof recentcloudPV
formationmodelsobserved(e.g., Lee for outflow lobes of W43Mostofthediagrams
timescale. The two timescale estimates given above are consis- & Hennebelle 2016; Vázquez-Semadeni et al. 2019), where sevprotostars
tent to within the uncertainties for 19 out of the 23 cores or mul- eral cloud andMM1
star formation
eventsdisplay
sum up toseveral
form theknots
final (see, e.g., Figs. 5a-c andstars. This
That said,
the anti-correlation
observed in of episodic ejection, with
tiples (see Fig. 10a). The dynamical timescales of outflows in population of 6f).
probablyis the signature
W43-MM1, before any correction of projection e↵ects, range Fig. 10b recalls the one found in Fig. 8b between Lmax and AV .
variation of the velocity and/or of the mass of the gas ejected,
from 400 yr to 104 yr, with a median value of ⇠1500 yr.
We recall that outflows in the central part of W43-MM1 generepisodic
massive could
protostarsbe
andrelated
develop intodense
environ- accretion. When observed
To estimate the real values of dynamical timescales, one ally arise fromwhichal. 2018b,
andspatial
Table 2).and/or
Therefore,
the trend resolutions
obwithlower
spectralor lower sensitivneeds to correct from projection e↵ects, process which is haz- ments (Motte et
be completely environmental (see Sect. 4.1).
ardous when the inclination angle remains unknown (see how- served here could
ity,
only the high-velocity part of the PV diagram stands out.
ever Sect. 4.2.3). Measured velocities and lengths in fact are ve- In any case we stress that outflow dynamical timescales should
Theseas velocity
peaks,
corresponding
rough estimatesof theprotostellar life-to internal shocks, are thenlocities projected along the line of sight, Vmax = Vreal ⇥cos (i), only be consideredreasonis thatstructures
outflows verythat
likelycan
get missedor be identified. When sevonlyclearly
where i is the angle to the line of sight, and lengths projected times. A first thecloud, if they do not simply fall below our detecon the plane of the sky, Lmax = Lreal ⇥ sin (i). The dynami- dispersed in theeralvelocity peaks appear, PV diagrams present a jagged procal timescales corrected for the projection angle would then be tion limit. A second important reason is that our angular resolufile sometimes
wedge’
to measure the calledvelocity ‘Hubbleof the atomic
and/or (Arce & Goodman 2001;
tdyn,corr = tdyn / tan (i). Assuming a random distribution of incli- tion does not allow
R ⇡/2
molecular jet Stojimirović
directly ejected byetthe
onlyettheal.
ve-2013; Plunkett et al. 2015),
al.protostar,
2006; butArcenation angles, the mean value is given by i = 0 i sin (i) di = locity of the entrained
molecular layer that could be a factor of athatshouldresemble to a smoothed version of, e.g., Fig. 5a.
1 rad = 57.3 . This leads to dynamical timescales that, once few to ten times
slower.

Our t sample should be complete as long as the referred two
knots can be separated in the PV diagrams. This will be thecase when their relative distance exceeds the spatial resolutionof the CO cube, FWHM = 0.013 pc, and their velocity o↵set the first quartile from which t values are measured in our
sample, Vlobe ⇠ 33 km s 1 (see Sect. 3.3 and Table 3). Whencorrected for an average projection angle of 57.3 , the mediantime span between ejection events and the 75% completenesslimit of our sample are ⇠500 yr and ⇠250 yr. The histogram ofFig. 11 clearly peaks at ⇠500 yr, well above the completenesslevel, suggesting that a typical timescale between two ejecta exists in W43-MM1.

議論2. Outflows characteristics in a dynamical environment(1) Dynamical Timescale vs. (a)コア質量，(b)埋れ度合い(2) 複数Knots間隔を時間間隔に換算• ~500yr にピーク，episodic accretionの間隔?

(b)はAvと逆相関＝深く埋れたアウトフローほどt• Environmental ?
• 深く埋れたYSOs (outflows)ほど若い？

小Fig. 11. Distribution of the timescale di↵erences between successive ejection events. The top and bottom axes show thetimescales before and after the correction by an homogeneous
inclination angle of 57.3 , tcorr and t respectively. The 75%
completeness limit, ⇠250 yr after deprojection, is represented
with a dashed line. Outflows display a variability with a typicaltimescale, corrected for projection e↵ects, of tcorr ⇠ 500 yr.

We evaluated the projection e↵ects of t measurements by• 次ページ：knots間の平均間隔と視線速度の関係からも
considering two extreme configurations presented in Figs. 12a-statistically corrected for homogeneous projection e↵ects, range
from tdyn,corr = 280 yr to 6300 yr.
Figures 10a–b display the dynamical timescale of protostel-b. First, we assumed that ejections in W43-MM1 have the same
periodicity throughout the whole cluster and that the observed
dispersion is entirely due to projection e↵ects. We recall thatmeasured distances and velocities are projected on the plane of同様の時間間隔が示唆

4.2.2. PV diagrams as open books on ejection variability

T. Nony et al.: Episodic accretion constrained by a rich cluster of outflows

lobe

Fig. 12. Mean projected distance between knots, r , vs. the mean velocity of knots, Vlobe , for 21 outflow lobes with at least
2 knots. Error bars represent the ±1 dispersion. Measurements and core numbers are indicated in red for outflow lobes discussed
below. In a, blue lines locate the measurements for the mean inclination angle of a random distribution of outflows (57.3 ) and forthe inclination angles of the first and last quartiles (41 , 76 ), assuming a constant variability timescale of tcorr = 500 yr. In b, bluelines locate the measurements for the median timescale di↵erence between knots, tcorr = 500 yr, and the first and last quartiles( tcorr = 400 and 800 yr), assuming a constant inclination angle of 57.3 .
median velocity ofVlobe = 45 km s 1 . A configuration withand western lobes (see Fig. 4a). The western jet and the western

The imaging and deconvolution was made with naturalweighting, and the images were cleaned down to the rms noiselevel with the hogbom algorithm Channel maps of the observedlines are presented in Figs. A.1 and A.2. In these data, themaximum recoverable scale is about 700 (1000 au at the source
distance). The images were not corrected for primary beam
attenuation: the beam size of 1700 yields attenuation of 12% at
500 au and 4% at 300 au.

22. Episodic accretion constrained GG
Tauri A: gas properties and dynamics
from the cavity to the outer disk
N.T. Phuong et al.
A&A, 635, A12
3. Results
3.1. Continuum emission

Figure 1 shows the continuum emission at 330 GHz derived【内容】GG Tauを囲むRingのALMA
COデータを解析。
from the Cycle 3 data. It reveals emission from the Aa disk andthe ring structure detected in previous observations (see Dutrey
et al. 2016, and references therein), but the ring is now clearly
resolved. The ring is not azimuthally symmetric (after taking thelimb-brightening effect along the major axis into account), butdisplays a ⇠15–20% stronger emission at PA ⇡ 240 260 . The
outer edge is clearly shallower than the steep inner edge, confirming that some dust remains beyond the ⇠260 au outer edgeof the ring, as initially mentioned by Dutrey et al. (1994). As inprevious studies (e.g., Guilloteau et al. 1999; Piétu et al. 2011),
compact, unresolved emission is detected in the direction of thesingle star Aa, but no emission originates from the Ab close
binary system.
A detailed study of the continuum emission is beyond thescope of this paper and is deferred to subsequent work.

一捻りした解析でリング内ガスの性質や運動を導出。

Kepler速度の約10%のinfall成分を検出，inner diskへの質量供給が星へのそれと同程度であることを示唆。

【GG Tau in Continuum】右図を参照• GG Tau Aa & Abからなる連星系(さらに離れたとこ
3.2. Images of line emission
ろにGG Tau B)

Figure 2 shows the integrated intensity and the velocity mapsof 13 CO(3–2) (left) and C18 O(3–2) (right). In these figures,the continuum has been subtracted. The velocity fields suggest
Keplerian rotation insideTable
the disk.
1. Observations.
13
The CO(3–2) emission extends out to 550 au, and the
C18 O(3–2) is mostly visible in the dense ring.
Spectral line Frequency Energy level
(GHz)
(K)
3.2.1. Intensity variations
12
CO(6–5)691.47333.2
Figure 3 (upper panels) shows
the radial345.795
profiles of theinte12CO(3–2)
16.6
grated brightness of the lines
13 and of the peak brightness of theCO(3–2)330.588
15.9
continuum emission, averaged
overtheentireazimuthaldirecC18 O(3–2)
329.330
15.8
tion, after deprojection to the disk plane. The deprojection wasContinuum
330.15–made assuming a position angle of the minor disk axis of 7 andan inclination of 35 (Dutrey et al. 2014; Phuong et al. 2018a).

• Abは近接連星系のため，星周円盤の放射が弱い。
• 用いたガス輝線のデータ→

Fig. 1. Continuum image at 330 GHz. The beam size of
0.1900 ⇥ 0.1400 , PA = 190 is indicated in the lower left corner. The contour levels are 0.5, 1, 2, 4, 8 and 16 mJy beam 1 . The noise level is
0.06 mJy beam 1 . The red ellipses indicate the inner and outer edges ofthe dense dust ring at 180 au and 260 au, respectively (e.g., Guilloteau
et al. 1999). The positions of the Aa star and Ab close binary are indicate by the crosses. The inset shows an enlarged view of the Aa and Ab
surroundings.
A&A 635, A12 (2020)

3.2.2. CO gas kinematics in the outer disk

For a thin disk, the line-of-sight velocity is given by
ChannelNoise
Vz = sin spacingi(Vrot sin ! + VfallBeamsizecos !) where i is thedisk inclination,
1)(K)
project
Vrot (kmthe srotation
velocity and Vfall the infall velocity,
and !ALMAisthe azimuth
disk00plane.
00 can neglect the infall motions,0.106 in the0.35⇥ 0.31We, PA = 104
1.8
2011.0.00059.S
becausethe studies0.34
of Phuong
et00 ,al.
(2018a)at angularresolution
00
0.106
⇥
0.28
PA
=
91
0.7
2012.1.00129.S00
of ⇠0.35 have placed 00an upper00limit of 9% on it with respect to0.110
0.22 ⇥ 0.16 , PA = 15
0.7 2012.1.00129.S, 2015.1.00224.S
the rotation. We used theintensity-weighted
imagesof the line0.110
0.1900 ⇥ 0.1400 , PA = 19
1.9
2015.1.00224.S
of-sight velocity Vz shown
in
Fig.
2
and
for
each
pixel
calculate–
0.1900 ⇥ 0.1400 , PA = 10
0.03
2015.1.00224.S
1
Vz (sin i sin !) . We then azimuthally averaged these values for
all pixels at the same radius (using a 0.1500 radial binning) such

13CO & C18O (3-2)のイメージ
A&A 635, A12 (2020)

表面輝度の動径＆方位角分布
N. T. Phuong et al.: Gas properties around GG Tauri A

Fig. 2. Upper panels: 13 CO(3–2) integrated
intensity map (left, contour spacing 0.01
Jy beam 1 km s 1 (3 ) with zero level omitted) and velocity map (right). Lower panels:
C18 O(3–2) integrated intensity map (left, contour spacing 0.01 Jy beam 1 km s 1 , 2 ) andvelocity map(right).The
beams
aredependence
indicated of the integrated brightness temperature (for lines) and brightness temperature (for continuum emission)
Fig. 3.
Upper panels:
radialin lower leftin eachintensity map.
the cornerdisk plane.
The horizontal
sticksThe
indicate the angular resolutions. Lower panels: azimuthal dependence of the same quantities averaged over t00
00
contour levelspacing
mapsplotis of the emission from the J = 6–5, 3–2 and 2–1 lines of 12 CO. For a description of the corresponding daring of
1.2 <inr <the2.0velocity. Left panels:
0.5 km s 1 see. Thecontinuum
has forbeen
Tang
et al. (2016)
thesubtracted.
J = 3–2 line and Dutrey et al. (2014) for the others. Right panels: less abundant CO isotopologs (J = 3–2) emissio
gray region
delineates
theanddustouterring in the upper panels. In the lower panels, black arrows show the location of the limb-brightening peaks a
The whiteTheellipsesindicate
the inner
thedustmagentaedges of thering.lines show the hot-spot location.

この研究：リングより外側(1.25” or 169au以遠)のケプラー円盤成
分を滑らかな円盤モデルで近似し，Visibility dataから円盤モデル
成分を引き去って改めてイメージング：リング内のガスを抽出

deviation as the error on the mean. Using the three indepen13

18

Levenberg–Marquardt
method,and the error bars were derived
about 160 au. The extended emission
from this model
is betfrom the covariance matrix.

spectral line data, so that the continuum does not introduce asignificant bias in the combined fit for spectral lines described外側の円盤成分を引き去った後で作ったCO, 13CO, C18O (J=3-2)
N. T. Phuong et al.: Gas properties around GG Tauri A

Fig. 5. Integrated intensity of (residual) emissions after subtracting the best ring+disk models. Left: 12 CO(3–2). The restoring beam of 0.3400 ⇥0.2800 , PA = 89 is indicated in the lower left corner. Middle: 13 CO(3–2). The restoring beam is 0.2200 ⇥ 0.1600 , PA = 16 . Right: C18 O(3–2). The
restoring beam is 0.1900 ⇥ 0.1400 , PA = 167 . The contour level is 0.07 Jy beam 1 km s 1 (5 ) for 12 CO(3–2) map and ⇠0.02 Jy beam 1 km s 1(3 ) for 13 CO(3–2) and C18 O(3–2) maps, the zero level is omitted. The ellipses show the inner and outer edges of the dust ring at 180 and 260 au.
Table 4. Temperature derived from 12 CO(3–2) and surface density from
13
18
CO(3–2)andC
O(3–2).
• リング内のガス成分がきれいに残っている
12COに見られる南東側のhot spot: 惑星の兆候？）• 外側の差し引き残り：非軸対称性に起因（(1) (2) (3)
(4)
(5)
(6)
(7)

r
(au)

Tk
(K)

13

(K)

(K)

CO
C18 O
1015 cm 2

Ratio

議論1. Dynamics inside the cavity
A&A 635, A12 (2020)(1) 円環内での方位角方向の視線速度分布
Table 5. Infall and rotation velocities ofA&A
the gas inside
cavity.
635, theA12(2020)Ring

VKep( km s 1 )Magentaの線

Vrot( km s 1 )Vfall( km s 1 )

Vfall
Vrot

(a)
(b)
(c)
(d)
(e)

–
–
3.63
3.07
2.71

0.34
0.79
0.98
1.08
1.27

0.04
0.21
0.30
0.38
0.48

12%
27%
31%
28%
38%

Kepler@1.36M⦿。

多くのパネルであってないのはバイナリの影
Notes. The rings (a)–(e) are defined in Fig. 8.
響＋有限分解能。
Greenの線 in (f)Infallをわずかに加えたもの。純粋な回転よりよくフィット
Fig. 8. Dependence of hVz i ( km s 1 ) on the azimuth ! ( ) in the cavity.
12
CO (3–2) is in black, 13 CO (3–2) in red and C18 O (3–2) in blue. The
red curve is a fit of a sine function to the 13 CO (3–2) data (see text). We
use a blank space when no good data are available (no emission) in thering. The magenta curves show the expected Keplerian velocity arounda single star of 1.36 M . The green curve in panel f shows the best-fit
velocity curve when infall motions are allowed, superimposed on the13
CO velocity.

Vfall
VKep

Table 5. Infall and rotation velocities of the gas inside the cavity.
––
8%Ring
12%
18%

(a)
(b)
(c)
(d)
(e)

VKep
( km s 1 )Vrot( km s 1 )Vfall( km s 1 )

Vfall
Vrot

Vfall
VKep

–
–
3.63
3.07
2.71

0.34
0.79
0.98
1.08
1.27

0.04
0.21
0.30
0.38
0.48

12%
27%
31%
28%
38%

–
–
8%
12%
18%

Notes. The rings (a)–(e) are defined in Fig. 8.

Notes. The rings (a)–(e) are defined in Fig. 8.

議論1. Dynamics inside the cavity(2) 長軸・短軸方向でのPV図

Fig. 8. Dependence of hVz i ( km s 1 ) on the azimuth ! ( ) in the cavity.
12
13
18
CO (3–2)Dust
is in black,
CO
(3–2)
in
red
and
C
O (3–2) in blue.
Theringのinneredge（横線）
でred curve is a fit of a sine function to the 13 CO (3–2) data (see text). We
use a blank space when no good data are available (no emission) in the
0.3-0.4km/sの落下速度を示唆
ring. The magenta
curves show the expected Keplerian velocity →arounda single star of 1.36 M . The green curve in panel f shows the best-fit
velocity curve when infall motions are allowed, superimposed on the13
CO velocity.

0.2500 each of 12 CO(3–2), 13 CO(3–2), and C18 O(3–2) emissions
in the region 0 < r < 1.2500 . Azimuth and radius are defined in
the disk plane, i.e. deprojected from the disk inclination. In each
ring, we fit the azimuthal dependency of hVz / sin(i)i of the 13 COwith a sine function Vz / sin(i) = Vz0 sin !. This sine function ispresented as the smooth red curve and the amplitude Vz0 is provided at the top of each panel. The good fit for the 13 CO(3–2)
indicates that the gas inside the cavity is dominated by rotation. Thep amplitude is smaller than that of the Keplerian velocityv(r) = GMtot /r, however, where Mtot is the total mass of theFig. 9. Position-velocity diagrams of the 13 CO(3–2) emission in the cavity along the major (upper panel) and minor axis (lower panel). The
black curves show the expected Keplerian velocity around a single starof 1.36 M . Contour levels are spaced by 10 mJy beam 1 ; the zero contour is omitted. The white lines indicate the position of the inner edge ofthe dust ring (180 au) and the black lines show the position of the inner
radius of the gas disk (169 au). The data were rotated by 7 to align withthe disk axis, so that cardinal directions are approximate.

5(0.45, 0.18)

0.48

2.5

6

(0, 0.1)

0.12

6.1

>1.0 ⇥ 104>1.0 ⇥ 104

17(2.2+1.0
0.8 ) ⇥ 10
16
(4.0+1.4
1.4 ) ⇥ 10

議論2. Gas Properties

40 ± 56(2.2+1.0(3.9 ± 0.1) ⇥ 10 7 (1.7 ± 0.1) ⇥ 10 6
0.8 ) ⇥ 10

780 ± 10 (4.0+1.4(3.4 ± 0.1) ⇥ 10 7 (2.8 ± 0.1) ⇥ 10 6
1.4 ) ⇥ 10

Notes. (1) Blob, (2) offset from ring center, (3) distance from center, (4) line-width ( km s 1 ), (5) H2 density, (6) CO column density, (7) kinetic
temperature, (8) H2 mass derived from the CO column density (non-LTE analysis), (9) H2 mass derived from the 12 CO flux and (10) H2 massderived from the 13 CO flux.

6箇所の代表点で，マルチラインnon-LTE解析

bright in CO(6–5) (see Fig. 10). Integrated line flux and line
width were derived for each blob by fitting a Gaussian to theline profile. Velocities and line width derived from the CO (3–2)
N. T. Phuong et al.: Gas properties around GG Tauri A
were used to determine the 13 CO and C18 O line intensities.
To determine the physical conditions, we used a non-LTE
escapeprobability
transfercode implemented inTable 6.
Propertiesradiativeof the brighter
blobs.
DiskFit. It uses the escape probability formulation of Elitzur(1992), = [1 exp( ⌧)]/⌧, a single collision partner, H2 , and12
13Gaussian
profiles. Non-LTE
best-fit
solutionswere foundBlob line
Position
Radiusdv
HdensityN
T
Mass(nLTE)
Mass
(
CO
Flux)
Mass
(
CO Flux)
2kin
2
by sampling thesurface,00which is definedas the 3quadratic(00 ,00 )( )( km s 1 )(cm )(cm 2 )
(K)
(M )
(M )
(M )
sum of the difference between the measured brightness tem(1)2) computed(3)(5)CO (3–2),
(6)
(7)
(8)
(9)
(10)
peraturesand (thevalues of(4)the CO (6–5),
13
18
CO(3–2),andC
O (3–2)0.45
transitions,for rangesof⇥H10
4
+0.6
17
+0.6
6
7
6
2 density
12(0.27,
0.36)
2.5
>5.0
(2.1
)
⇥
10
40
±
5
(2.1
)
⇥
10
(3.1
±
0.1)
⇥
10
(1.9
±
0.1)
⇥
10
0.7
of 10 1010 cm 3 , 12 CO column density of 1013 1019 cm 2 , and 0.7
4
17
+0.7
6
7
6
2 temperature( 0.09, 0.36)
0.37K using2.7⇥ 10param(1.4+0.7
)
⇥
10
50
±
5
(1.3
)
⇥
10
(4.5
±
0.1)
⇥
10
(2.3
±
0.1)
⇥
10
kinetic
of 3 100
50 steps>1.0
of each0.5
0.5
12 13 4eter.3 We (assumed
the standard
isotopicratiosC/
C = (2.6
70 +1.0 ) ⇥ 1017 40 ± 5 (2.5+1.0 ) ⇥ 10 6 (3.2 ± 0.2) ⇥ 10 7 (1.1 ± 0.1) ⇥ 10 6
0.45, 0.36)
0.58
2.1
>5.0
⇥
10
1.0
(Milam et al. 2005) and 16 O/18 O = 550 (Wilson 1999) for the 1.00
5
16
+2.1
7
7
6
12
4 abundances(0.09, 0.15)
0.17
6.2>1.0 ⇥ 10(6.3+2.1
)
⇥
10
80
±
10
(6.3
)
⇥
10
(6.9
±
0.1)
⇥
10
(2.5
±
0.2)
⇥
10
relativeof the isotopologs.
CO constrains
the tem1.4
1.4
13
17
+1.0
6
7
6
perature,
the column
to its
5 18 and(0.45,CO0.18)
0.48 densities.
2.5 Owing>1.0⇥ faintness,
104 (2.2+1.0
)
⇥
10
40
±
5
(2.2
)
⇥
10
(3.9
±
0.1)
⇥
10
(1.7
±
0.1)
⇥
10
0.8
0.8
the C O(3–2) data bring little information. Because the critical
4
+1.4
16
+1.4
7
7
6
6
(0,
0.1)
0.12
6.1
>1.0
⇥
10
(4.0
)
⇥
10
80
±
10
(4.0
)
⇥
10
(3.4
±
0.1)
⇥
10
(2.8
±
0.1)
⇥
10
densities of the observed transitions are low, we only obtain a 1.4
1.4
lower limit to the density. The blob parameters are presented inTable 6.
Notes. (1) Blob, (2) offset from ring center, (3) distance from center, (4) line-width ( km s 1 ), (5) H2 density, (6) CO column density, (7) kinetic
We typically find high CO column densities of about a few
12temperature,
(8)
H
mass
derivedfromtheCOcolumndensity(non-LTE
analysis),
(9)
H
mass
derivedfromtheCO flux and (10) H2 mass1722
2
⇠10 cm and temperatures
in
the
range
40–80
K,
with
a
lower
13
derivedfromthe
COabout
flux.105 cm 3 . For blob 6, the faintest
limit on the density ofregion analyzed with this method, the problem is marginallydegenerate, with two separate solutions: (i) a high column denbright in CO(6–5) (see Fig. 10). Integrated line flux and line
sity (⇠1017 cm 2 ) and low temperature (⇠20 K), and (ii) a low
width were derived for each blob by fitting a Gaussian to thecolumn density (⇠1015 cm 2 ) and high temperature (>80 K).
This region lies between Aa and Ab, therefore the secondline profile. Velocities and line width derived from the CO (3–2)
2
solution (which is also that of lowest ) is more probable.
were used to determine the 13 CO and C18 O line intensities.
5.3. Gas masses
Fig. 10. Integrated intensity map of 12 CO(3–2) (upper panel) (thiswork and Tang et al. 2016) and 12 CO(6–5) (lower panel) (from Dutrey
et al. 2014) and positions and sizes of the blobs. The crosses mark thepositions of Aa and Ab, and the ellipse is the inner edge of the dust ring(180 au).

(0.3500 ⇥ 0.3000 ), we identify 2 dominant features. We decom-

17 cm-2n(H2) > 104 cc-1; ToNdetermine~10thephysical conditions, we used a non-LTE

The lower limit on the density obtained from the non-LTE analysis was insufficient to provide any useful constraint on theblob masses, we therefore used another method. We estimatedthe blob mass from the derived molecular column density andblob size, assuming a molecular abundance relative to H2 , asdescribed below.
In the same way, we also derived the total amount of gas

escape probability radiative transfer code implemented in
DiskFit. It uses the escape probability formulation of Elitzur(1992), = [1 exp( ⌧)]/⌧, a single collision partner, H2 , andGaussian line profiles. Non-LTE best-fit solutions were foundby sampling the 2 surface, which is defined as the quadratic
sum of the difference between the measured brightness temperatures and the computed values of the CO (6–5), CO (3–2),

Table 7. Mass of gas inside the cavity.

議論3. Gas Masses & inner diskへの質量供給
Location

Integrated flux
(Jy km s 1 )H2 mass(M )LTEで質量を導出

3.8 ± 0.1
0.5 ± 0.2(†)
1.6 ± 0.1 ⇥ 10 44(‡)
1.6 ±A&A
0.8 ⇥ 10635,
A12 (2020)

Notes. (†) X[13 CO] = X[12 CO]/70, and (‡) X[C18 O] = X[12 CO]/550 (seetext).

Table 7. Mass of gas inside the cavity.

Streamers:
warmer CO

Inner Disks:
NIR dust, H2, warm CO
10 μm Si feature

NORTH
Near side

Abundance(w.r.t H2 )Cavity 内のintegrated fluxから温度を仮定した
Cavity (12 CO)
11.4 ± 0.8
6.1 ± 0.4 ⇥ 10 6 8.0 ⇥ 10 5
Cavity (13 CO)
Cavity (C18 O)

molecular rotational lines
CO, CS, DCO+, HCO+, H2S

800

260

–(
200*+
–(
!./$0 = 14200*+!"#$ = 27

180

decide300

side
Inner Disks:
Inner Far
Disks:
molecularCOrotational
Streamers:NIR dust,
snow-line lines
Disk Accretion,
H2, warm CO+ peak)+, HCO+, Hshocked
NIR
dust, H2, warm CO=20 K (DCO
Gas
&
Dust:
CO, CS,TkDCO
S
warmerCO
10 μm Si feature,2
SOUTH
molecular tracers, e.g. H2
sub-mm CO
dust Si feature
10& μmNORTH

Near side
In the optically thin approximation, the integrated flux and
CO streamers
NORTH
the column density of the upper level of a given transition are
Near side
Location
Integrated flux
Abundance
related by H2 mass
1
(Jy km s ) gu
(M )(w.r.t H2 )W = Nu ,
(1)
u
12
6
5
–(
Cavity ( CO)
11.4 ± 0.8
6.1 ±R 0.4 ⇥ 10
8.0 ⇥ 10!"#$ = 27
200*+
where W = T b dv is the integrated
brightness
inside the cav- 800
260
180
13
4
(†)
–(
Cavity ( CO)
3.8 ± 0.1 ity (R 1.6
±
0.1
⇥
10
decide300< 160 au), gu = 2J + 1 is the statistics weight and Nu is!./$0 = 14
200*+
18
4
(‡) hc3 Aul
Cavity (C O)
0.5 ± 0.2 the column
1.6 ±density
0.8 of⇥ the
10upper level, u = 8⇡ kB ⌫2 (the Einstein
Far side3
SOUTH
coefficient Aul is taken from Lamda database ). Guided by the
results of(‡)
the non-LTE
(†)
13
12
18 analysis,
12we assumed that the gas temperNotes. X[ CO] = X[ CO]/70,atureandT is 40X[CO] = X[
CO]/550(seecalculated the Fig. 11. Schematic view of the GG Tauri system. Colored text indicates
K everywhere
insidethe cavity andresults obtained from other publications. In particular, magenta shows
total column density Ntotal of a given molecule astext).
results from Phuong et al. (2018b); see Dutrey et al. (2016)
fora review Far side
Inner
Disks:!
ofCO
otherreferences. BlackDisk
text shows
results fromNIR
thisdust,
work.H , warm CO
snow-line
Accretion,
Nu
Eu
2
Ntotal =exp(2) Tk=20 K (DCO+ peak) shocked Gas & Dust:
10 μm Si feature,
SOUTH
Z
kB T
molecular tracers, e.g. H2
12COはopticallysub-mmCO&
dust
thickなのでlower
analysis of dust images between 3 and 0.45 mm using a simple
In the optically thinapproximation, limit
the integrated flux andwhere, Z is the partition function and Eu is the energy of theCavity内での平均的な柱密度と落下速度
power law.
Because this study
is based on multi-wavelength
CO streamers
NORTH
the13 column 18
density of the upper
levelTheof12 CO(3–2)
a giventransition
are optically continuum resolved observations from 3 up to Near
upper
state.
emission,which
ispartially
0.45sidemm, theCOとC Oは標準的な同位体組成比の下で，thick, yields a lower limit.
derived dust temperature is characteristicof the large
grainsrelated by
-8 M
-1
→
内側円盤へは~6.4×10yr
⦿The CO abundance was taken from the abundances measured that have likely settled down around the midplane.
It is then
良い一致を示した
in TMC-1 by Ohishi et al. (1992), and we assumed a standard reasonable to consider that this temperature traces the midplanegu
13
18
BLOBSsteep
Hαから求めたGG
Tau
への降着率
isotopic
ratio
for
the
isotopologs
(
CO
and
C
O).
The
results
dust
temperature.
ThisradialAa
slope&bof thetemperatures is
W = Nu ,(1)
40–80 Kare listed in Cols. 9–10 of Table 6 for the blobs, and Table 7 most likely due TNto ==thestellar light thatis blocked by the inner
10 cm
CAVITY
u
-8and
n yr
=rest
10 -1
cm
summarizes the results for the cavity.
dense ring,theof)と割と良い一致
the diskremains
in its shadow.
(~2×10M
M then=1.6×10M⦿
M =6.4×10 M /yr
The H2 mass derived from 13 CO(3–2) and C18 O(3–2) is simR
BLOBS

Tkin = 40–80 K
NCO = 1017 cm–2
nH2 = 107cm–3

CAVITYMgas=1.6×10–4 Msun
Macc=6.4×10–8 Msun/yr

Inner Disks:
sub-mm CO & dust

kin

CO

H2

17

7

–2

–3

gas
acc

–4

–8

sun

sun

#19

Protostellar Outflows at the EarliesT Stages (POETS). IV. Statistical properties of the
22 GHz H2 O masers
L. Moscadelli1, A. Sanna2,3 , C. Goddi4,5 , V. Krishnan1,6 , F. Massi1 , and F. Bacciotti11INAF-Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Firenze, Italy; 2 INAF-Istituto di Radioastronomia& Italian ALMA Regional Centre, Via P. Gobetti 101, I-40129 Bologna, Italy; 3 Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany; 4 Leiden Observatory, Leiden University, PO Box 9513, 2300
RA Leiden, The Netherlands; 5 Department of Astrophysics/IMAPP, Radboud University, PO Box 9010, 6500 GL,
Nijmegen, The Netherlands; 6 South African Radio Astronomy Observatory (SARAO), 2 Fir street, Black River Park,
Observatory, Cape Town, 7925, South Africa
E-mail contact: mosca at arcetri.astro.it

VLBAを用いた
水メーザーのサーベイ，スポットの統計的情報（固有運動含み）や熱放射点源との距離などを統計的に議論

We wish to perform a statistical study of the location and motion of individual 22 GHz water maser cloudlets,characterized by sizes that are within a few au, with respect to the radio thermal emission from young stellar objects(YSO). For this purpose, we have been carrying out the Protostellar Outflows at the EarliesT Stages (POETS)
survey of a sample (38) of high-mass YSOs. The water maser positions and three-dimensional (3D) velocities were
determined through Very Long Baseline Array observations with accuracies of a few milliarcsec (mas) and a few
km s−1 , respectively. The position of the ionized core of the protostellar wind, marking the YSO, was determinedthrough sensitive continuum Jansky Very Large Array observations with a typical error of 20 mas. The statistic of the
11
separation of the water masers from the radio continuum showsthat 84% of the masers are found within 1000 au fromthe YSO and 45% of them are within 200 au. Therefore, we can conclude that the 22 GHz water masers are a reliable
proxy for the YSO position. The distribution of maser luminosity is strongly peaked towards low values, indicating thatabout half of the maser population is still undetected with the current Very Long Baseline Interferometry detection
thresholds of 50–100 mJy beam−1 . Next-generation, sensitive radio interferometers will exploit these weak masers for
an improved sampling of the velocity and magnetic fields around the YSOs. The average direction of the water maser
proper motions provides a statistically-significant estimate for the orientation of the jet emitted by the YSO: 55% of
the maser proper motions are directed on the sky within an angle of 30◦ from the jet axis. Finally, we show that ourmeasurements of 3D maser velocities statistically support models in which water maser emission arises from planar
shocks with propagation direction close to the plane of the sky.

Accepted by A&A
http://arxiv.org/pdf/2002.06871

#20

Rocking shadows in broken circumbinary discs
Rebecca Nealon1 , Daniel J. Price2 and Christophe Pinte2
1

Department of Physics and Astronomy, University of Leicester, Leicester, LE1 7RH, United Kingdom; 2 School ofPhysics and Astronomy, Monash University, Clayton Vic 3800, Australia
E-mail contact: rebecca.nealon at leicester.ac.uk

We use three dimensional simulations with coupled hydrodynamics and Monte Carlo radiative transfer to show thatshadows cast by the inner disc in broken circumbinary discs move within a confined range of position angles on theouter disc. Over time, shadows appear to rock back and forth in azimuth as the inner disc precesses. The effect occurs
because the inner disc precesses around a vector that is not the angular momentum vector of the outer disc. We relate
our findings to recent observations of shadows in discs
Accepted by MNRAS Letters

2

R. Nealon et al.

http://arxiv.org/pdf/2002.02983

Episodic accretion constrained by a rich cluster of outflows
T. Nony1 , F. Motte1 , F. Louvet2,3,4 , A. Plunkett5 , A. Gusdorf4,6 et al.
1

Univ. Grenoble Alpes, CNRS, IPAG, 38000 Grenoble, France; 2 Departmento de Astronomia de Chile, Universidad
de Chile, Santiago, Chile; 3 AIM, CEA, CNRS, Universit Paris-Saclay, Universit Paris Diderot, Sorbonne Paris Cit,
91191 Gif-sur-Yvette, France; 4 Observatoire de Paris, PSL University, Sorbonne Universit, LERMA, 75014 Paris,
France; 5 National Radio Astronomy Observatory, Charlottesville, VA 22903, USA; 6 Laboratoire de Physique de lcole
mis-alignしたinnerdiskがouterdiskにNormale Suprieure, ENS, Universit PSL, CNRS, Sorbonne Universit, Universit de Paris, Paris, France

落とす影の時間変化をE-mailcontact: thomas.nony at univ-grenoble-alpes.fr
The シミュレーション
accretion history of protostars remains widely mysterious even though it represents one of the best ways to understand the protostellar collapse that leads to the formation of stars. Molecular outflows are here used to characterize
the protostellar accretion phase in W43-MM1. The W43-MM1 protocluster host a sufficient number of protostars tostatistically investigate molecular outflows in a single, homogeneous region. We used the CO(2–1) and SiO(5–4) line
1. Rocking
shadows
in broken circumbinary
over one precession
period of the inner disc, assuming
a binary orbital
plane misaligned to outerdatacubes, taken as part of an ALMA mosaic with Figure
a 2000AU
resolution,
todiscssearchfor protostellar
outflows,
evaluate
disc by 60 degrees. Observer is face-on to the outer disc. Upper row: column density. Middle row: synthetic scattered light images. White arrows indicate
instantaneous directionof motion of the shadow and andthe grey circle
in first
panel shows simulated beam
width
of 15mas. Lower
row: CO (3-2) Moment 1 map.
the influence that the environment has on these outflows’characteristics
put
constraints
onoutflow
variability
Times shown in panels correspond to dotted lines in Figures 2 and the panels in Figure 3. All panels are shown on the same spatial scale, indicated in theright-most panel ofeach row. driven by 27 protostars with masses of 1–100 M .
in W43-MM1. We discovered a rich cluster of 46 outflow
lobes,#
The complex environment inside which these outflow lobes develop has a definite influence on their length, limiting

Although the gas in the cavity is only a small fraction of the disk mass, the mass accretion rate throughout the cavity
is comparable to or higher than the stellar accretion rate. It is accordingly sufficient to sustain the circumstellar disks
on a long timescale.

#23

Accepted by Astronomy & Astrophysics
http://arxiv.org/pdf/2001.08147

Planet induced spirals in the circumbinary disk of GG Tau A
N.T. Phuong1 , A. Dutrey 2 , E. Di Folco2 , S. Guilloteau2 , A. Pierens2 , J. Bary3 , Tracy L. Beck4 , E.
Chapillon2,5 , O. Denis-Alpizar6 , P.N. Diep1 , L. Majumdar7 , , V. Piétu5 , and Y-W. Tang8 ,1

Department of Astrophysics, Vietnam National Space Center, Vietnam Academy of Science and Techonology, 18
Hoang Quoc Viet, Cau Giay, Hanoi, Vietnam

CO (J=2-1)の0.3”データを解析
Protoplanetの作るスパイラルらしき2

Laboratoire d’Astrophysique de Bordeaux, Université de Bordeaux, CNRS, B18N, Allée Geoffroy Saint-Hilaire, F33615 Pessac, France
13
3
Department of Physics and Astronomy, Colgate University, 13 Oak Drive, Hamilton, New York 13346, USA
4Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, Maryland 21218, USA
5
IRAM, 300 rue de la piscine, F-38406 Saint Martin d’Hères Cedex, France
6
Instituto de Ciencias Químicas Aplicadas, Facultad de Ingeniería, Universidad Autónoma de Chile, Av. Pedro deValdivia 425, 7500912 Providencia, Santiago, Chile
7
School of Earth and Planetary Sciences, National Institute of Science Education and Research, HBNI, Jatni 752050,
Odisha, India
8
Academia Sinica Institute of Astronomy and Astrophysics, PO Box 23-141, Taipei 106, Taiwan

兆候を「武藤公式」でフィットして議論
N. T. Phuong et al.: Planet-induced spirals in the circumbinary disk of GG Tauri A

E-mail contact: ntphuong02@vnsc.org.vn
ALMA high angular resolution observations of the dust and CO emission have already revealed signatures of protoplanets embedded in protoplanetary disks. These detections are around single T Tauri stars while exo-planet surveys
reveal that planets can also form in binary (or multiple) systems, either in circumstellar or circumbinary orbits. We
searched for indirect evidence for planet formation in the multiple system GG Tau A that harbors the most massive
circumbinary disk among T Tauri stars.
We performed CO(2–1) ALMA Cycle 6 observations of GG Tau A at 0.3!! resolution. The images confirm the “hotspot” detected at higher frequencies, but also reveal prominent spiral-like features. We model these features by using
the analytic prescription for the linear perturbation regime induced by low-mass planets.
The brightest spiral is well reproduced by a density wave excited by a protoplanet (GG Tau Ac) at the hot-spot
location (290 au), just outside the dust ring. The absence of a clear gap (in gas or dust) at the planet location implies
that its mass is significantly lower than that of Jupiter, i.e. of the order of the mass of Neptune or less. Furthermore,other prominent (trailing) spiral patterns can be represented by adding one (or more) planet(s) at larger orbital radii,
with the most obvious candidate located near the 2:1 mean-motion resonance with GG Tau Ac.
The (proto-)planet GG Tau Ac appears to externally confine the ring in a stable configuration, explaining its highmass. Our results also suggest that planets similar in mass to Neptune may form in dense circumbinary disks orbiting(wide) binary stars. In the GG Tau case, orbital resonances appear to play an important role in shaping this multiple
circumbinary planet system.
Accepted by Astronomy and Astrophysics (A&A Letters)

Fig. 4. Spirals induced by the three putative planets, GG Tau Ac (black), GG Tau Ad (cyan), and GG
p Tau Ae (blue). The white ellipses indicate
the approximate inner and outer edge of the dust ring. The peak brightness has been multiplied by r and the color scale is saturated to optimize
the visibility of spiral patterns. Contours are given where the color is saturated.

and are therefore more efficient in halting inward migration, thestability of GG Tau Ae may nevertheless be reinforced by the 3:2
resonance with the second planet GG Tau Ad. This resonance isby far the most frequent resonant configuration observed among
Kepler multiple exoplanet systems (Fabrycky et al. 2014).
Pierens & Nelson (2013) also reported that a planetary configuration similar to the observed one is stable only if the outerplanet(s) are less massive than the inner one. The lack of clear

Andrews, S. M., Huang, J., Pérez, L. M., et al. 2018, ApJ, 869, L41
Avenhaus, H., Quanz, S. P., Garufi, A., et al. 2018, ApJ, 863, 44
Bae, J., Zhu, Z., Baruteau, C., et al. 2019, ApJ, 884, L41
Benisty, M., Juhasz, A., Boccaletti, A., et al. 2015, A&A, 578, L6
Beust, H., & Dutrey, A. 2005, A&A, 439, 585
Beust, H., & Dutrey, A. 2006, A&A, 446, 137
Boehler, Y., Weaver, E., Isella, A., et al. 2017, ApJ, 840, 60
Boehler, Y., Ricci, L., Weaver, E., et al. 2018, ApJ, 853, 162
Christiaens, V., Casassus, S., Perez, S., van der Plas, G., & Ménard, F. 2014, ApJ,
785, L12

#24

Accepted by Astronomy and Astrophysics (A&A Letters)

Chemistry Along Accretion Streams in a Viscously-Evolving Protoplanetary Disk
Ellen M. Price1 , L. Ilsedore Cleeves2 , and Karin I. Öberg1
1

Center for Astrophysics | Harvard & Smithsonian, 60 Garden St., Cambridge, MA 02138; 2 University of Virginia,
Department of Astronomy, 530 McCormick Rd., Charlottesville, VA 22904
E-mail contact: ellen.price at cfa.harvard.edu
The composition of a protoplanetary disk is set by a combination of interstellar inheritance and gas and grain surface
chemical reactions within the disk. The survival of inherited molecules, as well as the disk in situ chemistry depends
on the local temperature, density and irradiation environment, which can change over time due to stellar and disk
evolution, as well as transport in the disk. We address one aspect of this coupling between the physical and chemical
evolution in disks by following accretion streamlines of gas and small grains in the disk midplane, while simultaneously
taking the evolving star into account. This approach is computationally efficient and enables us to take into accountchanging physical conditions without reducing the chemical network. We find that many species are enhanced in the
inner disk midplane in the dynamic model due to inward transport of cosmic-ray driven chemical products, resulting
in, e.g., orders-of magnitude hydrocarbon enhancements at 1 au, compared to a static disk. For several other chemicalfamilies, there is no difference between the static and dynamic models, indicative of a robust chemical reset, while yetothers show differences between static and dynamic models that depend on complex interactions between physics andchemistry during the inward track. The importance of coupling dynamics and chemistry when modeling the chemical
evolution of protoplanetary disks is thus depends on what chemistry is of interest.
Accepted by ApJ
http://arxiv.org/pdf/2002.04651

14
円盤内での動径移流を考慮しながら円盤ガス化学反応を考察，外域の宇宙線の影響が大きな領域から素早く内側へとガスが輸送されることで，分子種によっては，Staticな円盤モデルのケースに比べて円盤内域で大きくエンハンスされるものもある(e.g., hydrocarbon)。