星形成ゼミ 2020/9/11 片岡章雅（国立天文台）
36. “A planet within the debris disk around the pre-main-sequence star AU Microscopii” Peter
Plavchan et al.
37.“Variability of the Great Disk Shadow in Serpens” Klaus M. Pontoppidan et al.
38. “The circumstellar environment of EX Lup: SPHERE and SINFONI vimews” Elisabetta
Rigliaco etal.
39.Ring formation and dust dynamics in wind-driven protoplanetary discs: global
simulations A. Riols et al.
40. Characterizing the radio continuum nature of sources in the massive star-forming region
W75N (B) , Rodríguez-Kamenetzky et al.
41. The Role of Outﬂows, Radiation Pressure, and Magnetic Fields in Massive Star Formation,
Rosen and Krumholz
42. Planet Migration in Self-Gravitating Discs: Survival of Planets, Rowther and Meru

http://arxiv.org/pdf/2004.14122
36

Table 1 | Star parameters

Table 2 | Planetary parameters

Parameter

Parameter

A planet within the debris disk around the pre-main-sequence star AU Microscopii
Peter Plavchan , Thomas Barclay
12,5

68% credible interval Remarks, Jonathan Gagné , Peter GaoAU, Mic
BrysonCale et al.
(star)
3

4

1

68% credible interval

Remarks

Period

8.46321 ± 0.00004 days

TESS and Spitzer
transit light curve
analysisSemi-major axis

0.066+0.007−0.006 AU

TESS and Spitzer
transit light curve
analysis

Velocity
semi-amplitude, K<28 m s−1

RV analysis

Mass<3.4MNeptune <0.18MJupiter

RV analysis

Radius(1.08 ± 0.05)RNeptune(0.375 ± 0.018)RJupiter

TESS and Spitzer
transit light curve

Density<4.4 g cm−3

RV / TESS analysis

Time(s) ofconjunctiona2, 458, 330.39153+0.00070−0.00068BJDTESS and Spitzer
transit light curves

Transit duration, τ14

3.50+0.63−0.59hTESS and Spitzer
transit light curves

Rp/R*

0.0514 ± 0.0013

TESS and Spitzer
transit light curve

Impact parameter, b

0.16+0.14−0.11

TESS and Spitzer
transit light curve

a/R*

19.1+1.8−1.6

TESS and Spitzer
transit light curve

Eccentricity

0.10+0.17−0.09

TESS and Spitzer
transit light curveb.

AU Mic b

Distance from the Sun 2
9.79 ± 0.04 pc
Gaia mission parallax
Department of Physics and Astronomy, George Mason University, Fairfax,
VA, USA; Center
for Space
Sciences
Radius(0.75 ± 0.03)R☼
Directly measured with3
and Technology, University of Maryland Baltimore County (UMBC), Baltimore, MD, USA; Institute forinterferometry
Research
4
Mass(0.50 ± 0.03)M☼
Estimated fromof
on Exoplanets, Département de Physique, Université de Montréal, Montréal,
Quebec, Canada;
Department
spectral type and age
5
Astronomy, University of California, Berkeley, CA, USA; Exoplanets andT Stellar Astrophysics
Laboratory,NASA
3,700 ± 100
KSpectralenergydistribution modelling
Goddard Space Flight Center, Greenbelt, MD, USA

1

17

a

eff

E-mail contact: plavchan at gmail.com

15

Luminosity

0.09L☼

Spectral energy
distribution modelling15Age

22 ± 3 Myr

Ref. 1

period of 9.79 parsecs
4.863 ± 0.010 anddays with
RV analysis,
AU Microscopii (AU Mic) is the second closest pre main sequence star, at aRotation
distance
an TESSagelight
curve, SuperWASP
of 22 million years. AU Mic possesses a relatively rare and spatially resolved edge-on debris disk extending fromabout
light curve
Projected rotational velocity
8.7 ± 0.2 kms
Ref.of newly
35 to 210 astronomical units from the star, and with clumps exhibiting non-Keplerianmotion.
Detection
limb-darkening
TESS light curve
0.21
formed planets around such a star is challenged by the presence of spots, Linear
plage,
flaresandothermanifestationsofcoefficient (TESS)
magnetic activity on the star. Here we report observations of a planet transiting
Mic. The
AU
QuadraticAU
limb-darkening
TESS light curve
0.0 transiting planet,coefficient (TESS)
Mic b, has an orbital period of 8.46 days, an orbital distance of 0.07 astronomical
units, a radius of 0.4 Jupiterradii,
Linear limb-darkening
Spitzer
light
curve
0.17
coefficientand a mass of less than 0.18 Jupiter masses at 3σ confidence. Our observationsof(Spitzer)a planet co-existing with a debris
limb-darkening
Spitzer light curve
0.15
disk oﬀer the opportunity to test the predictions of current models of planetQuadratic
formation
andevolution.
coefficient (Spitzer)
18

−1

12

+0.20
−0.15

+0.18
−0.14

0.22
−0.12

+0.27
−0.21

Accepted by Nature

Visible stellar activity
amplitude−1
145+17−14 msRV analysis

http://arxiv.org/pdf/2006.13248

Near-infrared stellar activity
amplitude−1
80+16−12 msRV analysis; K band at
2.3 µmSpot decay half-life

110 ± 30 days

RV analysis

GP hyper-parameter 4
0.37 ± 0.02
RV analysis
・AU Micという天体の周りでトランジット惑星を発見した、という論文
Variability of the Great Disk Shadow in Serpens
Apparent magnitude
TESS = 6.76 mag
TESS light curve
・AU Micという天体について
Klaus M. Pontoppidan1 , Joel D. Green1 , Tyler A. Pauly1 , Colette Salyk2 and Joseph DePasquale1
・AU Micは距離9.79pcと地球に近い。
1
Space Telescope Science Institute 3700 San Martin Drive Baltimore, MD 21218, USA; 2 Vassar College 124 Raymond
an important test of migration models since we expect any obliquity in・AU Micはデブリ円盤を持っている。デブリ円盤の大きさは35-210au。
Ave Poughkeepsie,
NY 12604, USA
this young system to be unaffected by stellar tides and thus primordial.
AU Mic is a member of the β Pictoris Moving Group; the group’s archeE-mail・今回の結果
contact: pontoppi at stsci.edu
type β Pic is a much more massive (about 3.5×), luminous (about 100×)and hotterin(approximately2×) A-typestar, also possessing
a debris
We present
multi-epoch Hubble Space Telescope imaging of the Great Disk Shadow
the Serpens
star-forming
region.
・TESSを使った観測
disk. β Pic has a more massive Jovian planet β Pic b observed by direct
Also consistent with independently fitting the two transit events in TESS light curve for AU
Mic b.
a

The near-infrared images show strong variability of the disk shadow, revealing
of of
the
inner
disk
on
time
imagingdynamics
at a semi-major axis
about
9 AU, witha massof approximately
・軌道長半径0.07au、0.4木星半径と見積もられている。
with astrometryprotoplanetary. AU Mic and β Pic are of thescales of months.
The Great Shadow is projected onto the Serpens reflection(11±2)M
nebula determinedby an unresolved
same stellar age, but are very different exoplanet host stars. While AU
disk surrounding
the young intermediate-mass star SVS2/CK3/EC82. SinceMic
theshadowextendsout
totoaβ Pic
distanceof・デブリ円盤と惑星が同時にあるという天体の発見がすごい(とのこと)。
b possibly
formedat a distancesimilarb and then migrated
inwards to
its present
location,β Pic b has
not substantiallymigrated
at least 17,000 au, corresponding to a light travel time of 0.24 years, the images
mayrevealdetailedchangesin theJupiter

14

Candidate transit event
Period30 ± 6 days

TESS light curvetransit duration

Radius(0.60 ± 0.17)RNeptune =(0.21 ± 0.06)RJupiter

TESS transit lightcurve

Time(s) ofconjunction

2,458,342.22 ± 0.03 days

TESS transit light
curve

Rp/R*

0.028 ± 0.006

TESS transit lightcurve

Impact parameter, b

0.5 ± 0.3

TESS transit light
curve

a/R*

40 ± 8TESS transit lightcurve

Eccentricity

0.2 ± 0.2

TESS transit light
curve

Barycentric Julian Day.
Circular orbit assumed for RV analysis.

abof being ground down further by dust collisions until blown out of
the system by radiation pressure. For 1-µm-sized solid grains of dusty

Variability of the Great Disk Shadow in Serpens

1
1
1
2
1
Klaus
M.
Pontoppidan
,
Joel
D.
Green
,
Tyler
A.
Pauly
,
ColetteSalyk
andJosephDePasquale
37
1
Space Telescope Science Institute 3700 San Martin Drive Baltimore, MD 21218, USA; 2 Vassar College 124 Raymond
Ave Poughkeepsie, NY 12604, USA

E-mail contact: pontoppi at stsci.edu
We present multi-epoch Hubble Space Telescope imaging of the Great Disk Shadow in the Serpens star-forming region.
The near-infrared images show strong variability of the disk shadow, revealing dynamics of the inner disk on timescales of months. The Great Shadow is projected onto the Serpens reflection nebula by an unresolved protoplanetary
disk surrounding the young intermediate-mass star SVS2/CK3/EC82. Since the shadow extends out to a distance ofat least 17,000 au, corresponding to a light travel time of 0.24 years, the images may reveal detailed changes in thedisk scale height and position angle on time scales as short as a day, corresponding to the angular resolution of theimages, and up to the 1.11 year span between two observing epochs. We present a basic retrieval of temporal changes
in the disk density structure, based on the images. We find
23 that the inner disk changes position angle on time scalesof months, and that the change is not axisymmetric, suggesting the presence of a non-axisymmetric dynamical forcing
on ∼1 au size scales. We consider two diﬀerent scenarios, one in which a quadrupolar disk warp orbits the central
star, and one in which an unequal-mass binary orbiting out of the disk plane displaces the photo-center relative tothe shadowing disk. Continued space-based monitoring of the Serpens Disk Shadow is required to distinguish between
these scenarios, and could provide unique, and detailed, insight into the dynamics of inner protoplanetary disks notavailable through other means.
Accepted by ApJ• サーペンスの[Great Disk Shadow]という場所について、HST近赤で1年くらいおいて2回

http://arxiv.org/pdf/2006.05965

観測した。

The circumstellar environment of EX Lup: SPHERE and SINFONI views• 時間変動が観測された。これは内側になにかあって動いているを示している
2,3,4
1
1
2,4
1

Elisabetta Rigliaco , R. Gratton , Á. Kóspál
, D. Mesa , V. D’Orazi , P. Ábrahám , S. Desidera ,C. Ginski5,6 , R. G. van Holstein6,7, C. Dominik5 , A. Garufi8 , T. Henning9 , F. Menard10 , A. Zurlo11,12,13 ,D. Maurel10 , P. Blanchard14 and L. Weber•15shadowは（尋常じゃなく）広がっていて、17000au, 0.24光年にわたる。
1

1

INAF/Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, 35122 Padova, Italy; 2 Konkoly Observatory,• モデリングの結果、1auくらいの物体が非軸対称に動いているというモデルで説明できた。
Research Centre for Astronomy and Earth Sciences,
Konkoly-Thege Miklós út 15-17, 1121 Budapest, Hungary; 3 MPIA,
Königstuhl 17, D-69117 Heidelberg, Germany; 4 ELTE Eötvös Loránd University, Institute of Physics, Pázmány Péter
sétány 1/A, 1117 Budapest, Hungary; 5 Anton Pannekoek Inst for Astronomy, Univ of Amsterdam, Science Park 904,
1

The Great Shadow in Serpens
Variability in the Great Disk Shadow
EC 97

CK 2

• 照射源はEC82。Eastern LobeとWestern

3(Background Star)

Ea
ste
rn

Lobeとかかれたところに影が見えるLobe

EC 123• 天体までの距離はVLBIで436pcと改められた

EC
82
EC 122EC 86（以前は250pc）。これをもとにすると、
30Lsunで光ってる。1-2Myrの星だと思うと、
温度5400-6100K, 質量2.5-3Msunくらいの星
がいる。

EC 74

SMM 3

outflowWe• Lobeの大きさは40秒角くらい、約1700 au広

ste
rn

Lo
be

EC 90

EC 105

がっている。0.27光年。100日。
N• とはいえよく見えるのは西が45日、東が75日く

EC 117
10’’

E

らいの領域なので、ここより内側で議論する。
Figure 1. Annotated two-color WFC3 image of the Serpens core in the vicinity of EC 82, in F125W (blue) and F164N (red).
[A higher-resolution version of this figure is available in the journal article.]

Each image was processed through the Drizzlepac
package, using Tweakreg for individual dithered expo-

I binary EC 90 to the south and the Serpens Reflection
Nebula to the north, which is transversed by the EC

Date

Instrument

FilterExposure time

Variability in the Great Disk Shadow[seconds]
2017 Jul 222018 Aug 30
2018 Aug 30

WFC3
WFC3
WFC3

F160W
F164N
F125W

5

1597
5998
1798

2017 Jul 22
Both data sets were downloaded from the Mikulski
Archive for Space Telescopes (MAST) in FLT format, as
output from the calwf3 data reduction pipeline, which
includes basic calibration of the raw data, including bias
and dark current subtraction, linearity correction, flat
fielding, bad pixel masking, and cosmic ray removal.

• 1.6ミクロンイメージ。
404日あけて撮影

N

tary
riety
detionnvesle in
e-on
rved
feldt
riodSTIS

Table 1. Observing log

E

find
rees,
nner
cant
time

• 下は(August 2018/
July 2017)の比。

Ratio• 西のローブが時計回りに動いた。東のローブ
は反時計周りに動いた。

Figure 2. Top panel: Rotated and scaled 1.6 µm images of EC 82 for the two observing epochs, presented as an animation.
The color scale is logarithmic and scaled to the same arbitrary unit to emphasize the shadow morphology. Bottom panel: The
1.6 µm ratio image of the two epochs (Jul 2017/Aug 2018), with lighter colors indicating higher ratio. The resulting ratio image
highlights the change in angle of the two shadow lobes. It also demonstrates that the change in angle happens in opposite
directions for the two lobes. [This figure is animated in the published journal article.]

where R = Data

Model is the residual and C is thethe two epochs are overlaid as a function of light-travelformulation uses the approximation that the vertical coobes, R,
respectively,andwein units of R, and ✓ is the polar angle measured fromordinate z = R sin(⇡/2 ✓) ' R(⇡/2 ✓), appropriate
tativetheanalysis
to
45
days.
Pontoppidan
et al. height is parameterizedaxisモデリング
of rotation.
The
scale
for ✓ ' ⇡/2. The optical depth profile, ⌧ , is then a00of HST
at
1.6
µm
of
0.
151
hRthe(R)=source
hR (R
(R/R
The
dust
surfaceau (4000as) fromcentralalongeach) ⇥one-dimensionalpressure-supported
viscousdisk
modelof the form
outer
outer ) . andfunctionof✓:m),thetodata
in principle
al-years,
pondinga light-travel
time of 0.27
(Hartmann et al. 1998):
density
profile
is
also
a
power
law,
⌃
(R)
/
R
.
This
d
00
days.
Thus,
a
changeinthegeometryofZ
disk perturbations on time
Routerhisk, or itsformulation
central light source,ontimescales⌃dthat(R)
1
uses the approximation
thevertical
co2
⇢(R,
✓)C
dR
(2)
⇢d (R, ✓) = p ⌧ (✓) ⇥=
exp
((⇡/2
✓)/h
, (1)
R ) ext
100 days can lead to changes in the shadow22⇡hR R
inner
適当な円盤
を仮定す
ordinatez=Rsin(⇡/2✓)'
R(⇡/2✓), Rappropriate
ullwhen comparingwe
the comtwo epochs.
ngextent
for variability,
where hR = H/R is the disk scale height H at radiuschanges in the disk taking place on timeC
is theextinction
coefficientattheobserving waveext
for
✓
'
⇡/2.
The
opticaldepth
profile,
⌧
,
is
then
a
nd
F164N
(epoch2)filters,
R,
in
units
of
R,
and
✓
is
the
polar
angle
measuredfromhan 100 days will manifestる。すると、中心星からのoptiacal
as changes overdepthはlength,
.
This
model
assumessinglescattering,but intheaxisofrotation.
The
scaleheightisparameterized
nd
While
the
F164N
angepasses.
of radiiintheshadow,astheperturbaone-dimensional function of
✓:
as
h
hR (R
) . The dust
surface filled in either
R (R) =
outer ) ⇥ (R/R
outer
practicetheshadowisprobably
partlyoutwards at the speed of light. However, thee [Fe II] line at 1.64 µm, itdensity profile is also a power law, ⌃d (R) / R . This
est-defined within ⇠ 45 days and ⇠ 75 days
Z Routerbyformulation
multipleuses
scattering,or byscatteringof photons fromthe approximationthatthe vertical coledi↵erenceintheshadow
ern and eastern lobes, respectively, and weordinatez = R sin(⇡/2✓) ' R(⇡/2✓), appropriate of scatteringothersources.
Further,thedistribution
strictourquantitative
analysis
to
45
days.
⌧
(✓)
=
⇢(R,
✓)C
dR
(2)
ext
ission, and we interpret the 00
for
✓
'
⇡/2.
The
opticaldepth
profile,⌧, is then angular resolution of HST at 1.6 µm of 0. 151Rinnerdust
is not uniform.
We therefore model the intensity
one-dimensional
functionof✓:gh atduetoachangein
broad。
Half Maximum), the data in principle alprofile of the
shadowby addinga linear continuum, as
Cext
is perturbations
the extinction
at thewaveZ observing
measurementof disk
on time coefficientR⌧ (✓) =background⇢(R, ✓)CextdR(2)
ort as 10 hours.
well as asingle
linearcomponent:
length,. Thissubmodel assumes
scattering,
but
in
ud
contains
significant
R

•

•

outer

urpose of searching for variability, we comadow
morphology
most
60W (epoch1) and F164Nare(epoch2) filters,ve overlapping band passes. While the F164N
relativemeasurements bees coverage of the [Fe II] line at 1.64 µm, it
is anyillustrated
in Figurethatlarge scale di↵erencein the2,
shadowyio
is dueto line between
emission, andthe
we interpretimage
two the↵erences as being due to a change in broad
017).
The ratio image highillumination.
thereflectingsource
cloud contains
significant subminantof variabilhanges in the shadow morphology are mostt apparent change is thatn di↵erential or relative measurements bewo epochs.
This acrossis illustrated
Figure 2,
ositionangle
theinfullws the 1.6 µm ratio image between the two
eral degrees, and that theC highgust 2018/July 2017). The ratio image
ccurredindominant
oppositeanguhadow
as thesourceof variabilBGwesternregion.
The most
change BG
is thatbes. Thatis,apparent
theas changed its position angle across the full
irection,
the andeastheshadow bywhereas
several degrees,that thegle change has
occurred in
opposite anguclockwise
direction
betweenns for the two lobes. That is, the western
dow
lobes are therefore notin a clockwise direction, whereas the eastved being
in a counter-clockwise
direction betweenchmore out-of-plane018. The two shadow lobes are therefore nots unequivocal evidence forwith the 2017 epoch being more out-of-planeem,althoughcannotevidence
im- for
018 epoch.
This is we
unequivocalof the system, although we
nmetrya non-axisymmetryof cannot
the im-inner

Cext is the extinction coefficient at the observing wavepractice the shadow is probably
partlyfilledsingle
in scattering,either but inlength, I(✓)
. This =model
assumes(IC (✓) IBG ) ⇥ exp( ⌧ ) + IBG ,(3)practicethe shadow isofprobably
partly fromfilled in either
by multiple モデリングは
scattering, or byscattering
photonsby multiple
scattering,or by+scattering
photonsfrom the intensity
where
IBG
= ABG
BBG ⇥of ✓.
Finally,other sources. Further, theofdistribution
scatteringother distribution
sources. Further, theof scatteringprofileisconvolvedbya
one-dimensional
Gaussian ker（ちゃんと書いていないが）Backgroundを
dust
is
not
uniform.
We
therefore
model
theintensity
dust is not uniform. Weneltherefore
model
the
intensity
withof thea fullwidth
at half
maximumof 0.00 151 to simprofile
shadowby addinga linearcontinuum, as
profile of the
shadow byulate
addinglinearcontinuum,aswell as a linear
backgroundcomponent:
I̲BGとして、散乱光の強さは上で定義したτ
the aWFC3
point
spread function
at 1.6 µm. We
well as a linear backgroundcomponent:
generally
assumepriors for all paI(✓)= (IC (✓) uninformative,
IBG ) ⇥ exp( ⌧ ) + Ior
BG ,flat,(3)•で減衰して散乱された光を足すモデリング。
where I The= A single+ B exception⇥ ✓. Finally, isthethatintensity
rameters.
we let the linear
BG

BG

BG

profile is convolved by a one-dimensional Gaussian kercontinuum
be fixed,
definedby 151
thetosurface
brightness
IBG⇥ aexp(
⌧ at
)+
BG ,
nel )withfull width
halfImaximumof 0.00(3)simonulate
eachthespreadshadow,
to a highly
the side
WFC3ofpoint
functioncorrespondingat 1.6 µm. We
where I （ちゃんと書いてないが）スケールハイトは= A + Bconstrained⇥ ✓.assume
Finally,
theorintensity
BG
generally
uninformative,
flat, wepriorsall paprior.
However,
dofor assume
a flat priorrameters. The single exception
is that wekerlet the linear
profile is convolved by a for
one-dimensional
Gaussian
the backgroundcomponent.
continuum
be fixed, definedby the surface brightness
円盤外縁でのスケールハイトのようだ。
00
on eachof the (Foreman-Mackey
shadow,to et
a highlynel with a full width at half
maximumof 0. corresponding
151 to simWe
usesideemcee
al. 2013) to sample
constrained prior. However, we do assume a flat priorthefor posterior
probability
ulate the WFC3 point spreadfunction
at 1.6 distributionµm. We of the model pathe backgroundcomponent.
We use emcee(Foreman-Mackey
et to
al.
to sample
rameters.
we wishsearchfor temporal varigenerally assume uninformative,
orBecause
flat,
priorsfor
all2013)
pathe posterior probability distribution of the model pa大事なのは円盤のPAが全体でシフトしている
ationsin theshadowgeometry
within a single epoch, as
rameters. The single exception
that
the forlinear
rameters.is
Becausewewe
wishletto search
temporal variwellasinbetween
we retrieve
independent
probaations
the shadowepochs,
geometry within
a single epoch,ascontinuum befixed, definedbytheepochs,
surfacebrightness
こと
well
as
betweenweretrieveproba-distance of thebility distributions of slicesindependentof constantI(✓) = (I (✓)••

Figure 5. Summary of constrained parameters as a function of light-travel time. The top panels show the epoch 2 image(F164N), while the lower three rows show the retrieved position angles, disk scale heights, and background intensities, respectively. Each epoch is indicated by a di↵erent color and symbol. The western lobe is much better defined, as the eastern lobe isstrongly a↵ected by the presence of a bright source (EC 86).

Variability in the Great

fea
fla

• featureはたぶんscattering cloundのintrinsicなもの• position angleをみるべき• negative changeが時計、positivechangeが反時計• この変化は、1auから a few auくらいだろう• 1イメージの中での変化が少ないから、45
日は変化がないと思おう。

• もし周期的なら、4 45日=180日以上だろう。

• これは、ケプラータイムで言うと0.85au
より外側

Figure 6. Relative change in position angle between thetwo epochs as a function of Epoch 2 date. The values foreach lobes are o↵set relative to each other by one day, forclarity.

radii, based on the short time scale for the shadow variability, EC 82 disk is likely self-shadowed at large radii.
Is EC 82 a classical, full disk, or a more evolved,
transitional disk? The system appears to have significant amounts of hot material (up to 1500 K) near thecentral star, as evidenced by strong veiling at 1–2 µm(Doppmann et al. 2005; Gorlova et al. 2010), strong
mid-infrared silicate emission, and strong rovibrational
CO emission lines (Banzatti & Pontoppidan 2015). The
mid-infrared slope of the SED is flat (see Figure 5 andDunham et al. 2015), indicative of a full disk. Pontoppidan, & Dullemond (2005) found that it was diffi-• puﬀed up inner rimの変化と捉
えよう

Fi
Th
et
al.
20
&
wa

bo

議論
• シナリオ１：これが伴星や惑星だったらどうか？
10
Pontoppidan et al.
disk warp rather than a bipolar warp, which would proangleの変化が違うのがおかし• 東と西でPosition
duce a symmetric “wobble”. It is not presently clear if
い。このような変化はbipolarではなくquadrupole
there is a theoretical basis for such a quadrupolar warp,
and further modeling is required to test this idea.

ShadowDisk

likeな構造が良さそう。warp?

4.1. An alternate explanation for the variability
We have discussed a model in which the variability is
シナリオ2：伴星がいて、中心星自体が重心に対して運•
caused by an orbiting, or precessing, quadrupolar disk
動しているというシナリオ。
warp. However, there is an alternative explanation: A
low-mass companion to EC 82, orbiting out of the planeof the disk, may move the photo-center of the source• 図を参照relative to the disk plane. This would also lead to nonplanar (flapping) change in observed shadow position
angle. This scenario
is illustrated in Figure 8. It requires
diskをはいてしまって円盤が• ただ、これだとinneran unequal-mass (or unequal-luminosity) binary, as an
無くなりそう。SEDからはinnerdiskが示唆されて
equal-mass binary would not move the photo-center at
any point in its orbit. The main argument against a stelいるので難しいか？
lar binary may be that this scenario is at odds with thepresence of an optically thick inner disk, as evidenced by
the spectral energy distribution of the system (see Sec-Disk
Shadow

Figure 8. Sketch showing an alternative explanation forthe shadow variability, in which an unequal-mass binary orbiting out of the disk plane shifts the photo-center of theillumination. This predicts that the inner disk is cleared-out
of material and that the position angle variation is strongly
periodic. The sketch is not to scale.
5. CONCLUSION

http://arxiv.org/pdf/2006.05965

38 The circumstellar environment of EX Lup: SPHERE and SINFONI views
Elisabetta Rigliaco1 , R. Gratton1 , Á. Kóspál2,3,4 , D. Mesa1 , V. D’Orazi1 , P. Ábrahám2,4 , S. Desidera1 ,C. Ginski5,6 , R. G. van Holstein6,7, C. Dominik5 , A. Garufi8 , T. Henning9 , F. Menard10 , A. Zurlo11,12,13 ,D. Maurel10 , P. Blanchard14 and L. Weber15
INAF/Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, 35122 Padova, Italy; 2 Konkoly Observatory,
Research Centre for Astronomy and Earth Sciences, Konkoly-Thege Miklós út 15-17, 1121 Budapest, Hungary; 3 MPIA,
Königstuhl 17, D-69117 Heidelberg, Germany; 4 ELTE Eötvös Loránd University, Institute of Physics, Pázmány Péter
sétány 1/A, 1117 Budapest, Hungary; 5 Anton Pannekoek Inst for Astronomy, Univ of Amsterdam, Science Park 904,
Amsterdam, The Netherlands; 6 Leiden Observatory, Leiden University, 2300 RA Leiden, The Netherlands; 7 ESO,
Alonso de Córdova 3107, Casilla 19001, Vitacura, Santiago, Chile; 8 INAF/Osservatorio Astrofisico di Arcetri, Largo
Enrico Fermi 5, 50125 Firenze, Italy; 9 MPIA, Königstuhl 17, D-69117 Heidelberg, Germany; 10 Univ. Grenoble Alpes,
IPAG, F-38000 Grenoble, France; 11 Núcleo de Astronomía, Facultad de Ingeniería y Ciencias, Univ Diego Portales,
A&A proofs: manuscript no. main
Santiago, Chile; 12 Escuela de Ingeniería Industrial, Facultad de Ingeniería y Ciencias, Univ Diego Portales, Santiago,
Chile; 13 Aix Marseille Univ., CNRS, LAM, Marseille, France; 14 Aix Marseille Univ., CNRS, CNES, LAM, Marseille,
France; 15 Geneva Observatory, Univ of Geneva, Chemin des Mailettes 51, 1290 Versoix, Switzerland
1

EXLupというoutburstを起こす天体の、静穏期をちゃんと観測して様子を明らかにした、という論文
中心から0.3-0.55秒角のところに、何らかの構造を発見

E-mail contact: elisabetta.rigliaco at inaf.it
EX Lup is a well-studied T Tauri star that represents the prototype of young eruptive stars known as EXors. They
are characterized by repetitive outbursts that are due to enhanced accretion from the circumstellar disk onto thestar. In this paper, we analyze new adaptive optics imaging and spectroscopic observations to study EX Lup and its
circumstellar environment in near-infrared in its quiescent phase. We aim to provide a comprehensive understandingof the circumstellar environment around EX Lup in quiescence, building upon the vast store of data provided by
the literature. We observed EX Lup in quiescence with the high contrast imager SPHERE/IRDIS in the dual-beampolarimetric imaging mode to resolve the circumstellar environment in near-infrared scattered light. We complemented
the data with earlier SINFONI spectroscopy, which was also taken in quiescence. We resolve, for the first time inscattered light, a compact feature around EX Lup azimuthally extending from ∼280◦ to ∼360◦ and radially extending
from ∼0.3′′ to ∼0.55′′ in the plane of the disk. We explore two diﬀerent scenarios for the detected emission. The
first one accounts for the emission as coming from the brightened walls of the cavity excavated by the outflow whose
presence was suggested by ALMA observations in the J = 3 − 2 line of 12 CO. The second attributes the emission
to an inclined disk. In this latter case, we detect, for the first time, a more extended circumstellar disk in scattered
light, which shows that a region between ∼10 and ∼30 au is depleted of µm-size grains. We compare the J−, H−, andK−band spectra obtained with SINFONI in quiescence with the spectra taken during the outburst, showing that all
the emission lines result from the episodic accretion event. Based on the morphology analysis, we favor the scenariothat assumes the scattered light is coming from a circumstellar disk rather than the outflow around EX Lup. We
determine the origin of the observed feature as either coming from a continuous circumstellar disk with a cavity, fromthe illuminated wall of the outer disk, or from a shadowed disk. Moreover, we discuss the potential origins of thedepleted region of µm-size grains, exploring the possibility that a sub-stellar companion may be the source of thisfeature.
Accepted by A&A
http://arxiv.org/2006.09787

Fig. 1. Reduced SPHERE-IRDIS polarimetric Q , U and Q –r2 scaled images. The coronagraph is marked with a blue circle. We assumed a39

Ring formation and dust dynamics in wind-driven protoplanetary discs: global simulations
A. Riols1 , G. Lesur1 and F. Menard1
Univ. Grenoble Alpes, CNRS, Institut de Planétologie et d’Astrophysique de Grenoble (IPAG), F-38000, Grenoble,
France1E-mail contact: antoine.riols at univ-grenoble-alpes.fr
Large-scale vertical magnetic fields are believed to play a key role in the evolution of protoplanetary discs. Associatedwith non-ideal eﬀects, such as ambipolar diﬀusion, they are known to launch a wind that could drive accretion in
the outer part of the disc (R > 1 AU). They also potentially lead to self-organisation of the disc into large-scale
axisymmetric structures, similar to the rings recently imaged by sub-millimetre or near-infrared instruments (ALMAand SPHERE). The aim of this paper is to investigate the mechanism behind the formation of these gaseous rings,
but also to understand the dust dynamics and its emission in discs threaded by a large-scale magnetic field. To
this end, we performed global magneto-hydrodynamics (MHD) axisymmetric simulations with ambipolar diﬀusion
using a modified version of the PLUTO code. We explored diﬀerent magnetisations with the midplane β parameter
ranging from 105 to 103 and included dust grains – treated in the fluid approximation – ranging from 100 µm to 1
cm in size. We first show that the gaseous rings (associated with zonal flows) are tightly linked to the existence of
MHD winds. Secondly, we find that millimetre-size dust is highly sedimented, with a typical scale height of 1 AU at
R = 100 AU for β = 104 , compatible with recent ALMA observations. We also show that these grains concentrate
into pressure maxima associated with zonal flows, leading to the formation of dusty rings. Using the radiative transfer
code MCFOST, we computed the dust emission and make predictions on the ring-gap contrast and the spectral index
that one might observe with interferometers like ALMA.
Accepted by A&A
http://arxiv.org/pdf/2006.01194

Characterizing the radio continuum nature of sources in the massive star-forming region
W75N (B)A. Rodríguez-Kamenetzky1 , C. Carrasco-González2 , J. M. Torrelles3 , W. H. T. Vlemmings5, L. F.

• 計算手法•

10 9 cm3 s 1 is the ion-neutral colspectively, the neutral and ion mas
やったこと：円盤のグローバルMHDシミュレーション
dividual mass. By introducing thealso define the dimensionless amb

v2A
i ⇢i
A. Riols,
G.formation
Lesur andandF. dustMenard:Ringinformation
Am =in PP discs=.
• 軸対称を過程（2.5D）、PLUTOがもと。
A. Riols, G. Lesur and F. Menard:
Ring
dynamics
PP discs and dust dynamics⌦⌘A⌦

エルサッサー数•
We note that the unit of time is taken as the inner ⌦ 1 and the up to a certain height z . The height z corresponds to th1
io
io generally need
We
note
that
the
unit
of
time
is
taken
as
the
inner
⌦
and
the
up
to
a
certainheightz.
Theheightz
corresponds
to
the
base
To
calculate
⇢
,
one
io
io
i
非理想MHDは両極性拡散のみいれた• unit of length corresponds
unit toofRlength
corresponds= 10AU.
Theionisation
unit of massof theionisation
above
which
the
FUV
can
penetr= 10AU.
The unit toof Rmass
is of thelayer,isabove
which
the FUVlayer,canpenetrate.
If
ious ionisation sources (X rays, csuch thatat⇢R==110AU.
in the midplane at R = 10AU.
we note⌃i (z) ascolumn
the horizontally
averaged column densi
such that ⇢ = 1 in the midplane
we note ⌃i (z) as the horizontally
averageddensity,
inand
FUV
from
the
star)
as
well
a
簡単のため、赤道面からとある高さ
The magnetic
fieldwhichforce
appearstegrated
in the fromLorentz
force boundary,tegrated fromthe verticalboundary,
The magnetic field B, whichappears in
theB,
Lorentz
the vertical
we assumethat FUV
are we assume that FU• 流体としてダストも入れた。
This
requiresthe2(r ⇥ B) ⇥ B, is governed
inductionfor equation,⌃i (z) > ⌃ic =までAm=const（order
0.0005 g/cm
This>tioncorresponds
to g/cm(rby⇥ theB) non-ideal⇥ B, is governed
byequation,
the non-idealblocked
induction
blockedfor ⌃2 .i (z)⌃ic mechanisms.
= 0.0005.
Thiscorrespo
unity）とchemistryoccurring
the gasphzio ' 3.5H, where H is thezio
gas'disc
scaleheight.
avoid
3.5H,
where
H To
is thegasanydisc scale inheight.
To avo
⇥
⇤
@B
⇥
⇤
@B
thetheas muchtoasmap
discontinuity
at z = z(4)use a smoothatfunction
make• パラメータ
discontinuity
z =simplify
zioto
, we
use
aproblemsmooth function
io , we= r ⇥ (v ⇥ B) + r ⇥ ⌘A=(Jr⇥⇥eb(v) ⇥⇥eB)) ⇥ eb .
b . + r ⇥ ⌘A (J ⇥ eb (4)して、ある高さから上は大きくした@t@ttransition between the midplane
regions
and thethe
ionisedlayer, regions
so the essential
transition
between
midplaneand the ionisedlaythatcapturesphysics
⇥
⇤
⇥
⇤
that
the
Ambipolar
Elsassernumberisthat
theAmbipolar
Elsassernumber
is
The
term
r
⇥
⌘
(J
⇥
e
)
⇥
e
corresponds
toambipolardi↵uthe
e↵ect
of
FUV
in
the
disc
coro
A
b
b
The
term
r
⇥
⌘
(J
⇥
e
)
⇥
e
corresponds
toambipolardi↵uAbbプラズマβが1e5,1e4,1e3にふった。
• sion, with J = r ⇥ B the current density, eb = B/kBk the unit
2
3!
!
2 stant22!2 and of!2 3order unity in the midp
sion, with J = r ⇥ B the current density, eb = B/kBk the666 ⌃icunit ⌃ic 777
666 ⌃ic
777⌃
6
7
vector parallel to the field
line,
and
⌘
the
ambipolardi↵usivity.
ic
A
Am
=
Am
exp
+
,
(6)
6
7
0 di↵usivity.
775 ,vector parallel to the field line, and ⌘A the ambipolar
4 ⌃+5 Am0 exp 664
Am
=
+
⌃
+
To simplify the problem,
weassumethatthedi↵usion
of
mag⌃
⌃
1
mm,
1
cm
。流体とした。
• ダストサイズは100μm,
To simplify the problem, we assume that the di↵usion of magnetic field is only due to ion-neutral collisions (ambipolar di↵u+
netic field is only due to ion-neutral collisions(ambipolardi↵uwhereAmistheconstant
midplanevalue,⌃
is the column den0+sion) and is assumed to be the dominant non-ideal e↵ect in the
where
Am
is
the
constant
midplanevalue,⌃
is the colum
0
sion)
and
is
assumedtobethedominant
non-ideal
e↵ect
in
the
sity
integrated
from
the
top
boundary,
and⌃thecolumndenregions considered in this paper, that is R 10 AU (see justificasity boundary.
integratedThis
fromthe top profile
boundary, and ⌃ the column• エルサッサー数の鉛直分布は手で与えた。
regions
consideredin
thispaper,
that
is
R
10
AU
(see
justificasity
integratedfromthebottomdi↵usivity
tions in Simon et al. 2015). Therefore, the Hall e↵ect and ohmic
sity integrated
from thebottomboundary. This di↵usivity p
isdesigned
to
mimic
the
conductivity
obtainedfromcomplete
tions
in
Simon
et
al.
2015).
Therefore,
the
Hall
e↵ect
and
ohmic
物理的には、FUVが侵入して来れるdi↵usion are neglected. The di↵usivity ⌘A is not uniform, and its
isprotoplanetary
designed to mimic
the conductivity
obtained from com
thermochemical
modelsofdiscs
includingdust
di↵usionareneglected.
The
di↵usivity⌘
isnotuniform,
anditsA
vertical profile is prescribed in Section 2.1.3.
• 円盤モデル
thermochemical
modelsprotoplanetary discs includin
grains, with Am0 ' 1. (Thi場所に対応
et al. 2018). We notethat ofto convert
vertical profile is prescribed in Section 2.1.3.
Am0 the' 1.disc(Thi
et al.of2018). We note that to cnumerical values of ⌃i intograins,
g/cm2 ,with
we usedmodel
2numerical
values
of
⌃
into
g/cm
, we used the disc mo
2.1.2.
Surfacedensity
profileEq.
(5).
i= 200 AU
• Rin = 10 AU、Rout
2.1.2. Surface density profile
Eq. (5).
We assume a disc model with surface density following
We2 assume a disc model with surface density
2.1.4.following
Temperature and hot corona
0.5⌃ = 200 (R/1 AU) g/cm .
(5)
2.1.4.
Temperatureandhotcorona0.5
2•Complex thermo-chemical
models of protoplanetary discs⌃ = 200 (R/1 AU) g/cm .
(5)This corresponds to a rather flat density profile that has been (Kamp & Dullemond 2004;Complex
Thi et al. thermo-chemical2018) suggest that models
thereof protoplanetary•••

A&A proofs: manuscript no. ambi_global

averaged between t = 250T 0 and t = 1000T 0 .

0.00.00112

7.5

0.2

7.5
0.9

5.0

0.00084

0.0

5.0

vR/vK

1.8
2.5

0.00056
2.5

8

2.7

6

4

2

0.00028

0
z/H

2

4

6

8

3.6
2.5

0.00000

24

5.4

0.00056

7.5
6.3
0
3.0

2

4

6

8

2.5 2.0 1.5 1.0 0.5 0.0
Stream lines (in log of sonic mach)

10

0
7.2

BR/B0

0.00028

5.0

7.5

0

2.5

4.5
5.0

Bf

0.0
Density r

0.0

vz/vK

0.2

2

4

6

8

10

3
2
1
0
1
2
Field lines (colorbar in log of Alfvenic mach)

0.000840.00112

Bf /B0864

左がガス密度と流線。右が磁力線。Bφはカラー。
点線の位置がz=3.5H: そこより赤道面側では両極性拡散が効いている。
温度も同じ位置でtransitionするモデルになっている
全領域積分したwind のoutﬂow rateはṀ w (rout )6

10-8M⊙ /yrs.

0
z/H

2

4

6

8

Fig. 2. Top: Vertical profiles of the radial and vertical velocity. Bottom:
Vertical profiles of the radial and toroidal magnetic fields. Quantities
are averaged between R = 25AU and R = 50AU and in time betweent = 250T 0 and t = 1000T 0 .
Article number, page 6 of 17

Fig. 1. Left: Average gas density and streamlines in the poloidal plane. Right: Average toroidal field and poloidal field lines. All quantities areaveraged between t = 250T 0 and t = 1000T 0 .

2

上は、オレンジの線が鉛直方向のガス速度。3.5Hより外側で外に向かってものが流れている、すな
わちウインドが吹いていることを示している

Fig.
integ
angu

spee
cool
nally

nd F. Menard: Ring formation and dust dynamics in PP discs

0.2
0.1
20

30

40

50

60

70
Bz

0.001

0.000

20

30

40
50
R in AU

60

70

Fig. 5. Surface density ⌃ defined as ⇢ (top) and vertically-integrated↵ M (defined magnetic
Fig. 5. fieldSurface⌃ definedas ⇢ (top)and vertically-integrated
Bz density(bottom)
as a functionof radius
R. Quantities are av↵
and↵(defined
R
M. The dashed eraged
magnetic
field between
Bz (bottom)
as aandfunction
oforbits,
radius R.
Quantities
are avin
time
250
1000
and
verticallybetweenorbits. The dashed eraged in time between 250 and 1000 orbits, and vertically between✓ = ✓ and ✓ = ✓+ (z ' ±3.5H). The dashed red lines are the location ofents.
✓ = ✓ and ✓ = ✓+ (z ' ±3.5H). The dashed red lines are the location of
the thedensitymaxima.
density maxima.

図5：時間平均した面密度と縦磁場の強さ。
ガス面密度に、リング・ギャップ構造ができている。
シミュレーションの間（一番内側で１０００回転）、このリングギャッ
プ構造はsteadyであった。
図６：350回転時のポロイダル磁場分布。ギャップのところで縦磁場が強いのがわかる。
図７：流線も書いてみた。南半球では降着、北半球でその反対。ギャッ
プに流れ込むような流線は、観測 Teague et al. 2019とよく一致している
A&A proofs: manuscript no. ambi_globaloward
zerozeroocity
toward1. 1.
l equalto 0.5⌦
to 0.5⌦7.5

z in AUand pure gaseous
ure simulation
gaseous is
hisital period is
at themulationme-averaged
eriod at thedeneen R =den10 anderagedbase of the disc10 and✓=d (respectivelyof the discdelimitedby therespectively
o the frontiers at
gion
ited andby thethefullysoto the jump
frontiers
at inFig. 6. Snapshot of the simulation for = 103 at t = 350T 0 , zoomed in
the
cold
disc
and
nd the fully
emitted and ac- the inner part of the disc. The colour map represents the density, while
the jump in/2 ± ✓d , with su- the white lines are the poloidal magnetic field lines.
Fig. 6. Snapshot of the simulation for = 103 at t = 350T 0 , zoomed in
old
disc
and
e total mass outtedandinnerac- andthe inner part of the disc. The colour map represents the density, while
een thewhite lines
arethetheformation
poloidal ofmagnetic
field lines.
4.3 isrings, which
are visible in Fig. 1✓ned , particularity
with su- isthe Section
midplane,le mass
out-with (top) and appear after 200 inner orbits in the simulation.
spherew beingheinner( ˙and

0.00
0.75

5.0

1.50

2.5

2.25
3.00

0.0

3.75

2.5

4.50

5.0

5.25

7.5

6.00

10

20

30

R in AU

40

50

60

Fig. 7. Snapshot of the simulation for = 103 at t = 350T 0 , zoomed in on the inner part of the disc. The colour map represents the density, while
the black lines are the poloidal streamlines. The dashed white lines delimit the z = ±3.5H surface.

like in the box simulations. This term is actually associated withthe laminar magnetic field that supports the wind structure. We

Teague et al. (2019) in 12 CO lines in the disc surface aroundHD 163296. Above this altitude, the flow is mainly vertical(pointing outward)
and dominated by the wind component.
リング構造の起源についての議論

R zd

is a net accretion flow z vR dz/(2zd ) (with zd = 3.5H), whichdvaries from 0.002 vK0 at R = 15AU to 0.01 vK0 at R = 100AU.
This accretion flow is located mainly at the disc’s southern surface (see Figs. 2 and 7) and should transport the magnetic flux
radially over several AU. However, we do not see any motionof the magnetic structures during the course of the simulation;
they rather stay at a fixed position. To better understand this, we
examine the induction equation for B✓ ' Bz

One important question is to understand the origin of thesestructures (rings + magnetic shell) when ambipolar di↵usion
is the dominant non-ideal MHD e↵ect. Recently, Riols &Lesur (2019) proposed that the rings appear spontaneously inwind-emitting discs and are the result of a linear and seculari 1 @ hiinstability, driven by the MHD wind. The process works as @( B✓ ) 1 @ h=rE =r(ur B✓ u✓ Br + ⌘A J? ) .
(28)
follows: a small initial radial perturbation of density generates
@t
r @r
r @r
ion equation
for
B
'
B
(Eq.
28),
integrated
verticallybetween
✓
=
✓
and
✓
z
a radial flow
directed
towards the gaps, because of the radial
500T 0 . The red dashed lines correspond to the local maxima of Bz (averaged
viscous stress. The magnetic field, initially uniform, is radially Figure 9 shows the radial profile of each term in the right-handtransported towards the gaps (by the radial flow) and the excess side of the induction equation. We see that the term related to theof poloidal flux induces a more efficient mass ejection in the vertical stretching of the field 1r @r@ ( ru✓ Br ) plays little role in the
ur- gaps, which reinforces the initial density perturbation. In this
Article number, page 9 of 17
ith instability model, the gas is then depleted vertically by wind
Bai plumes rather than being radially accumulated into rings. This15; process, however, has only been shown in a local shearing box
de- model, and its persistence in global configurations remains to be
ity demonstrated. In particular, global models include a net vertical
the stress and a net radial accretion velocity that were not present indes the local framework.
owout 小さな動径方向の摂動が加わる
We then explore how these rings form and evolve in theize global configuration. Firstly, in the snapshots of Fig. 6 and 7,→ギャップにガスが流れ込む（拡散）ion it seems obvious, in particular in the southern hemisphere, thatig- thewind is not uniform and homogeneous in the radius. Actu→初期に一様だった磁場が、引きづられる
nes ally, it is composed of wind plumes that emanate from the re→このpoloidal
ﬂuxがギャップでの質量放出を加速する
the gionswhere the poloidal
field is maximum (gap regions). We
can checkedthat streamlines connected to the gap regions have the→よりガスがギャップに流れ込む
nto largest vertical (outward) velocity component. In addition to
ide that, in Fig. 8 (top) we show the radial divergence of the flow
our 1/r @/@r(rvr ), averaged in time and in the vertical direction ini & the disc, as a function of the radius. We see that the divergence18) is positive within the rings and negative within the gaps. Thismin
8 0 R !m>literature,
we
also
label
each
dust
componentbyits
Stokesnum>>
0
fixed in space and time. To make an easier comparison with the
Rmin
>
>
<
ber
measuredfor ⇢label= ⇢0each= 1,dust
c s =componentc s0 = 0.05byand= ⌦0 num= 1, B = >>B0
if R < Rmin>
literature,
we alsoits⌦Stokes
z
>
R0!m
>
which
we
label
St
and
refer
to
as
the
’reference
Stokesnumber’.
>0<
ber measured for ⇢ = ⇢0 = 1, c s = c s0 = 0.05 and ⌦ = ⌦0 = 1, B = >
R
>
>
>
z
We
note,
however,
that
the
local
Stokesnumber(whichcontrols
B0>!m if R > Rmin:>
which we label St0 and refer to as the ’reference Stokes number’.
>
RR0
>
>
the
dynamics)
dependsontheposition
in
the
disc
for
one
species.
>>
We note, however, that the local Stokes number (which controls
B0
if R > Rmin>:
Therefore,
we depends
need to finda correspondence
betweengrain with m = R
0
2.5/2
andG.RLesur=R= 10Ring
AU.
Again
the
prescription
the dynamics)
on theposition in the discfor onethe
species.
min and
0Menard:A.
Riols,
F.
formationanddust
dynamics in PP discs
size
and
the
Stokesnumber.
For
that,
we
assume
that
in
the
limit
R<
R=min2.5/2
is to avoid
the=divergenceof theverticalfield at the
Therefore, we need to find a correspondence between the grain at
with
m
and
R
R
=
10
AU.
Again
the
prescription
min
0
of
small
magnetisation,
the
disc
is
in
hydrostatic
equilibrium
and
p
With
such
dependence,
it is possible
toverticaldefine fielda constantsize and the Stokes number. For that, we assume that in the limit pole.
at
R
<
R
is
to
avoid
the
divergenceoftheat theminitscolumn(orsurface)density
is
⌃
=
⇢
H
2⇡,
where
⇢
is
the
0
0
ordependence,
plasma parameterin thetomidplanefor R >
of small magnetisation, the disc is in hydrostatic
equilibrium and magnetisation
p
pole.
With
such
it
is
possibledefineaconstant
midplaneand Hdensity= c s /⌦
height.
its columndensity(or surface)
is is
⌃ the
= ⇢0disc
H scale2⇡, where⇢0 Thus,is the Rmin :magnetisationor plasma parameter in the midplane for R >combiningthesedi↵erent
relations,weobtain:midplane density and H = c s /⌦ is the disc scale height. Thus,150 R :1502
2
min 2⇢c
2⇢0 c s04
0 p
combiningthese1 di↵erent relations, we obtain:s= 22 =BBB ⇢ s 2⇡ CCC
Bz s
2⇢B020c2s02⇢c
St = ak B@0 p CA1
100 =
100
=
BBB ⇢ s ⌃ 2⇡ CCC
B2z
B20
CA
St = ak B@⌃
6
in the midplane. To estimate the Stokes number as a function 2.4. Averages
St0 = 0.001 (a = 100µm)
St0 = 0.03 (a = 3 mm)
50
50
ofsize, we To
assumethat thethe Stokes
disc surface
densityfollows 2.4. Averages
in grainthe midplane.
estimate
numberasafunction
We define a vertical integral of a quantity Q between angles ✓3
Eq.
5. Bysize,
considering⇢ s =that
2.5 theg/cmdisc, we
obtainthe followingof grain
we assume
surfacedensityfollows and ✓+ as
We define a vertical integral of a quantity
Q0 between angles ✓conversion
inthe
midplane:380
Eq. 5. By considering ⇢ s = 2.5 g/cm , we obtain the following
and ✓Z+ as✓+
conversion
inthe
midplane:✓ R ◆1/2✓ a ◆Q = Z ✓ r sin ✓Qd✓,(16)k+
St ' St0 ✓
with
St
=
0.1
.
(12)
✓
0
50
50
✓1 acm ◆10 R
AU ◆1/2
Q
=
r
sin
✓Qd✓,
(16)
k
10
St
with St0 = 0.1.
(12) and, respectively,✓
1 ' St0a radial integral and temporal average:
We note 10
thatAUin principle, dust particles
couldbeporous
and
have
a
1 cm
more
complex geometry. This would not change the dynamics described
100T average:
Z routZ
1
a radial integral and temporal
We note that in principle, dust particles could be porous and have a100 and, respectively,f1
here, but only the relation between the Stokes number and the physicalmore complex geometry. This would not change the dynamics described hQiR = Z r Qdr and hQi =(17)
Z T f Qdt.
12
size of the particles (11).
out
T f 1 Ti T irinhere, but only the relation between the Stokes number and the physical150 hQiR =150 Qdt.
Qdr and hQi =(17)
size of the particles (11).
T f Ti T irin
Article number, page 4 of 17

ダストをいれた• ダストサイズを固定していである（＊）ことから、その動径依存性はとなるはずである（ガス面密度はr^-0.5を仮定している）。
Article number, page 4 of 17

•

4

6

z

るので、ストークス数は

2

0

50

100
R (AU)

150

140Fig. 10. Dust density distribution for grains of size 100 µm (left) and 3 mm (right), forinner orbits*ストッピングタイムの定義のときに、分母のガスの熱速度のところを間違えて音速にしていて、その結果
π/8 間違ってる50

100
R (AU)

150

8

10

12

14

= 103 . This is averaged in time between 400 and 650

A&A proofs: manuscript no. ambi_global

g. 10. Dust density distribution for grains of size 100 µm (left) and 3 mm (right), for
ner orbits

b = 105100= 103 . This is averaged in time between 400 and 650

b = 104b = 103
0.05 St 1/2

1000.016 St

0.03 St 1/2

1/2

b = 103b = 104
Gas pressure

0.00020.00020.0001100

0.0000
10

20

30

40

50

60

0.0000
10

20

30

40

50

60

20

30

40

50

60

20

30

40
R in AU

50

60

Hd /Hg10 2
10

a = 3 mm

3

10 5

10 5

10 1

10

10 3

10 2
10 1
Effective St

10 7
10

1

10 1

10 3

10 2
10 1
Effective St

20

30

40

10 3
10 3

10 2
10 1
Effective St

g. 11. Dust to gas scale height ratio as a function of the e↵ective Stokes number for = 105 (left), = 104 (centre) andolour accounts for a given particle size (purple: 1cm, blue: 3 mm, yellow: 1mm, red: 0.3 mm and green: 0.1mm.)= 103 (right). Each5060

a = 100µm

10 4
10 5
10

10 8
10

10 4

20

30

40
R in AU

50

60

10

Hd/Hg。
Article number, page 11 of 17
Fig. 12. Top panels shows the gas pressure as a function of radius R for = 103 (left) and = 104 (right). The central and bottom panels show the
ダスト巻き上げの動径依存性の関係式は、傾きは良さそう

dust surface density for grains of 3 mm and 100 µm, respectively. Quantities are averaged vertically between ✓ = ✓ and ✓ = ✓+ (z ± 3.5H) and intime between 400 and 650 inner orbits.

＊αに直したらもっと議論できる気がするが…

the gravitational settling and the di↵usion due to the turbulence. time, for two di↵erent plasma beta parameters and two di↵erent
In that theory,ダストサイズに違いとトラッピングの違い
the turbulence is assumed to be homogeneous and grain sizes. For = 103 , both large and small grains concentrateisotropic, and can be described through a uniform, e↵ective dif- into pressure maxima (with position delimited by the red dashedfusion coefficient. We note that the local simulations of Riols & vertical lines). However, 3mm grains, the contrast in density, atLesur (2018) found a deviation from the simple power law St 1/2 some radii, can be 200 to 500 times larger than the contrast meaat = 103 and for small particles. In our case, we do not see such sured for 0.1 mm grains. Although the density contrast is pretty
a deviation. Moreover, for = 104 and = 103 , the ratio Hd /Hg high for a = 3 mm, the dust-to-gas ratio remains bounded andis two times larger than those measured in the local shearing box smaller than 0.04. We also note that the radial dust structures are(see comparison with Fig. 7 of Riols & Lesur (2018)). This is thinner in the case of larger particles. For = 104 , the results are
due to a larger vertical rms turbulent velocity measured in the similar, except that the pressure maxima are less pronounced andmidplane in global simulations (vzrms ⇠ 0.07c s ) compared to lo- the dust is less concentrated on average. This is particularly obcal simulations (where vzrms ⇠ 0.04c s , see Fig. 2 of Riols & Lesur vious for 0.1 mm grains, for which the contrast in density never(2018) for = 103 ).
exceeds ten. Also, the typical separation between the dust ringsin the case = 104 is about 2.5 H at R ' 20 AU, while it isabout 4 H for = 103 . By fitting the radial density profiles with
5.3. Concentration in pressure maxima
gaussian functions, we find that the half-width of the dust rings
at R = 20AU is ⇠ 0.3 0.4 AU (0.3-0.4 H) for = 103 andAs we showed in Section 4.3, rings of matter arise spontaneously a = 3 mm, while it is ⇠ 0.15 AU for = 104 . For comparison,in wind emitting discs, and they correspond to pressure maxima

A. Riols, G. Lesur and F. Menard: Ring formation and dust dynamics in PP discs

2

50
100

1

150

0

100

0
x in AU

100
10 20

100
50

4

100

0
2
1

100
150

100

0
x in AU

100

100

150

1.0

50

0.6

50

150

100

0.8

0

5

3

50

50

1.2

100

0
x in AU

100

0.4
0.2

100

0.0

150

0.6

50

0.4

50

100

0.2

100

150

0.0

150

0
x in AU100

Fig. 13. Top: Flux density ⌫F⌫ in W/m2 /pixel calculated by MCFOST, with the disc viewed face-on (left panel) and edge-on (right panel) for = 1
mm. This is calculated for t = 500T 0 and = 103 . Bottom: same images but convolved with a Gaussian kernel with a full width half maximum(FWHM) of 3 AU. The white dashed lines in the right panels shows the surface z = Hg for R = 200 AU.

1
100

0
x in AU

0

100
10 19

8

0

4
3
2
1
100

0
x in AU

100

0

3

50

2

100

1

150

0

100

0
x in AU

100
10 19

6
5

50

4

0

3

50

2

100
150

7

1
100

0
x in AU100

Fig. 15. Top: Flux density ⌫F⌫ in W/m2 /pixel calculated by MCFOST, with the disc viewed face-on (left panel) and edge-on (right panel) for= 1.65 microns. This is calculated for t = 500T 0 and = 103 . Bottom: Same images but convolved with a Gaussian kernel with a full width half
maximum (FWHM) of 10 AU. The white dashed lines in the right panels shows the surface z = Hg for R = 200 AU.

ALMA観測シミュレーション。B7.
gas-to-dust ratio is 100. Such a re-normalisation proceFor that we use the 3D continuum and line radiative transfer code that the
下の図は、3auのガウシアンカーネルで畳み込み
SPHEREの観測シミュレーション
dure canbe achieved because the dust has almost no feedback
MCFOST based on the Monte Carlo method (Pinte et al. 2006).
Assuming a given gas and dust density distribution, the code
first computes the temperature structure assuming that the dust
is in radiative equilibrium with the local radiation field. This isdone by propagating photon packets, originally emitted by thecentral star (assumed to be of solar-type), through a combinationof scattering, absorption and re-emission events until they exit
the computation grid. During this step, we use 1 280 000 photon
packets in total. It is assumed that the dust opacities are fixed anddo not depend on the temperature during this stage. Observables’
quantities (images) are then obtained via a ray-tracing method,

4

100

6
5

5

50

150

7

50

0

100

2

100

1.0
0.8

3

150

1.2

50

4

7
6

100

6

0
50

10 20

7
5

10 19

150

z in AU0

1.4

8z in AU3150

y in AU

4

100

10 19

150

y in AU

50

150

10 20

y in AUy in AU100

5z in AU150

z in AU

10 20

A&A proofs: manuscript no. ambi_global

which are indeed contained in a layer of ⇡ 0.5Hg at R = 200
AU according to Fig. 11. We note that the emission is lower inon the gas (since the dust-to-gas ratio remains low in all simu- the midplane of the disc than in the layers above (see Appendix
下の図は10auガウシアンカーネルで畳み込み
lations).
We note that during the interpolation, we assume that A showing the vertical profile of the flux density). This is due
the smallest grains of 0.03 µm follow the gas distribution exactly. to the smaller temperature in the midplane, resulting from thestrong sedimentation of mm to cm grains.
We also show, in the bottom panel of Fig. 14, the spectral inIn the upper panels of Fig. 13, we show the image of our discfor = 103 calculated at a given wavelength ( = 1 mm), corre- dex ↵ s = d log F(⌫)/d log ⌫ as a function of the radius obtained
sponding to ALMA band 7, in a face-on (left) and edge-on (right) by comparing the image at = 1 mm with the image at = 3
configuration. To take into account the resolution of instruments mm. The spectral index ranges between 2.5 and 3.5, suggesting
like ALMA, in the lower panels we show the same images con- that the disc deviates slightly from a black body (whose specvolved with a Gaussian kernel with a full width half maximum tral index is 2 in the Rayleigh-Jeans regime). The spectral index

Accepted by A&A
40

http://arxiv.org/pdf/2006.01194

Characterizing the radio continuum nature of sources in the massive star-forming region
W75N (B)A. Rodríguez-Kamenetzky1 , C. Carrasco-González2 , J. M. Torrelles3 , W. H. T. Vlemmings5, L. F.
Rodríguez2 , G. Surcis6 , J. F. Gómez7 , J. Cantó8 , C. Goddi9,10 , J. S. Kim11 , S. -W. Kim11 , N. AñezLópez3,4 , S. Curiel8 and H. J. van Langevelde12,13
1 Instituto de Astronomia Teorica y Experimental, (IATE-UNC), X5000BGR Cordoba, Argentina; 2 Instituto deRadioastronomia y Astrofisica (IRyA-UNAM), 58089 Morelia, Mexico; 3 Institut de Ciencies de l’Espai (ICE, CSIC),Can Magrans s/n, E-08193, Cerdanyola del Valles, Spain; 4 Institut d’Estudis Espacials de Catalunya (IEEC), E08034, Barcelona, Spain; 5 Department of Earth and Space Sciences, Chalmers University of Technology, SE-43992
Onsala, Sweden; 6 INAF-Osservatorio Astronomico di Cagliari, Via della Scienza 5, 09047 Selargius (CA), Italy; 7
Instituto de Astrofisica de Andalucia, CSIC, Glorieta de la Astronomia s/n, E-18008 Granada, Spain; 8 Institutode Astronomia, Universidad Nacional Autonoma de Mexico (UNAM), Apartado Postal 70-264, DF 04510 Mexico; 9
ALLEGRO/Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands; 10 Departmentof Astrophysics/IMAPP, Radboud University Nijmegen, P.O. Box 9010, 6500 GL Nijmegen, The Netherlands; 11
Korea Astronomy and Space Science Institute, 776 Daedeokdaero, Yuseong, Daejeon 305-348, Republic of Korea; 12
Joint Institute for VLBI ERIC (JIVE), Oude Hoogeveensedijk 4, 7991 PD Dwingeloo, The Netherlands; 13 Leiden
Observatory, Leiden University, PO Box 9513, 2300 RA, Leiden, The Netherlands
1E-mail contact: adriana.rodriguez at unc.edu.ar
The massive star-forming region W75N (B) is thought to host a cluster of massive protostars (VLA 1, VLA 2, andVLA 3) undergoing diﬀerent evolutionary stages. In this work, we present radio continuum data with the highest
sensitivity and angular resolution obtained to date in this region, using the VLA-A and covering a wide range of
frequencies (4-48 GHz), which allowed us to study the morphology and the nature of the emission of the diﬀerent
radio continuum sources. We also performed complementary studies with multi-epoch VLA data and ALMA archive
data at 1.3 mm wavelength. We find that VLA 1 is driving a thermal radio jet at scales of ≈0.1 arcsec (≈130 au), but<also shows signs of an incipient hyper-compact HII region at scales of <1
arcsec(∼∼ 1300 au). VLA 3 is also drivinga thermal radio jet at scales of a few tenths of arcsec (few hundred of au). We conclude that this jet is shock-exciting25https://arxiv.org/pdf/2006.08566v1
41

The Role of Outflows, Radiation Pressure, and Magnetic Fields in Massive Star Formation
Anna L. Rosen1 and Mark R. Krumholz2,3Center for Astrophysics | Harvard & Smithsonian, 60 Garden St, Cambridge, MA 02138, USA; 2 Research Schoolof Astronomy and Astrophysics, Australian National University, Canberra, ACT 2611 Australia; 3 ARC Centre of
Excellence for Astronomy in Three Dimensions (ASTRO-3D), Canberra, ACT 2611 Australia

1E-mail contact: anna.rosen at cfa.harvard.edu
Stellar feedback in the form of radiation pressure and magnetically-driven collimated outflows may limit the maximum
mass that a star can achieve and aﬀect the star-formation eﬃciency of massive pre-stellar cores. Here we present a
series of 3D adaptive mesh refinement radiation-magnetohydrodynamic simulations of the collapse of initially turbulent,
massive pre-stellar cores. Our simulations include radiative feedback from both the direct stellar and dust-reprocessedradiation fields, and collimated outflow feedback from the accreting stars. We find that protostellar outflows punches
holes in the dusty circumstellar gas along the star’s polar directions, thereby increasing the size of optically thinregions through which radiation can escape. Precession of the outflows as the star’s spin axis changes due to theturbulent accretion flow further broadens the outflow, and causes more material to be entrained. Additionally, thepresence of magnetic fields in the entrained material leads to broader entrained outflows that escape the core. We
compare the injected and entrained outflow properties and find that the entrained outflow mass is a factor of ∼3larger than the injected mass and the momentum and energy contained in the entrained material are ∼25% and ∼5%of the injected momentum and energy, respectively. As a result, we find that, when one includes both outflows andradiation pressure, the former are a much more eﬀective and important feedback mechanism, even for massive stars
with significant radiative outputs.
Accepted by ApJ
http://arxiv.org/pdf/2006.04829

http://arxiv.org/pdf/2006.04829
42

Planet Migration in Self-Gravitating Discs: Survival of Planets
Sahl Rowther1,2 and Farzana Meru1,2
1

Centre for Exoplanets and Habitability, University of Warwick, Coventry CV4 7AL, UK; 2 Department of Physics,
University of Warwick, Coventry CV4 7AL, UK
E-mail contact: sahl.rowther at warwick.ac.uk
We carry out three-dimensional SPH simulations to study whether planets can survive in self-gravitating protoplanetary
discs. The discs modelled here use a cooling prescription that mimics a real disc which is only gravitationally unstablein the outer regions. We do this by modelling the cooling using a simplified method such that the cooling time inthe outer parts of the disc is shorter than in the inner regions, as expected in real discs. We find that both giant(> MSat ) and low mass (< MNep ) planets initially migrate inwards very rapidly, but are able to slow down in the inner
gravitationally stable regions of the disc without needing to open up a gap. This is in contrast to previous studies
where the cooling was modelled in a more simplified manner where regardless of mass, the planets were unable toslow down their inward migration. This shows the important eﬀect the thermodynamics has on planet migration. In
a broader context, these results show that planets that form in the early stages of the discs’ evolution, when they arestill quite massive and self-gravitating, can survive.
Accepted by MNRAS
https://arxiv.org/pdf/2006.03077
26