SFN #376 31-45 片岡

38.

Dynamics Near the Inner Dead-Zone Edges in a Proprotoplanetary Disk

Kazunari Iwasaki, Kengo Tomida, Shinsuke Takasao, Satoshi Okuzumi, Takeru K. Suzuki F We perform three-dimensional global non-ideal magnetohydrodynamic simulations of a protoplanetary disk containing the inner deadzone edge. We take into account realistic diﬀusion coeﬃcients of the Ohmic resistivity and ambipolar diﬀusion basedon detailed chemical reactions with a single-size dust grains. We found that the conventional dead zone identified
by the Elsässer numbers of the Ohmic resistivity and ambipolar diﬀusion is divided into two regions: "the transitionzone" and "the coherent zone". The coherent zone has the same properties as the conventional dead zone, and extends
outside of the transition zone in the radial direction. Between the active and coherent zones, we discover the transitionzone whose inner edge is identical to that of the conventional dead-zone. The transition zone extends out over theregions where thermal ionization determines diﬀusion coeﬃcients. The transition zone has completely diﬀerent physicalproperties than the conventional dead zone, the so-called undead zone, and zombie zone. Combination of amplificationof the radial magnetic field owing to the ambipolar diﬀusion and a steep radial gradient of the Ohmic diﬀusivity causes
the eﬃcient evacuation of the net vertical magnetic flux from the transition zone within several rotations. Surface gas
accretion occurs in the coherent zone but not in the transition zone. The presence of the transition zone prohibits
mass and magnetic flux transport from the coherent zone to the active zone. Mass accumulation occurs at both edgesof the transition zone as a result of mass supply from the active and coherent zones.

39. Meteorites and Planet Formation

Rhian H. Jones F Meteorites are a remarkable resource. They capture the imagination of people worldwide with theirspectacular entry through Earth’s atmosphere as fireballs, and their exotic character of being pieces of other worlds.
Scientifically, they are critical to interpreting the early stages of formation of the Solar System, as well as the geologicalevolution of asteroids, the Moon, and Mars, and they are vital to understanding planetary formation processes. With•

term indicates that the electrons are provided by photoionization of carbon and sulfur when ⌃r < ⌃FUV (Bai
原始惑星系円盤の非理想MHD計算。
2017).

1.00

O =1
A =1

10

0.75

z = zatm (R)global。

••

chargedgrains⇣(r, ✓)dust= ⇣CR (r,✓) + ⇣X (r, ✓) + ⇣radio .

•

(20)

The first source is the cosmic rays whose ionization rate is
初期にガス面密度が1auで500
g/cm2,
r^-1.
given by✓◆✓◆
中心星2.5太陽質量（大きめ？）。温度1au
+
⌃✓
⌃✓
17
1
⇣CR = 10
s
exp
+ exp
,
で1520Kになってそう。（f̲boost=2）。
⌃CR
⌃CR

(21)

where ⌃CR = 96 g cm 2 (Umebayashi & Nakano 1981),
+

0.50
5

z [au]

0.25
z [au]•For the non-thermal ionization, we use the data tableof the three conductivities based on a chemical reaction
+
+
+
+
+
+
network
of
e
,
H
,
He
,
C
,
H
,
HCO
,
Mg
in theOhm散逸、両極性拡散が入ってる。簡易ケ
3
gas phase and the charged dust grains using the methods
ミストリーとシングルサイズダストも入れて
described by Nakano et al. (2002) and Okuzumi (2009).
ある。 The table is a function of ⇢/B, T , and ⇢/⇣, where ⇣ is theionization rate that will be defined in Equation (20) and T1
is replaced with T fboost. For simplicity, the single size dust
ケミストリー、高温ではカリウムの熱電
grain with a radius of a = 0.1 µm and an internal density of
離。炭素と硫黄のphotoionization
2 g cm 3 is considered. In this paper, the dust-to-gas mass
ratio is set to 10 4 , assuming that dust growth reduces the
非熱的電離ではachemical
reactionamount of dust grains smaller than the interstellar value
networkofThee­,H+, He+,C+,ofH+3,⇠ 0.01.
following
three kinds
non-thermal
ionization
sourcesarein
taken
account,
HCO+,
Mg+
theintogas
phase and the

0

0.00
0.25

5
0.50
0.75

10
0

5
R [au]

6

4

10

1.00
0.0

2
log10 ini

0.5
R [au]
0

1.0

2Fig. 1. (Left) Initial setup of the simulation. The color map indicates the initial density profile in the code units, and the black lines indicate the streamlines of the magnetic fields. The red (blue) line corresponds to the condition⇤O = 1 (⇤A = 1), which show the OR (AD) dead zone boundary. The magenta lines show z = ±zatm (R) above which the warm atmospheres exist.
The right panel is the same as the left panel but it shows the zoom-in of theregion enclosed by the black rectangle shown in the left panel.

tial density profile in the code units, and the black lines indicate the streamtion rate6is
Publications of the Astronomical Society of Japan, (2018), Vol. 00, No. 0lines of the magnetic fields. The red (blue) line corresponds to the condition⇤O = 1 (⇤A = 1), which show the OR (AD) dead zone boundary. The maデッドゾーンの定義にはオーム散逸、両極性拡散そ
RMRI
R
Rch
◆
genta lines show z = ±zatm (R) above which the warm atmospheres exist.
6
Publications of the Astronomical Society of Japan, (2018), Vol. 00, No. 0
れぞれのエルサッサー数を利用。これが１を超える(a), (21) 2 The right panel is the same
( O,Aas)T the left panel but it shows the zoom-in of theと、磁場が強く、MRI
enclosedby the black rectangle shown in the left panel.
( O,A
)NT
0 regionlog ( O,A)

ano 1981),
integrated
t a given r.
-ray radiaula (Bai &by Igea &4.0

2O (t = 0)
A (t = 0)

4

O

⇤O ⌘
A

6

2
(t = 500 t )v(t = 500 t ) A
K0

K0⌘O ⌦ K, ⇤A ⌘

2vA⌘A ⌦ K

0.5

3.2

2.4

,

10

1.6

0.8

(a) Ohmid Resistivity

0.4
z [au]

log10( O,A)

•

0.0

O =1

(24)

zatm(R)
2H(R)

0.3

p
0.2
8 1
H(R)
10 are smaller than
100 unity, where10v1 A = B/ 4⇡⇢ is the Alfvén0.1
speed (Sano&Miyama
1999;
Blaes
&
Balbus1994;
7.5 (b)
RMRI
R
Rch
0.0
0.5
1999; Kunz & Balbus 2004;
Bai Diffusion& Stone(b) Ambipolar
5.0 Wardle et al.
= 1 z (R)
0.4
2.5 2011).
2H(R)
=1
0.3
In order for the MRI to occur at a given
height
z,
the
0.0
0.2
H(R)
measured lo0.1
2.5 wavelength of the maximum growing mode
Fig. 3. Two-dimensional distributions of (a) ⌘ and (b) ⌘ in the code units
0.0
at
t
=
500
t 0.25when the
MRI turbulenceis saturated.
dashed2.00
callyshouldbesmaller
than
the
scale
(Sano
5.0
(22)
0.00height
0.50
0.75 &1.00
1.25 The
1.50white1.75
lines show the contours of ⌘ and ⌘ at R
10[au] and 10 . The conditions of
101
10 Miyama 1999).
10 Using the expressions
10
the⇤ most
unstable⇤ of= 1 and= 1 are displayedby the red and blue lines in Panel (a) andz [au]

log10( O,A)A

O (t = 0)
A (t = 0)

O (t = 500 tK0 ),A (t = 500 tK0 )

A

4

2

K0

O

2

O

Radius [au]

O

A

A

(b),Fig.
respectively.
In each panel,
the twoofblacklines
3. Two-dimensional
distributions(a) ⌘dashed⌘A correspondin the code tounits
O and (b)atm

ここは縦磁場の消失に着目。コヒーレントゾーンは動径ガスフラックスがない

Publications of the Astronomical Society of Japan, (2018), Vol. 00, No. 0

10• 従来のデッドゾーンは「遷移ゾー

R⌘

RMRI

Rch

funnell field
funnellfield
formed by
B flux accretion

turbulent magenetic field

failed winddisk windfrom the transition-zone
inner edge
the inner boundary of
the computational
domainreflection of
acoustic waves

asymmetricsurface gas accretion

mid-plane gas accretionthe central starring formation
B flux concentration

efficient outward
B flux transport! hBz i ⇠ 0

quasi-periodic
inversion of the sign of hB i
due to the MRI activity

two density peaks

Active Zone

Transition Zone

Coherent Zone

R-axisFig. 6. A schematic picture showing our main findings with respect to the time-averaged gas dynamics of the disk. The orange arrows represent the gas
motion, and the blue lines show the poloidal magnetic field lines.

ン」と「コヒーレントゾーン」の2つの領域に分かれる。
• 遷移ゾーンは、活発なゾーンとコ
ヒーレントゾーン間の質量と磁束の輸送を妨げる。
• 遷移ゾーン内では、急激なオーミック抵抗の勾配とアンビポー
ラー拡散によるラジアル磁場の増幅が原因で、数回の回転以内に垂直磁束が効率的に排出される。
• コヒーレントゾーン内では表面ガスの降着が発生するが、遷移ゾー
ン内では発生しない。
• 質量は遷移ゾーンの両端で蓄積され、質量供給がアクティブゾーン
とコヒーレントゾーンから供給されるため、リング構造が形成される。

• 磁束密度が高いところで
ガス面密度が下がる。これによりリングのように見える。

enhancing the azimuthal velocities of small particles. This dynamic behavior arises from the modification of dust
trajectories due to turbulent eddies.

TIPSY:
TrajectoryofInfalling
Particlesin
StreamersaroundYoungstars.
Dynamical
31. analysis of the streamers around S CrA and HL Tau

Aashish Gupta, Anna Miotello, Jonathan P. Williams, Til Birnstiel, Michael Kuﬀmeier, Hsi-Wei Yen F Context.

Elongated trails of infalling gas, often referred to as "streamers," have recently been observed around young stellar
objects (YSOs) at diﬀerent evolutionary stages. This asymmetric infall of material can significantly alter star andplanet formation processes, especially in the more evolved YSOs. Aims. In order to ascertain the infalling nature ofobserved streamer-like structures and then systematically characterize their dynamics, we developed the code TIPSY(Trajectory of Infalling Particles in Streamers around Young stars). Methods. Using TIPSY, the streamer molecularline emission is first isolated from the disk emission. Then the streamer emission, which is eﬀectively a point cloudin three-dimensional (3D) position-position-velocity space, is simplified to a curve-like representation. The observed
streamer curve is then compared to the theoretical trajectories of infalling material. The best-fit trajectories are used
to constrain streamer features, such as the specific energy, the specific angular momenta, the infall timescale, andthe 3D morphology. Results. We used TIPSY to fit molecular-line ALMA observations of streamers around a Class
II binary system, S CrA, and a Class I/II protostar, HL Tau. Our results indicate that both of the streamers are
consistent with infalling motion. TIPSY results and mass estimates suggest that S CrA and HL Tau are accreting
material at a rate of & 27 Mjupiter Myr 1 and & 5 Mjupiter Myr 1 , respectively, which can significantly increase
the mass budget available to form planets. Conclusions. TIPSY can be used to assess whether the morphology andkinematics of observed streamers are consistent with infalling motion and to characterize their dynamics, which is
crucial for quantifying their impact on the protostellar systems.

32. Maser Flares Driven by Isothermal Shock Waves

M. D. Gray, S. Etoka, B. Pimpanuwat, A. M. S. Richards F We use 3D computer modelling to investigate thetimescales and radiative output from maser flares generated by the impact of shock-waves on astronomical unit-scale- 若い星の周囲に観測されるガスの流入経路、いわゆる「ストリーマー」が、星と惑星の形成過程に及ぼす影響を理解するために開発されたコードTIPSY（Trajectoryof Infalling Particles in Streamers around Youngstars）の紹介- クラスIIの連星系であるS CrAと、クラスI/IIの原始星であるHL Tauのストリーマーに対するTIPSYの適用例を紹介。

- S CrAのストリーマーは、流入運動と一致しており、動的パラメータが良好に制約されている 。

- HL Tauのストリーマーも流入運動と一致しているが、観測範囲が狭いため、微妙Aashish Gupta et al.: TIPSY: Trajectory of Infalling Particles in Streamers around Young stars- 物理モデルはUlrich-Cassen-Moosmana)(UCM) model (e.g., Ulrich 1976;
Cassen & Moosman
1981)
Aashish Gupta et al.: TIPSY: Trajectory of Infalling Particles in Streamers around Young stars-

Table 1: Parameters used to isolateb)and fit the HL Tau and S CrA
streamers

Line of sight
Decl.

𝜙
𝜽
R.A.

Parameter
Stellar mass [M ]Distance [pc]Systemic velocity [km / s]Min. RV o↵set [km / s]Max. RV o↵set [km / s]Min. RA o↵set [arcsec]
Max. RA o↵set [arcsec]
Min. Decl. o↵set [arcsec]e)Max. Decl. o↵set [arcsec]
Significance ( ) level

S CrA
2
160
5.86
4.5
7
2
15
-7
7
3

HL Tau
2.1
147
7.14
7
10
-3
-1
-3
0.5
4

c)

Intensity-weighted meansd)Best Fit

Best Fit

See Appendix D for more details on the stellar parameters used.

Fig. 1: Schematic diagram of coordinate axes used to computethe theoretical trajectories of infalling gas (green cloud) arounda protostellar system (orange star), as discussed in Sect. 2.1. !
r0
and !
v0 denote the initial position and velocity vector of infallinggas, respectively. !
r represents the position vector of gas at a future point in its trajectory (circumference of blue ellipse), with ✓and denoting the polar and azimuthal angles, respectively. The
dashed red arrows show the unit vectors x̂, ŷ, and ẑ, defined using the directions of !
r0 and !
v0 . Together they set the coordinate
frame in which TIPSY solves the Mendoza et al. (2009) equations. The gray plane represents the POS, with the overlaid darkgray arrows denoting the coordinate frame of our observations.

sources of emission, such as disks. Given the wide range of morphologies exhibited by disks and streamers in di↵erent sources,
which depend on the observational parameters and molecularlines used, it is hard to automate this step. Therefore, we visually examine the emission maps and define a boundary for a subcube that encompasses the streamer emission using RA, Decl.,
13
and RV limits. Next, we eliminate
pixelswithvaluesbe- pipeline. Panel a: Intensity-weighted velocity (moment 1) map in colors,
Fig. 2: Flowof S CrA
CO flux(2–1) datain the TIPSY
with contours representing the integrated intensity (moment 0; see Fig. E.1). The red segments in the bottom-left corners
low a specified noise ( ) level.overlaid
Table
1 scalelistsof the
depict a length
1,000 values
au. The pinkused
ellipsesfor
in the bottom-right corners depicts the beam size of the data. Panel b: Isometricprojection
the 3D
PPVAs
diagramof pixels
with seintensity > 5 in the whole field of view. Panel c: Isometric projection of theselecting streamers around S CrA
and ofHL
Tau.
long
as the
PPV diagram of an isolated and cleaned streamer. The red square and its error bars represent intensity-weighted means and standard
lected boundaries fully capturedeviations,the observed
respectively.streamer
Panel d: Sameemission,as Panel c, but with the best-fit trajectory, as represented by the black line. Black circles
denote the interpolated values of the theoretical trajectory, which are directly compared to the intensity-weighted means. Panel e:
the final results are not very sensitive
to thevaluesof these
Same as Panelb, but exact
with the best-fit
trajectory,
as represented by the black line. Note: 3D interactive versions of panels d and e areavailable online.
limits. This is because we primarily
rely on the central brighter
emission throughout the structure for the fitting, as described inthe molecule relative to H2 ,2 and fu uis the fraction of molecules
G andM⇤ Mrepresent
universal
mass
G and
universalgravitational
gravitationalconstantconstant andandA&Amass⇤ represent
the upper
upper
energy state
stateofofthe
thetransitiontransition(e.g.,(e.g.,
Berginet al.
proofs: inmanuscript
no. 48007corr
energyBerginet al.
Eu /Tof protostellar
system,respectively.
ofsystem,respectively.We
Wecan
cansumsumthemthemtotoget
get thethe 2013). Here,u /Ta) protostellar
Here, ffuucan
canbe
befurtherfurthercomputed
computed= E3e
asasfu f=u 3e/Q/Q(T ),(T ),mol.mol.
2: Fitting results for S CrA and HL Tau.
specific
totalenergy(T.E.),
which
tellususififthis
thegas toinsurrounded
specific
totalenergy(T.E.),
which
tell
the
gas
isisbeinaa where EEuu isisTable
thermore,
Gupta
etcan
al.can(2023)
found
system
the
upper
state
energyforforthethetransition,
T is
the
the
upper
state
energytransition,Tistheby(T.E.
0.1 <parsec-scale
clouds,appearing
as reflection
nebulae,andbound
ellipticalorbit<similarthestreamer
streameraroundS
boundelliptical
orbit
(T.E.
0, 0,
similartotothearoundSgas temperature,and(T(T) )isisthefunction
forforthe thetemperature,
andQQmol.
thepartition
Quantity
Spartition
CrA
HLfunction
Taumol.
⇠ 1000 au elongated structures revealed by 12 CO (2–1) Ata13 13 -227±12
CrA),
a boundparabolic
orbit
(T.E.
⇡
0,
similartothestreamerCrA),
a boundparabolic
orbit
(T.E.
⇡
0,
similartothestreamer
RA
o↵set
[AU]
1149±76
molecule.
We
took
A
and
E
valuesfor(2–1)(S CrA)cama Large Millimeter/submillimeter Array (ALMA) observau u values for CO
molecule. We took Atrans.
CO(2–1)(S CrA)trans. and E
Decl. o↵set [AU]618±123 + -159±15
77 1 1aroundTau),or tions,
anunboundhyperbolic
orbit(T.E.
(T.E.
0). However,+suggesting
infall of material
ontothe>>system.
around
HLHL
Tau),
or an
unboundhyperbolic
orbit
0).
to
be
6.038⇥10
s
and
15.87
K,
and
HCO
(3–2)
(HL
Tau)
to
be
6.038⇥10
s
and
15.87
K,
and
HCO
(3–2)
(HL
Tau)
LOS
o↵set
[AU]
300±150-1400±573
12
these CO!
observationssu↵ered !
fromcontamination from tosur-be 1.453⇥10 33 s 1 1and 25.68 K, respectively, from the Lei!
!
RA speed [km / s] s0.3±0.6from the LeiUsing
initialpositionr0()r0and)gas
and(seevelocity
vectorof inin-and toand 25.68 K,0.37±0.18
respectively,
Using
thetheinitialposition
velocityv(0vGupta)0 )vector
rounding (di↵use
Fig. F.1(inet al.of2023)
alsobe 1.453⇥10
Decl. speed[km/ s]-0.47±0.15
0.6±1.2
den
AtomicandMolecular
Database(Schöieretal.
2005).
The
falling
gas,
we
can
also
estimatethespecific
angularmomentumaprominent
absorption
featureclose
tosystemic
velocityof
theden
AtomicandMolecular
Database(Schöieretal.
2005).
The
falling
gas,
we
can
also
estimatethespecific
angularmomentum
LOS
speed
[km
/
s]
0.02±0.17
1.79±0.20
6
13
!
!
b)
2be2 1.45 ⇥ 10xvalues
were
taken
to
for
CO
(e.g.,
Huang
source,where ac)lot toof boundmaterial is
expected to be.
6
13
mol.
!
!
as
r
⇥
v
.
This
can
be
comparedtheangularmomentumofSpecific
kinetic
energy[km/s]
0.18±0.10
1.8±0.900xmol. values were taken
to 2be2 1.45
⇥ 10 for CO (e.g., Huang
Intensity-weighted
means
as r0 ⇥ v0 . This can be comparedtotheangularmomentumof
13
9+
Forroleourofanalysis
we material
used the inCO(2–1) observations,
Specific potential
energy[kmHCO/ s ] +(e.g.,-1.33±0.09-1.3±0.5
et
al.
2020)
and
10
Jørgensenetal.
2004).
9 for
the
disks
to
quantify
theinfalling
misaligning
etId.:al. 2020)
and 10for[AU
HCO(e.g.,
Jørgensen
et al. 2004).
the diskstoquantify
the
role
of
infallingmaterial
inmisaligning
Specific
angularmomentumkm/s]
791±218606±1803
taken
as
a
part
of
the
same
ALMA
project(Project
Gupta
et
al.:
TIPSY:
TrajectoryofInfalling
Particlesin
StreamersaroundYoungstars
For bothsources,
weassumedarepresentative
temperatureofthe Aashish
protoplanetarydisks,
ashas
beensuggestedbysomehydroInfallsources,
time [yr] we assumed a representative
8301±1358 2724±1237
temperature of
2019.1.01792.S),show a bymuchcleaner⇠ 1300For
au boththe protoplanetary disks,
as has beenwhichsuggestedsomehydro25
K, which is typical for gas at these ⇠100-1000 au scales (e.g.,
dynamic simulations(e.g.,(Fig.
Thieset E.1a).
al. 2011;
Ku↵meieral. 25streamer
2 andWe usedthe standardetpipeline
cal-K,
whichistypical
for
gas
at
these
⇠100-1000
au
scales(e.g.,
dynamic
simulations(e.g.,
Thies
et
al.
2011;
Ku↵meieretal.
Jørgensen
et al. 2005).
At thistemperature, Qmol. (T = 25 K) valdata Tau
and theimaging was
using"Briggs" weight2021). For the S CrAibrated
and HLstreamers,
wedonefindthespecific
infallingtimeが出るので、質量が見積
Jørgensenetal.
2005).
At
thistemperature,
Q
(T
=
25
K)
val00 the specific
TIPSY
results,
as
shown
in
Fig.
3,
demonstrate
that
all
mol.
2021).
For
the
S
CrA
and
HL
Tau
streamers,
we
find
ing
with
robust=0.5,
a
cell
size
of
0.05
,
and
"auto-multithresh"1ues, interpolated from values provided in Cologne Database forangularmomentummagnitudes
to
be
791
±
218
AU
km
s
and
1 streamer
thepoints of thefrom
simplified
streamercurve (intensity-weighted
ons,angularcanprovidecompleteinformation
about
the
tegratedflux(F),
assuming
opticallmasking
withtodefault
parameters.
We
detectedtheat interpolatedues,
valuesprovidedin
CologneDatabaseforstreamer
momentummagnitudes
be
791
±
218
AU
km
s
and
1
Molecular
Spectroscopy(Endresetal.
2016),
were
19.6
and
12.0
606 ± 1803 AU km &1s 5 , level,
respectively.
For reference,
もれていたのはちょっと面白いか？
means) can be fit within the error bars (intensity-weighted stanin all the relevant
channels. the specific Molecular
13+
Spectroscopy(Endresetal.
2016),
were
19.6
and
12.0
606
±
1803
AU
km
s
,
respectively.
For
reference,thespecific
for
CO
and
HCO
,
respectively.
We
computedF
tobedeviations) by+an infalling trajectory. The fitting resultsstreamer
areation ofinfallinggas
relativeto thattheprotostar.
angularmomentum(l)TIPSY
in
theresults
outer
part
of thea 100
austreamer
Kepleriansuggestoverallis consistent 13dard 20
2 For,HL
21energy is2confor
CO
and
HCO
respectively.
We
computedF
tobelisted
in
Table
2.
Tau,
the
specific
kinetic
streamerangularmomentum
(l)
in
the
outer
part
of
a
100
au
Keplerian
1.3
⇥
10
W
m
for
S
CrA
and
4.2
⇥
10
W
m
for
HL
Tau
with
being
a
trail
of
infalling
gas.
For
the
best
fits,
infalling
tradisk around a 2 M star p
protostar, similar to HL Tau, should be
221 within the
20with the specific22sistent
gravitational
potential
energy
1.3
⇥
10
W
m
for
S
CrA
and
4.2
⇥
10
W
m
for
HL
Tau
2.37m
4⇡D
F
1star
jectories
could
fit
all
the
ten
pointsinthesimplified
streamerbyintegrating
the
flux
of
the
isolatedstreameremission(panelc
disk⇠
arounda2
M
protostar,similarto
HL
Tau,shouldbeHstreamermeters
can
be
used
to
deriveotherphysicallyrel(l =p GM⇤ Rd)d , where G, M⇤ , and Rd are thee) 421 AU km serror bars, suggesting that the gas is roughly in a zero-energy(intensity-weighted means, red squares in Fig. 2) within theby
er1
M
&
,
integrating
the
flux
of
the
isolatedstreameremission(panel
c
in
Figs.
2
and
3)
over
both
the
positionandthevelocity.
streamer
⇠ 421
AU
km
s
(l
=
GM
R
,
where
G,
M
,
and
R
are
the
gravitational constant,the
protostellar
mass,
and
the
disk
radii,
parabolicorbit.
However,
TIPSY
could
not
constrainthetrajec⇤d⇤dBest Fit The best-fit paFit(intensity-weighted standard deviations).
rorBestbars
s. Forexample,specific
(per
unit
mass)
kinetic
A
h⌫x
f
in
Figs.
2
and
3)
over
both
the
positionandthevelocity.
toryof
infalling
particles
for
HL
Tau
well
(see
bottompanels,trans.
mol.
u
Using
these
values,weestimatedthestreamermassesto begravitational
constant,
theprotostellar
mass,
and
the
disk
radii,
respectively).
rameters, as given in Table 2, suggest that the material is strongly
4 further
5
2
2
2
Fig.
4),theseasdiscussed
in⇥
Sect.
4.1.
Using
values,weestimatedthestreamermasses
to
be
bound
to
the
protostars,
withthespecific
(per
unit
mass)
total
&
2.1
⇥
10
M
and
&
1.2
10
M
for
S
CrA
and
HL
Tau,
respectively).
stimatedas0.5⇥(v+v+v).
Similarly,
asAs TIPSY provides
complete
trajectoryof theinfalling see
RA theLOSgravitational
4Although these estimates5
Decl.
energy(kinetic
energy plus
potential energy;&
respectively.
notSinclude
parts
2.1
⇥
10
M
and
&
1.2
⇥
10
Mdofor
CrA and
HLofTau,
2gravitational
2
gas,
until
the
motionisdominatedbytheforce,weAs
TIPSY
providesthecomplete
trajectoryoftheinfallingSect.
4.2) of 1.1 ± 0.1is
kmdominateds . This suggests thatby
the observed
m
is
the
mass
of
a
hydrogenat
local
gravitational
potential
the4.whereDiscussion
Hthestreamers
beyondtheprimary
beams
of
the
interferometrespectively.
Althoughtheseestimates
donotincludepartsofis not anby
ejection
ofdistance)
unbound material,
which
also
the 3D(RA,
Decl.,
andLOS
morphology
gas,canuntil
theinfer
motionisstructure
dominatedthegravitational
force,
weshould be
ric observations,they canstill
be usedtoestimate
mass
infall A c
on hyperbolic
trajectories.
The
sizeTheseofobserved
streamer,which
thestreamers
beyondtheprimary
beams
of
the
interferomet4.1.
Data
requirementstothesource,A
isthe
Einsteinoftheinfallingstreamer,
as
shown
in
Fig.
B.1.
morpholotellar
system,specific
gravitational
potential
encan also infer the 3D (RA,Decl.,
andLOS
distance)morphology
trans.
is at qleast an order of magnitude larger than the protoplanetaryratesobservations,as Ṁin f =M/Tstill
where
T(Sect.
toHLtheinfallstreamerin f ,for
into
f refers
ric
they
can
be
used
estimate
mass
infallgies
canfurtherbevalidatedusing
near-infrared
polarization
obComparing
TIPSY
fitting
results
S
CrA
3.1)
and
of the infalling streamer,
as further
shownindicate
in Fig.
Thesemorpholodisks,
thatB.1.
this isnot a spiralarm inducedtime
by forline
transition,
⌫
is
the
line
frequency,xtheobserved
streamer.
Massinfall
rates
are
found
to
be
2
2
2
m
Tau
(Sect.
3.2),
we
can
see
that
the
uncertainties
aremuchhigher
rates
as
Ṁ
=
M
/T
,
where
T
referstotheinfallservations,asthedegreeofpolarizationinsuchobservations
caninfstreamerinf
inf
gravitationallyunstabledisk.
gies canfurtherbe validated
using
near-infrared
polarizationob- & 2.5for⇥ HLimated
as
G
⇥
M
/
d
+
d
+
d
,
where
8 (see Table
1 2). Moreover, Fig. 4b shows that
1 the dis⇤
Tau
10
M
yr
(or
&
27
M
Myr
)
for
S
CrA
and
RA
LOS
jupiter
Decl.
time
for
the
observed
streamer.
Massinfall
rates
are
found
to
be
Moreover,
we
also
find
that
the
velocity
profileofobservedbecorrelated
to
the
3D
orientationof
dust
structures(e.g.,
Ginski2themolecule
relative
to
H
,
and
f
is
the
servations,asthedegreeofpolarizationinsuchobservations
can2u
tribution9of
best-fit
parameters(higher
fitting
fractions,
lower
1
1
8M yr
1(or & 4.7 M jupiter Myr ) for1 HL Tau. In&
4.5⇥
10
Fig. 3: Same procedure as described in Figure
2, but for HL
Tau insteadcloseof S CrA.
Note:
3Dprotostars,
interactive versionsof ispanelsdsay,
and e the LOS
streamerchanges
to
the
that
to
⇥ 10 areM notyrwell (or& 27 by
M simple
Myr ) foretuniversalal.online.
2021).
betterunderstandingof the3D morphology
of mass
the & 2.5deviations)jupiter symmetricalrepresented
er- S CrA andavailable
esent
gravitational
constant
andbearecorrelated
toAthe3D
orientationof
dust
structures(e.g.,
Ginskivelocities stop decreasing and start increasing, which was notterestingly,
rethesevaluesarecomparable
tomassaccretion
rates
in
the
energystate
thetransiti
9 Tau.upper
1is likely
1 of
rors
for
HL
This
becausethe
HL
Tau
observations,streamercanbeusefulinconstraining
thelocationandvelocity
&
4.5
⇥
10
M
yr
(or
&
4.7
M
Myr
)
for
HL
Tau.
Injupiter
produced in our best-fit(panel d, Fig. 2).
et al. 2021). A better understandingof themodels
3D morphology
ofThisthesuggests
A&A proofs: manuscript no. 48007corr-

37. On the magnetic field properties of protostellar envelopes in Orion

Bo Huang, Josep M. Girart, Ian W. Stephens, Manuel Fernandez-Lopez, Hector G. Arce, John M. Carpenter,Paulo Cortes, Erin G. Cox, Rachel Friesen, Valentin J. M. Le Gouellec, Charles L. H. Hull, Nicole Karnath,
Woojin Kwon, Zhi-Yun Li, Leslie W. Looney, Tom Megeath, Philip C. Myers, Nadia M. Murillo, Jaime E.
Pineda, Sarah Sadavoy, Alvaro Sanchez-Monge, Patricio Sanhueza, John J. Tobin, Qizhou Zhang, James M.
Jackson, Dominique Segura-Cox F We present 870 um polarimetric observations toward 61 protostars in the Orion

molecular clouds, with 400 au (1") resolution using the Atacama Large Millimeter/submillimeter Array. We successfully detect dust polarization and outflow emission in 56 protostars, in 16 of them the polarization is likely produced
by self-scattering. Self-scattering signatures are seen in several Class 0 sources, suggesting that grain growth appears
to be significant in disks at earlier protostellar phases. For the rest of the protostars, the dust polarization traces themagnetic field, whose morphology can be approximately classified into three categories: standard-hourglass, rotatedhourglass (with its axis perpendicular to outflow), and spiral-like morphology. 40.0% (+-3.0%) of the protostars exhibita mean magnetic field direction approximately perpendicular to the outflow on several 100–1000 au scales. However,in the remaining sample, this relative orientation appears to be random, probably due to the complex set of morphologies observed. Furthermore, we classify the protostars into three types based on the C17O (3–2) velocity envelope’s
gradient: perpendicular to outflow, non-perpendicular to outflow, and unresolved gradient (<1.0 km/s/arcsec). In
protostars with a velocity gradient perpendicular to outflow, the magnetic field lines are preferentially perpendicular
to outflow, most of them exhibit a rotated hourglass morphology, suggesting that the magnetic field has been overwhelmed by gravity and angular momentum. Spiral-like magnetic fields are associated with envelopes having large
velocity gradients, indicating that the rotation motions are strong enough to twist the field lines. All of the protostars
with a standard-hourglass field morphology show no significant velocity gradient due to the strong magnetic braking.

On the magnetic eld properties of protostellar envelopes in Orion
BoHuang,1JosepM.Girart,1,2IanW.Stephens,3ManuelFernándezLópez,4HectorG.Arce,5JohnM.Carpenter,6PauloCortes,7,8 Erin G. Cox,9 Rachel Friesen,10 Valentin J. M. Le Gouellec,11, 12 Charles L. H.
Hull,6, 13 Nicole Karnath,14, 15, 16 Woojin Kwon,17, 18 Zhi-Yun Li,19 Leslie W. Looney,20 Tom Megeath,21 Philip C. Myers,16 Nadia M. Murillo,22, 23 Jaime E. Pineda,24 Sarah Sadavoy,25ÁlvaroSánchez-Monge,1,2PatricioSanhueza,26,27JohnJ.Tobin,7QizhouZhang,16JamesM.Jackson,7andDominiqueSegura-Cox24,28• 観測対象: オリオン大星形成領域（OMCs）に• サンプルのうち2/5は平均的な磁場がアウトフある61個の若い低質量原始星。ALMA870µmのダスト偏光で調査。

ローに対して垂直である一方、残りの3/5はランダムな相対方向を持つ。強い速度勾配を持つ天体では、磁場もアウトフロー軸に対して垂直であることが多い。

• 16個の天体で自己散乱が観測され、その多くはClass 0の段階にあり、ディスク進化の非常に早い段階での粒子成長が示唆される。

• 40個の原始星でエンベロープスケール（約3000 auまで）での磁場をトレース。

• 磁場形態は主に三つに分類される: 砂時計型、

fi

回転砂時計型、スパイラル型。

• 速度勾配と磁場形態の関係: 強い速度勾配を持つ原始星は、回転砂時計型または螺旋型の磁場形態を示す傾向がある。 回転砂時計型の磁場では、磁場の強度は重力や角運動量に比べてそれほど重要でないと考えられる。 螺旋型の構造では、回転運動が磁場を捻じ曲げるのに十分な強さを持つ。

• BOPS（B- eld Orion Protostellar Survey）により、オリオン分子雲内の61個の若い低質量原始星の870 µmダスト偏光を観測。

• 観測フィールド: VLA/ALMA Nascent Disk and Multiplicity (VANDAM)
Surveyによって特定された57の原始星フィールド。観測された原始星の大半
はClass 0で、一部に明るいClass I原始星も含まれる。

fi• 最終的なストークスI画像の平均ノイズレベルは約0.1 mJy/ビーム。

4
Perp-Type: velocity gradient direction ? outflow direction (67.5  |✓Out✓v |  90 )• 典型的な観測例• 左から、870 µmダスト偏Nonperp-Type: velocity gradient direction is not perpendicular to outflow direction (|✓Out✓v | < 67.5 )

光強度、速度場、その拡大図、PV diagram• 上から速度勾配とアウトフ
Unres-Type: velocity gradient is unresolved

ローが垂直な場合、揃っている場合、よく見えなかった場合• グレーは偏光からわかる磁場の平均方向。
Figure 1. Example maps in BOPS survey. First column: 870 µm dust polarization intensity in color scale overlaid with the redshifted andblueshifted outflow lobes (obtained from the 12 CO (3–2) line), polarization segments, and dust continuum emission (Stokes I) contours. Blue• 速度勾配が強い天体は、
7

B-field morphology

NumberVelocity gradient VS outflow direction
Perp-Type (8)Nonperp-Type (6) Unres-Type (25)Standard-hourglass

8

0

0

8

Rotated-hourglass

13

5

1

7

Spiral9243Complex

5

1

0

4

Not enough data

5*

0

1

3

Self-scattering

16

9

3

4

Figure 3. Number of protostars for di↵erent types of B-field morphology. In the first column, green rings indicate protoplanetary envelope (diskfor self-scattering), star symbols located at the center indicate protostars. Red and blue arrows indicate the outflow directions. Black curves andsegments indicate the B-field orientation, while magenta segments indicate polarization direction. HOPS-250, does not have 4 polarization
detection within annular region from 400 to 1000 au scale, is included in the type of Not enough data, which is marked by the asterisk symbol.

between 30 and 250 au (0.00 08 and 0.00 63) (Tobin et al. 2020).
To avoid possible contamination from disk self-scattering polarization (see Section 3.1), we select an annular region from100 (⇠400 au) to 2.00 5 (⇠1000 au). Within this region the di↵erence between the outflow and the intensity-weighted average
B-field direction appears to be very similar to the one obtained using all polarization data (as shown in bottom panels

and Tobin et al. (2011). Further details regarding envelope
kinematics are discussed in Appendix C. Column 7 in Table
3 lists the velocity gradient position angle. In cases withouta significant velocity gradient, we use the spectral resolution
divided by the angular resolution as an upper limit. We use
the velocity gradient position angle ✓v and outflow direction✓Out (as depicted in Figure 1) to classify the protostars into

回転砂時計かスパイラルが多かった。磁場はあまり強くない？

• スパイラルの場合は磁場を十分曲げてそう

OMC3-MMS5
OMC3-MMS6
VLA 1623
IRAS 16293 Aa,c

05:35:22.4
05:35:23.4
16:26:26.4
16:32:22.9–05:01:14.5–05:01:30.6–24:24:30.4–24:28:36.3• メインの結果：40%の天体は磁場とアウトフローが18:29:48.1
18:29:48.7
18:29:49.8
20:39:06.2
21:17:38.7

01:16:43.6
01:16:55.8
01:15:20.3
68:02:16.0
68:17:32.4

垂直、60％はランダムだった。

Ser-emb 8
Ser-emb 8(N)Ser-emb 6
L1157
CB 230

52(9)
45(2)
36(9)
5(9)
5(9)
39(6)
88(7)
157(3)
146(4)89(6)80(6)
171(8)
120(5)
77(9)143(9)
107(1)
129(2)
135(3)
146(7)
172(4)

Notes. Coordinates are fitted positions of dust emission peaks. The outflow angle χo and inferred magfrom north. The angle difference θo−B between the outflow and the B-field is always between 0
perpendicular to the direction of the polarized dust emission. d is the distance to the source. θbm isthe synthesized beam.
a Source has two outflows.
b Coordinates from Takahashi et al. (2009).
c Results from Rao et al. (2009).
d Distance references. (1) Hirota et al. 2011; (2) Loinard et al. 2007; (3) Menten et al. 2007; (4) Lo
2010; (7) Looney et al. 2007; (8) Launhardt et al. 2010.

参考：Hull et al. 2013, TADPOL
survey with CARMA
10probability thatmisaligned pop
is 0.64. Althou
misaligned casewith the data.

Figure 5. Cumulative distribution function of the projected angles. The red polyline is the observed angle di↵erence between the outflow
direction ✓Out and uncertainty-weighted mean B-field direction h✓B i . The black solid line, dashed curve, and dotted-dashed curve indicate
Monte Carlo simulations of the expected projected angle for two vectors that are 3D random (di↵erence angle of two vector is between 0 and90 ), parallel (0 -22. 5), and perpendicular (67. 5-90 ), respectively.

speculate that these could be accretion streamers (e.g., Alveset al. 2020; Pineda et al. 2020; Fernández-López et al. 2023),
however, higher spectral resolution observations are needed

sources with a strong velocity gradient perpendicular
to the outflow axis (Perp-Type) have a B-field that is
also perpendicular to the outflow axis.

Figure 2. Thick solid curve shows the cumulative distribution function (CDF) ofthe (projected) angles between the mean magnetic field and outflow directions
for the sources in Table 1. The upper dashed curve is the CDF from a Monte
Carlo simulation where outflow and B-field directions are oriented within 20◦
of one another (tightly aligned). The lower dot-dashed curve is the CDF froma simulation where outflow and B-field directions are separated by 70◦ –90◦
(preferentially misaligned). The straight line is the CDF for random orientation.

and result in a CDF that is similar to that expected for randomalignment, shown by the solid curve.

Most analyti
rotation axis ofand its outflow d& Pudritz 2000
picture.
We have assu
their long axes
that B-fields arThis may not
are mechanicall
expected to be palignment shoucore, however, ain our sample
volume. Lazaria
that grains alignparallel to the B
in the densest p
Outflows macause of magnethe origin of the first hydrostatic core. This depends weakly on large scale environmental conditions such as radiation,
magnetic field, turbulence or metallicity. This likely constitutes one of the reason of the relative universality of the
IMF.

35. Protoplanetary Disk Polarization at Multiple Wavelengths: Are Dust Populations Diverse?

Rachel E. Harrison, Zhe-Yu Daniel Lin, Leslie W. Looney, Zhi-Yun Li, Haifeng Yang, Ian Stephens, ManuelFernández-López F Millimeter and sub-millimeter observations of continuum linear dust polarization provide insight

into dust grain growth in protoplanetary disks, which are the progenitors of planetary systems. We present the resultsof the first survey of dust polarization in protoplanetary disks at 870 µm and 3 mm. We find that protoplanetarydisks in the same molecular cloud at similar evolutionary stages can exhibit diﬀerent correlations between observing
wavelength and polarization morphology and fraction. We explore possible origins for these diﬀerences in polarization,

including diﬀerences in dust populations and protostar properties. For RY Tau and MWC 480, which are consistent
with scattering at both wavelengths, we present models of the scattering polarization from several dust grain size
13 of the observational results for these disks: (1) both disks
distributions. These models aim to reproduce two features
have an observable degree of polarization at both wavelengths and (2) the polarization fraction is higher at 3 mm thanat 870 µm in the centers of the disks. For both disks, these features can be reproduced by a power-law distribution of
spherical dust grains with a maximum radius of 200 µm and high optical depth. In MWC 480, we can also reproduce
features (1) and (2) with a model containing large grains (amax = 490 µm ) near the disk midplane and small grains(amax = 140 µm) above and below the midplane.

Ineﬃcient star formation in high Mach number environments I. The turbulent support
36. analytical model
Patrick Hennebelle, Noé Brucy, Tine Colman F The star formation rate (SFR), the number of stars formed per unitof time, is a fundamental quantity in the evolution of the Universe. While turbulence is believed to play a crucial roleProtoplanetary Disk Polarization at Multiple Wavelengths: Are Dust Populations Diverse?
Rachel E. Harrison,1, 2 Zhe-Yu Daniel Lin,3 Leslie W. Looney,1 Zhi-Yun Li,3 Haifeng Yang,4 Ian Stephens,5 and Manuel Fernández-López6

Multi-wavelength Dust Polarization

7• 赤が870µm偏光、青が3mmの偏光• ２天体では方向が一致→この二波長
(a)

(b)

で自己散乱が卓越• モデリングは、140, 490ミクロン
c)

(d)

e)

Figure 2. Images of Haro 6-13, MWC 480, RY Tau, V892 Tau, and DL Tau showing the polarization angles at 870 µm (redlines) and 3 mm (blue lines). The black contours represent Stokes I at levels of -3 (dashed), 3, 10, 50, 100, 250, 325, 500, 750,
1000, and 1500 , where is the RMS of the smoothed 870 µm images. Vectors are plotted in regions where the total intensityand polarized intensity are both > 3 . The length of the vectors is scaled with percent polarization up to a threshold of 3% tobetter show the variation at low polarization fraction. The blue ellipse represents the beam size.

の２つのダストモデルの組み合わせで説明した。

Figure 3. Polarization P at the scattering angle of 90° and the albedo w = ksca (kabs + ksca ) as a function of maximum grain size. The size distribution is assumed tohave a power law of n (a ) µ a-3.5 . The wavelengths are assumed to be 3.1 mm for the left panel and 870 m m for the right panel. The arrows indicate the maximum
grain size that has the most efﬁcient polarization by 90° scattering.

• ここは片岡の補足スライド• コンパクトな場合、a̲maxを決めて-3.5

乗則でダストを考えると、偏光を出す波
The second condition for producing the polarization due toscattering is the asymmetry in the distribution of the light
長レンジが極めて狭い
source. In the case of sub-millimeteremission from proto2.5. Anisotropic Radiation Field

planetary disks, the light source is the thermal emission of dust
particles. Therefore, its spatial distribution should be anisotropic to produce polarized scattered light. For example, if theincident radiation from one direction is stronger than that fromthe direction different by 90 , the ﬁnal scattering is partiallypolarized. Such asymmetry or anisotropy may occur inprotoplanetary disks showing ring-like or lopsided surface
brightness (e.g., Casassus et al. 2013; Fukagawa et al. 2013;
Isella et al. 2013; van der Marel et al. 2013; Pérez et al. 2014).
In this section, we demonstrate that polarization due to selfZhang et al.
scattering
occurs in such anisotropic radiation ﬁelds by using asimple toy model. Note that the polarization due to scattering of
the anisotropic radiation ﬁeld has been discussed in the context
of E-mode polarization of the cosmic microwave backgroundradiation (e.g., Rees 1968; Kamionkowski et al. 1997;
Seljak 1997).
Hereafter, we ﬁx the maximum grain size and wavelengths
to be a max = 100 m m and l = 870 m m , which is one of thebest combinations of the efﬁcient polarization, in order toinvestigate possibilities of detecting millimeter-wave polarization from protoplanetary disks. Note that the calculated
absorption and scattering opacities are kabs = 0.51 andksca = 1.22 cm2 g-1. We stress here that we consider the same

• よって870µmと3mmの波長の療法で偏光を出すためには２つの分布を考える。
が、それはもはやpower-lawではない。

Figure 4. Polarization degree times the albedo Pw against the maximum grain
size. This ﬁgure represents a grain size that contributes most to the polarized
intensity. Each line corresponds to the wavelengths of 0.34, 0.87, 3.1, and
7 mm. The band numbers correspond to the ALMA band numbers for eachwavelength.

The Astrophysical Journal, 953:96 (31pp), 2023 August 10• porous aggregatesを考えると、比較的自然に説明できる(Zhang et al.
2023 for HL Tau, Ueda et al. 2024
for IM Lup)

Table 1
The Detectable Grain Size for Observed Wavelengths
Wavelengths λ

7 mm
3.1 mm
870 m m
340 m mThe Detectable Grain Size a max

1 mm
500 m m
150 m m
70 m m3• とすると、天体に依ってporosityが違うことを示唆しているのか？（M1北出くんがこれから明らかにしてくれる。）

Figure 3. From top to bottom: the size (a max f ) dependence of the κabs (absorption opacity), Iν/Bν (intensity reduction due to scattering compared to the blackbody
radiation), αthin, αthick (spectral indices at optically thin and thick regimes), and Pωeff (the product of the polarization fraction at 90°, P and the effective albedo, ωeff).
In the bottom panels, solid lines are Pωeff, dotted lines are ωeff, and dashed–dotted lines are P. From left to right, the ﬁlling factors are f = 1, 0.3, 0.1, and 0.01. Red,

crucial for quantifying their impact on the protostellar systems.

32. Maser Flares Driven by Isothermal Shock Waves

M. D. Gray, S. Etoka, B. Pimpanuwat, A. M. S. Richards F We use 3D computer modelling to investigate thetimescales and radiative output from maser flares generated by the impact of shock-waves on astronomical unit-scale
clouds in interstellar and star-forming regions, and in circumstellar regions in some circumstances. Physical conditions
are derived from simple models of isothermal hydrodynamic (single-fluid) and C-type (ionic and neutral fluid) shockwaves, and based on the ortho-H2 O 22-GHz transition. Maser saturation is comprehensively included, and we findthat the most saturated maser inversions are found predominantly in the shocked material. We study the eﬀect on theintensity, flux density and duration of flares of the following parameters: the pre-shock level of saturation, the observer’s
viewpoint, and the shock speed. Our models are able to reproduce observed flare rise times of a few times 10 days,
specific intensities of up to 105 times the saturation intensity and flux densities of order 100(R/d)2 Jy from a source ofradius R astronomical units at a distance of d kiloparsec. 12
We found that flares from C-type shocks are approximately
5 times more likely to be seen by a randomly placed observer than flares from hydrodynamically shocked clouds ofsimilar dimensions. We computed intrinsic beaming patterns of the maser emission, finding substantial extension of
the pattern parallel to the shock front in the hydrodynamic models. Beaming solid angles for hydrodynamic models
can be as small as 1.3 ⇥ 10 5 sr, but are an order of magnitude larger for C-type models.

33. Radial and vertical constraints on the icy origin of H2 CO in the HD 163296 Protoplanetary Disk

Claudio Hernández-Vera, Viviana V. Guzmán, Elizabeth Artur de la Villarmois, Karin I. Öberg, L. Ilsedore
Cleeves, Michiel R. Hogerheijde, Chunhua Qi, John Carpenter, Edith C. Fayolle F H2 CO is a small organic

molecule widely detected in protoplanetary disks. As a precursor to grain-surface formation of CH3 OH, H2 CO isconsidered an important precursor of O-bearing organic molecules that are locked in ices. Still, since gas-phase reactions can also form H2 CO, there remains an open question on the channels by which organics form in disks, andsimilar dimensions. We computed intrinsic beaming patterns of the maser emission, finding substantial extension of
the pattern parallel to the shock front in the hydrodynamic models. Beaming solid angles for hydrodynamic models
can be as small as 1.3 ⇥ 10 5 sr, but are an order of magnitude larger for C-type models.

33. Radial and vertical constraints on the icy origin of H2 CO in the HD 163296 Protoplanetary Disk

Claudio Hernández-Vera, Viviana V. Guzmán, Elizabeth Artur de la Villarmois, Karin I. Öberg, L. Ilsedore
Cleeves, Michiel R. Hogerheijde, Chunhua Qi, John Carpenter, Edith C. Fayolle F H2 CO is a small organic

molecule widely detected in protoplanetary disks. As a precursor to grain-surface formation of CH3 OH, H2 CO isconsidered an important precursor of O-bearing organic molecules that are locked in ices. Still, since gas-phase reactions can also form H2 CO, there remains an open question on the channels by which organics form in disks, andhow much the grain versus the gas pathways impact the overall organic reservoir. We present spectrally and spatiallyresolved Atacama Large Millimeter/submillimeter Array observations of several ortho- and para-H2 CO transitions
toward the bright protoplanetary disk around the Herbig Ae star HD 163296. We derive column density, excitation
12
temperature, and ortho-to-para ratio (OPR) radial profiles for H2 CO, as well as disk-averaged values of NT ⇠ 4 ⇥ 10
cm 2 , Tex ⇠ 20 K, and OPR ⇠ 2.7, respectively. We empirically determine the vertical structure of the emission,
finding vertical heights of z/r ⇠ 0.1. From the profiles, we find a relatively constant OPR ⇠ 2.7 with radius, but still
consistent with 3.0 among the uncertainties, a secondary increase of NT in the outer disk, and low Tex values that
decrease with disk radius. Our resulting radial, vertical, and OPR constraints suggest an increased UV penetration
beyond the dust millimeter edge, consistent with an icy origin but also with cold gas-phase chemistry. This Herbig
disk contrasts previous results for the T Tauri disk, TW Hya, which had a larger contribution from cold gas-phasechemistry. More observations of other sources are needed to disentangle the dominant formation pathway of H2 CO inprotoplanetary disks.

34. The Physical Origin of the Stellar Initial Mass Function

Patrick Hennebelle, Michael Y. Grudić F Stars are amongst the most fundamental structures of our Universe.

34. The Physical Origin of the Stellar Initial Mass Function

Patrick Hennebelle, Michael Y. Grudić F Stars are amongst the most fundamental structures of our Universe.

They comprise most of the baryonic and luminous mass of galaxies, synthethise heavy elements, and injec t mass,
momentum, and energy into the interstellar medium. They are also home to the planets. Since stellar properties areprimarily decided by their mass, the so-called stellar initial mass function (IMF) is critical to the structuring of ourUniverse. We review the various physical processes, and theories which have been put forward as well as the numericalsimulations which have been carried out to explain the origin of the stellar initial mass function. Key messages fromthis review are: (1) Gravity and turbulence most likely determine the power-law, high-mass part of the IMF. (2)Depending of the Mach number and the density distribution, several regimes are possible, including IM F ' 0, -0.8, -1
or -1.3 where dN/d log M / M IM F . These regimes are likely universal, however the transition between these regimes
is not. (3) Protostellar jets can play a regulating influence on the IMF by injecting momentum into collapsing clumps
and unbinding gas. (4) The peak of the IMF may be a consequence of dust opacity and molecular hydrogen physics atthe origin of the first hydrostatic core. This depends weakly on large scale environmental conditions such as radiation,
magnetic field, turbulence or metallicity. This likely constitutes one of the reason of the relative universality of the
IMF.

35. Protoplanetary Disk Polarization at Multiple Wavelengths: Are Dust Populations Diverse?

Rachel E. Harrison, Zhe-Yu Daniel Lin, Leslie W. Looney, Zhi-Yun Li, Haifeng Yang, Ian Stephens, ManuelFernández-López F Millimeter and sub-millimeter observations of continuum linear dust polarization provide insight

into dust grain growth in protoplanetary disks, which are the progenitors of planetary systems. We present the resultsof the first survey of dust polarization in protoplanetary disks at 870 µm and 3 mm. We find that protoplanetarydisks in the same molecular cloud at similar evolutionary stages can exhibit diﬀerent correlations between observing
wavelength and polarization morphology and fraction. We explore possible origins for these diﬀerences in polarization,

at 870 µm in the centers of the disks. For both disks, these features can be reproduced by a power-law distribution of
spherical dust grains with a maximum radius of 200 µm and high optical depth. In MWC 480, we can also reproduce
features (1) and (2) with a model containing large grains (amax = 490 µm ) near the disk midplane and small grains(amax = 140 µm) above and below the midplane.

Ineﬃcient star formation in high Mach number environments I. The turbulent support
36. analytical model
Patrick Hennebelle, Noé Brucy, Tine Colman F The star formation rate (SFR), the number of stars formed per unitof time, is a fundamental quantity in the evolution of the Universe. While turbulence is believed to play a crucial rolein setting the SFR, the exact mechanism remains unclear. Turbulence promotes star formation by compressing the gas,
but also slows it down by stabilizing the gas against gravity. Most widely-used analytical models rely on questionableassumptions, including: i) integrating over the density PDF, a one-point statistical description that ignores spatial
correlation, ii) selecting self-gravitating gas based on a density threshold that often ignores turbulent dispersion, iii)
assuming the freefall time as the timescale for estimating SFR without considering the need to rejuvenate the density
PDF, iv) assuming the density PDF to be lognormal. Improving upon the only existing model that incorporates thespatial correlation of the density field, we present a new analytical model. We calculate the time needed to rejuvenate
density fluctuations of a given density and spatial scale, revealing that it is generally much longer than the freefall time,rendering the latter inappropriate for use. We make specific predictions regarding the role of the Mach number, M ,and the driving scale of turbulence divided by the mean Jeans length. At low to moderate Mach numbers, turbulence
does not reduce and may even slightly promote star formation by broadening the PDF. However, at higher Mach
numbers, most density fluctuations are stabilized by turbulent dispersion, leading to a steep drop in the SFR as theMach number increases.

37. On the magnetic field properties of protostellar envelopes in Orion

Bo Huang, Josep M. Girart, Ian W. Stephens, Manuel Fernandez-Lopez, Hector G. Arce, John M. Carpenter,of the radial magnetic field owing to the ambipolar diﬀusion and a steep radial gradient of the Ohmic diﬀusivity causes
the eﬃcient evacuation of the net vertical magnetic flux from the transition zone within several rotations. Surface gas
accretion occurs in the coherent zone but not in the transition zone. The presence of the transition zone prohibits
mass and magnetic flux transport from the coherent zone to the active zone. Mass accumulation occurs at both edgesof the transition zone as a result of mass supply from the active and coherent zones.

39. Meteorites and Planet Formation

Rhian H. Jones F Meteorites are a remarkable resource. They capture the imagination of people worldwide with theirspectacular entry through Earth’s atmosphere as fireballs, and their exotic character of being pieces of other worlds.
Scientifically, they are critical to interpreting the early stages of formation of the Solar System, as well as the geologicalevolution of asteroids, the Moon, and Mars, and they are vital to understanding planetary formation processes. With
the burgeoning exploration of extrasolar planetary systems, knowledge of the fundamental process of planetary growthfrom protoplanetary disks has taken on a new significance. Meteorites provide essential and detailed insight into theformation of planetary systems, although we must bear in mind that they only represent one reference point (our own
Solar System) in what is clearly a wide spectrum of possible chemical and physical parameters governing the diverse
realm of extrasolar planets. This chapter summarises the nature of our meteorite collections, and the ways in which
meteorites contribute to our understanding of the formation and evolution of our own Solar System, with broader
implications for planetary systems in general.

The
CO-to-H
Conversion
FactorintheCentral
Molecular
Zone
of
the
Milky
Way
using
2
40.
CO isotopologues

Mikito Kohno, Yoshiaki Sofue F We performed correlation analyses between the 12 CO and 13 CO J =1-0 line

intensities in order to derive the variability of the CO-to-H2 conversion factor (XCO,iso ) in the central molecular zone
(CMZ) of our Galaxy. New high-resolution XCO,iso maps at a resolution of ⇠ 30" and the longitude-velocity diagram(LVD) at resolution ⇠ 30" ⇥ 2 km s 1 are presented using the 12 CO and 13 CO archival survey data obtained by
the Nobeyama 45 m telescope. We revealed the variation of XCO,iso in the CMZ within the range of XCO,iso ⇠Solar System) in what is clearly a wide spectrum of possible chemical and physical parameters governing the diverse
realm of extrasolar planets. This chapter summarises the nature of our meteorite collections, and the ways in which
meteorites contribute to our understanding of the formation and evolution of our own Solar System, with broader
implications for planetary systems in general.

The
CO-to-H
Conversion
FactorintheCentral
Molecular
Zone
of
the
Milky
Way
using
2
40.
CO isotopologues

Mikito Kohno, Yoshiaki Sofue F We performed correlation analyses between the 12 CO and 13 CO J =1-0 line

intensities in order to derive the variability of the CO-to-H2 conversion factor (XCO,iso ) in the central molecular zone
(CMZ) of our Galaxy. New high-resolution XCO,iso maps at a resolution of ⇠ 30" and the longitude-velocity diagram(LVD) at resolution ⇠ 30" ⇥ 2 km s 1 are presented using the 12 CO and 13 CO archival survey data obtained by
the Nobeyama 45 m telescope. We revealed the variation of XCO,iso in the CMZ within the range of XCO,iso ⇠(0.2 1.3) ⇥ 1020 cm 2 (K km s 1 ) 1 , if we assume the normalization value of 0.59 ⇥ 1020 cm 2 (K km s 1 ) 1 . The
mean value is obtained as XCO,iso = (0.48 ± 0.15) ⇥ 1020 cm 2 (K km s 1 ) 1 in the CMZ of our Galaxy.

41. Scales of Stability and Turbulence in the Molecular ISM

Eric Keto F From the archival data of the BU-FCRAO 13 CO GRS, we measure the radial profiles of column density

and turbulent velocity dispersion from the centers of molecular clouds outward into the surrounding diﬀuse molecular
ISM. Averaged across spatial scales, the profiles are consistent with turbulence that is on average in hydrostatic
equilibrium. We measure the turbulent kinetic energy (KE) and the gravitational potential energy (GE) within
and the pressure energy (PE) outside the clouds. The distribution of the sum 2KE-|GE|-PE has a mean near zero
indicating approximate virial equilibrium. The average pressure energy is consistent with estimates of the pressure of
the multiphase ISM at the Galactic mid-plane. If the dynamical time scale of the turbulence scales with its crossing
time, this apparent equilibrium may result from the rapid virialization of the KE and GE with respect to changes in
PE. In a snapshot, the clouds appear to be in virial equilibrium within a confining pressure. However, the durationof the equilibrium is approximately the crossing time. An analysis of Larson’s scaling relationships of line width andmean value is obtained as XCO,iso = (0.48 ± 0.15) ⇥ 1020 cm 2 (K km s 1 ) 1 in the CMZ of our Galaxy.

41. Scales of Stability and Turbulence in the Molecular ISM

Eric Keto F From the archival data of the BU-FCRAO 13 CO GRS, we measure the radial profiles of column density

and turbulent velocity dispersion from the centers of molecular clouds outward into the surrounding diﬀuse molecular
ISM. Averaged across spatial scales, the profiles are consistent with turbulence that is on average in hydrostatic
equilibrium. We measure the turbulent kinetic energy (KE) and the gravitational potential energy (GE) within
and the pressure energy (PE) outside the clouds. The distribution of the sum 2KE-|GE|-PE has a mean near zero
indicating approximate virial equilibrium. The average pressure energy is consistent with estimates of the pressure of
the multiphase ISM at the Galactic mid-plane. If the dynamical time scale of the turbulence scales with its crossing
time, this apparent equilibrium may result from the rapid virialization of the KE and GE with respect to changes in
PE. In a snapshot, the clouds appear to be in virial equilibrium within a confining pressure. However, the durationof the equilibrium is approximately the crossing time. An analysis of Larson’s scaling relationships of line width andcolumn density finds no correlation. Tthe inference of constant column density with cloud size is a misinterpretation

15

42. Kinematic age of the -Pictoris moving group

Jinhee Lee, Inseok Song F Accurate age estimation of nearby young moving groups (NYMGs) is important asthey serve as crucial testbeds in various fields of astrophysics, including formation and evolution of stars, planets,
as well as loose stellar associations. The -Pictoris moving group (BPMG), being one of the closest and youngest
NYMGs, has been extensively investigated, and its estimated ages have a wide range from ⇠10 to 25 Myr, dependingon the age estimation methods and data used. Unlike other age dating methods, kinematic traceback analysis oﬀers
a model-independent age assessment hence the merit in comparing many seemingly discordant age estimates. In thisstudy, we determine the kinematic ages of the BPMG using three methods: probabilistic volume calculation, mean
pairwise distance calculation, and covariance matrix analysis. These methods yield consistent results, with estimatedages in the range of 14 to 20 Myr. Implementing corrections to radial velocities due to gravitational redshift andconvectional blueshift increases the ages by ⇠ 2 4 Myr. Conversely, considering data uncertainties decreases theestimated ages by 1 to 2 Myr. Taken together, our analysis determined the kinematic age of BPMG to be 16.3+3.4
2.1
Myr. This age is significantly younger than the commonly accepted age of the BPMG (⇠24 Myr) determined primarily
from the lithium depletion boundary analysis. This younger kinematic age may point to the discrepancy between theluminosity evolution and lithium depletion models or the presence of unaccounted systematic error in the method.
This result underscores the necessity for systematic reevaluations of age-dating methods for nearby, young moving
groups.

43. Formation of the four terrestrial planets in the Jupiter-Saturn chaotic excitation scenario: fundamental properties and water delivery

Patryk Sofia Lykawka, Takashi Ito F The Jupiter-Saturn chaotic excitation (JSCE) scenario proposes that the

protoplanetary disk was dynamically excited and depleted beyond 1-1.5 au in a few Myr, oﬀering a new and plausible explanation for several observed properties of the inner solar system. Here, we expanded our previous work byconducting a comprehensive analysis of 37 optimal terrestrial planet systems obtained in the context of the JSCE
scenario. Each optimal system harbored exactly four terrestrial planets analogs to Mercury, Venus, Earth, and Mars.

groups.

43. Formation of the four terrestrial planets in the Jupiter-Saturn chaotic excitation scenario: fundamental properties and water delivery

Patryk Sofia Lykawka, Takashi Ito F The Jupiter-Saturn chaotic excitation (JSCE) scenario proposes that the

protoplanetary disk was dynamically excited and depleted beyond 1-1.5 au in a few Myr, oﬀering a new and plausible explanation for several observed properties of the inner solar system. Here, we expanded our previous work byconducting a comprehensive analysis of 37 optimal terrestrial planet systems obtained in the context of the JSCE
scenario. Each optimal system harbored exactly four terrestrial planets analogs to Mercury, Venus, Earth, and Mars.
We further investigated water delivery, feeding zones, and accretion history for the planet analogs, which allowed us tobetter constrain the water distribution in the disk. The main findings of this work are as follows: 1) the formation offour terrestrial planets with orbits and masses similar to those observed in our solar system in most of our sample, as
evidenced by the dynamically colder and hotter orbits of Venus-Earth and Mercury-Mars analogs, and the high successrates of similar mutual orbital separations ( 40-85%) and mass ratios of the planets ( 70-90%) among the 37 systems;
and 2) water was delivered to all terrestrial planets during their formation through the accretion of water-bearing
disk objects from beyond 1-1.5 au. The achievement of Earth’s estimated bulk water content required the disk tocontain suﬃcient water mass distributed within those objects initially. This requirement implies that Mercury, Venus,
and Mars acquired water similar to the amount on Earth during their formation. Several of our planet analogs also
matched additional constraints, such as the timing of Moon formation by a giant impact, Earth’s late accretion massand composition, and Mars’s formation timescale.

44. Disk dissipation, giant planet formation and star-formation-rate fluctuations in the past
three-million-year history of Gould’s Belt

Mingchao Liu, Jinhua He, Jixing Ge, Tie Liu, Yuping Tang, Xuzhi Li F Although episodic star formation (SF) had

been suggested for nearby SF regions, a panoramic view to the latest episodic SF history in the solar neighborhoodand Mars acquired water similar to the amount on Earth during their formation. Several of our planet analogs also
matched additional constraints, such as the timing of Moon formation by a giant impact, Earth’s late accretion massand composition, and Mars’s formation timescale.

44. Disk dissipation, giant planet formation and star-formation-rate fluctuations in the past
three-million-year history of Gould’s Belt

Mingchao Liu, Jinhua He, Jixing Ge, Tie Liu, Yuping Tang, Xuzhi Li F Although episodic star formation (SF) had

been suggested for nearby SF regions, a panoramic view to the latest episodic SF history in the solar neighborhood
is still missing. By uniformly constraining the slope ↵ of infrared spectral energy distributions (SEDs) of youngstellar objects (YSOs) in the 13 largest Gould’s Belt (GB) protoclusters surveyed by Spitzer Space Telescope, wehave constructed a cluster-averaged histogram of ↵ representing YSO evolution lifetime as a function of the ↵ value.
Complementary to the traditional SED classification scheme (0, I, F, II, III) that is based on diﬀerent ↵ values, astaging scheme (A,B,C,D,E) of SED evolution is advised on the basis of the ↵ statistical features that can be better
matched to the physical stages of disk dissipation and giant planet formation. This has also allowed us to unravel
the fluctuations of star formation rate (SFR) in the past three-million-year (3 Myr) history of these GB protoclusters.
Diverse evolutionary patterns such as single peak, double peaks and on-going acceleration of SFR are revealed. The
SFR fluctuations are between 20% ⇠ 60% (⇠ 40% on average) and no dependence on the average SFR or the number of
SFR episodes is found. However, spatially close protoclusters tend to share similar SFR fluctuation trends, indicatingthat the driving force of the fluctuations should be at size scales beyond the typical cluster sizes of several parsec.
16

Can near-to-mid infrared spectral energy distribution quantitatively trace protoplane45. tary disk evolution?

Mingchao Liu, Jinhua He, Zhen Guo, Jixing Ge, Yuping Tang F Infrared (IR) spectral energy distribution (SED)

is the major tracer of protoplanetary disks. It was recently proposed to use the near-to-mid IR (or K-24) SED slope↵ defined between 2-24µm as a potential quantitative tracer of disk age. We critically examine the viability of this

SFR episodes is found. However, spatially close protoclusters tend to share similar SFR fluctuation trends, indicatingthat the driving force of the fluctuations should be at size scales beyond the typical cluster sizes of several parsec.

Can near-to-mid infrared spectral energy distribution quantitatively trace protoplane45. tary disk evolution?

Mingchao Liu, Jinhua He, Zhen Guo, Jixing Ge, Yuping Tang F Infrared (IR) spectral energy distribution (SED)

is the major tracer of protoplanetary disks. It was recently proposed to use the near-to-mid IR (or K-24) SED slope↵ defined between 2-24µm as a potential quantitative tracer of disk age. We critically examine the viability of thisidea and confront it with additional statistics of IR luminosities and SED shapes. We point out that, because thestatistical properties of most of the complicated physical factors involved in disk evolution are still poorly understoodin a quantitative sense, the only viable way is to assume them to be random so that an idealized ‘average disk’ can be
defined, which allows the ↵ histogram to trace its age. We confirm that the statistics of the zeroth order (luminosity),
first order (slope ↵) and second order characteristics (concavity) of the observed K-24 SEDs indeed carry useful information upon the evolutionary processes of the ‘average disk’. We also stress that intrinsic diversities in K-24 SED
shapes and luminosities are always large at the level of individual stars so that the application of the evolutionary
path of the ‘average disk’ to individual stars must be done with care. The data of most curves in plots are provided
on GitHub.

New
Herbig-Haro
Objects
associatedwithEmbedded
Sources46.

T. A. Movsessian, T. Yu. Magakian, B. Reipurth, H. R. Andreasyan F We continue to present the results of the

Byurakan Narrow Band Imaging Survey (BNBIS). The main goal of this survey is to search for Herbig-Haro (HH)
objects and jets in Galactic dark clouds. In this work we present the results of the search in the vicinity of infrared
sources that are bright in the WISE survey and embedded in the dark clouds. The survey is performed with the
1 m Schmidt telescope of Byurakan Observatory, lately equipped with a new CCD detector, which allows to obtain