星形成ゼミ
2020/09/04
SFN 331 #22-#28

長谷川哲夫（国立天文台）

1

New Hot Corinos with ACA
1,222
Shih-Ying
HsuSurvey, Sheng-Yuan, TieGalacBc
Liu3 , Dipen, NaomiHirano2 etal.
331 ALMA
of OrionLiu
PlanckColdSahu
Clumps(ALMASOP):
I. DetecBonof New
1#22
Hot
Corinoswith (NTU),ACA Taiwan (R.O.C.); 2 Institute of Astronomy and Astrophysics, Academic Sinica,
National
TaiwanUniversity

Shih-Ying(R.O.C.)
Hsu1,2, Sheng-Yuan Liu2, Tie Liu3, Dipen Sahu2, Naomi Hirano2 et al.
Taipei, Taiwan
ApJ 898, 107H

E-mail contact: syliu at asiaa.sinica.edu.tw
We report the detection of four new hot corino sources, G211.47−19.27S, G208.68−19.20N1, G210.49−19.79W and
G192.12−11.10 from a survey study of Planck Galactic Cold Clumps in the Orion Molecular Cloud Complex with the
Atacama Compact Array (ACA). Three sources had been identified as low mass Class 0 protostars in the Herschel Orion
Protostar Survey (HOPS). One source in the lambda Orionis region is firstly reported as a protostellar core. We haveobserved abundant complex organic molecules (COMs), primarily methanol but also other oxygen-bearing COMs
(in G211.47−19.27S and G208.68−19.20N1) and the moleculeof prebiotic interest NH2 CHO (in G211.47−19.27S),
16
signifying the presence of hot corinos. While our spatial resolution is not suﬃcient for resolving most of the molecularemission structure, the large linewidth and high rotational temperature of COMs suggest that they likely reside in thehotter and innermost region immediately surrounding the protostar. In G211.47−19.27S, the D/H ratio of methanol([CH2 DOH]/[CH3 OH]) and the 12 C/13 C ratio of methanol ([CH3 OH]/[13 CH3 OH]) are comparable to those of otherhot corinos. Hydrocarbons and long carbon-chain molecules such as c-C3 H2 and HCCCN are also detected in the four
sources, likely tracing the outer and cooler molecular envelopes.
Accepted by ApJ
http://arxiv.org/pdf/2006.15850

BV (RI)c photometric study of three variable PMS stars in the field of V733 Cephei
Sunay Ibryamov1 , Evgeni Semkov2 , Stoyanka Peneva2 and Kristina Gocheva1
Department of Physics and Astronomy, University of Shumen, 115, Universitetska Str., 9700 Shumen, Bulgaria;
Institute of Astronomy and National Astronomical Observatory, Bulgarian Academy of Sciences, 72, Tsarigradsko
Shose Blvd., 1784 Sofia, Bulgaria

1

2

E-mail contact: sibryamov at shu.bg
This paper report results from the first long-term BV (RI)c photometric CCD observations of three variable pre-mainsequence stars collected during the period from February 2007 to January 2020. The investigated stars are located2in the field of the PMS star V733 Cep within the Cepheus OB3 association. All stars from our study show rapid

3

Department of Physics and Astronomy, University of Shumen, 115, Universitetska Str., 9700 Shumen, Bulgaria;
331
photometric
studyof three variable
PMS
stars inAcademy
the ﬁeld
V733 72,
Cephei
InstituteBVof (RI)c
Astronomy
and National
Astronomical
Observatory,
Bulgarianof of
Sciences,
Tsarigradsko
Sunay
Ibryamov1,
EvgeniSemkov2,
Stoyanka
Peneva2
and
KrisBna
Gocheva1
Shose#23Blvd., 1784 Sofia, Bulgaria

1

2

A&Ap accepted

E-mail contact: sibryamov at shu.bg
This paper report results from the first long-term BV (RI)c photometric CCD observations of three variable pre-mainsequence stars collected during the period from February 2007 to January 2020. The investigated stars are locatedin the field of the PMS star V733 Cep within the Cepheus OB3 association. All stars from our study show rapid
photometric variability in all-optical passbands. In this paper, we describe and discuss the photometric behavior ofthese stars and the possible reasons for their variability. In the light variation of two of the stars we found periodicity.
Accepted by Research in Astronomy and Astrophysics
https://arxiv.org/pdf/2006.02498

The ionized warped disk and disk wind of the massive protostar Monoceros R2-IRS2
seen with ALMA
Izaskun Jiménez-Serra1 , Alejandro Báez-Rubio1 , Jesús Martín-Pintado1 , Qizhou Zhang2 and Víctor
M. Rivilla3
Centro de Astrobiología (CSIC-INTA), Ctra. de Torrejón a Ajalvir km 4, 28850, Torrejón de Ardoz, Spain; 2 Center
for Astrophysics, Harvard & Smithsonian, 60 Garden St., Cambridge, MA 02138 USA; 3 INAF-Osservatorio Astrofisico
di Arcetri, Largo E. Fermi 5, I-50125 Firenze, Italy

1E-mail contact: ijimenez at cab.inta-csic.es
Theories of massive star formation predict that massive protostars accrete gas through circumstellar disks. Althoughseveral cases have been found already thanks to high-angular resolution interferometry, it remains unknown the internal
physical structure of these disks and, in particular, whether they present warps or internal holes as observed in lowmass proto-planetary disks. Here, we report very high angular resolution observations of the H21α radio recombinationline carried out in Band 9 with the Atacama Large Millimeter/submillimeter Array (beam of 80 mas×60 mas, or
70 au×50 au) toward the IRS2 massive young stellar object in the Monoceros R2 star-forming cluster. The H21α line
shows maser amplification, which allows us to study the kinematics and physical structure of the ionised gas aroundthe massive protostar down to spatial scales of ∼1-2 au. Our ALMA images and 3D radiative transfer modelling 4revealthat the ionized gas around IRS2 is distributed in a Keplerian circumstellar disk and an expanding wind. The H21α

Izaskun Jiménez-Serra , Alejandro Báez-Rubio , Jesús Martín-Pintado , Qizhou Zhang and Víctor
M. Rivilla3

331 The ionized warped disk and disk wind of the massive protostar Monoceros R2-IRS2
Centro de Astrobiología (CSIC-INTA), Ctra. de Torrejón a Ajalvir km 4, 28850, Torrejón de Ardoz, Spain; 2 Center#24
seen with
ALMA
forAstrophysics,
Harvard& Smithsonian, 60 Garden St., Cambridge, MA 02138 USA; 3 INAF-Osservatorio Astrofisico

1

IzaskunJiménez-Serra1
Báez-Rubio1, Jesús Maren-Pintado1 , Qizhou Zhang2 and Víctor M. Rivilla3
di Arcetri,
Largo
E. Fermi 5,, Alejandro
I-50125 Firenze,
Italy
ApJ 897, L33J

E-mail contact: ijimenez at cab.inta-csic.es
Theories of massive star formation predict that massive protostars accrete gas through circumstellar disks. Althoughseveral cases have been found already thanks to high-angular resolution interferometry, it remains unknown the internal
physical structure of these disks and, in particular, whether they present warps or internal holes as observed in lowmass proto-planetary disks. Here, we report very high angular resolution observations of the H21α radio recombinationline carried out in Band 9 with the Atacama Large Millimeter/submillimeter Array (beam of 80 mas×60 mas, or
70 au×50 au) toward the IRS2 massive young stellar object in the Monoceros R2 star-forming cluster. The H21α line
shows maser amplification, which allows us to study the kinematics and physical structure of the ionised gas aroundthe massive protostar down to spatial scales of ∼1-2 au. Our ALMA images and 3D radiative transfer modelling revealthat the ionized gas around IRS2 is distributed in a Keplerian circumstellar disk and an expanding wind. The H21αemission centroids at velocities between −10 and 20 km s−1 deviate from the disk plane, suggesting a warping for the
disk. This could be explained by the presence of a secondary object (a stellar companion or a massive planet) within
the system. The ionized wind seems to be launched from the disk surface at distances ∼11 au from the central star,consistent with magnetically-regulated disk wind models. This suggests a similar wind launching mechanism to thatrecently found for evolved massive stars such as MWC349A and MWC922.

17

5

Submillimeter Array (SMA) revealed that MonR2-IRS2 is a
UC H II region with its RRLs at λ0.85 mm experiencing
maser ampliﬁcation. However, these observations were unable
to resolve the internal structure of the source (Jiménez-Serra
et al. 2013).
We present high angular-resolution observations(80 mas×60 mas, 70 au×50 au) of the H21α RRL carriedout with the Atacama Large Millimeter/submillimeter Array(ALMA) toward the MonR2-IRS2 UC H II region. Our H21α
images show that the ionized gas in MonR2-IRS2 is distributed
in a warped Keplerian disk and an ionized wind that islaunched at radii ∼11 au.
2. Observations
The H21α RRL at 662.40416 GHz was observed on 2015
September 25 with ALMA in Band 9 (project 2012.1.00522.S)
using the 12 m array (baselines from 43 m to 2.27 km).
Observations were performed in dual polarization mode using
a spectral bandwidth of 1.875 GHz and a spectral resolution of
976 kHz (0.44 km s−1). As calibrators, J0522−3627 was used
for bandpass calibration, J0423−013 for ﬂux calibration, andJ0607−0834 for phase and amplitude calibration.
The calibrated data set was obtained by running the original
pipeline reduction scripts, and the data were additionally selfcalibrated using the CASA (the Common Astronomy Software
Applications) package version 5.0.0. The MonR2-IRS2 spectra
are clean from molecular emission, thus the 0.4 mm continuum
map was generated using the CASA task uvcontsub across5.61 GHz of line-free bandwidth. The angular resolution of the
0.4 mm continuum image was 60 mas×50 mas, P.A.=78°
and the rms noise level was 1.9 mJy beam−1. The H21α image
was obtained from the continuum-subtracted data cube aftersmoothing the velocity resolution to 1 km s−1. The beam andrms noise level of the H21α image were 80 mas×60 mas, P.
A.=−88°, and 30 mJy beam−1 in 1 km s−1-channels,
respectively.

3. Results
3.1. 0.4 mm Continuum Emission
Figure 1 (left panel) reports the 0.4 mm continuum map
obtained toward MonR2-IRS2 (black contours and grayscale).
The source is barely resolved at the resolution of ALMA
(deconvolved size 62 mas×43 mas, or 55 au×38 au). Its
morphology clearly deviates from Gaussianity at the lowemission level (12σ, with 1σ=1.9 mJy beam−1), likely due
to the presence of a wind (Section 4). By ﬁtting a 2D Gaussian
to the continuum emission, the peak coordinates of MonR2IRS2 are α(J2000)=06h07m45 8034 (±0 0008), δ(J2000)=−06°22′53 5155 (±0 0005), which are consistent
with the previously measured values (Jiménez-Serra et al.
2013).
The integrated continuum ﬂux at 0.4 mm is (345 ± 10) mJy.
This value is larger than that expected from the spectral indexof α=−0.16±0.2 measured by the SMA (with S ∝ ν α;Jiménez-Serra et al. 2013).5 The free–free contribution is130±30 mJy; thus while the continuum emission atλ 0.85 mm is entirely due to ionized gas, the λ=0.4 mm
continuum ﬂux is dominated by dust. The mass for the neutral
disk can then be estimated using a dust ﬂux of 215±32 mJy
and by assuming dust temperatures Tdust=500–1200 K(Howard et al. 1994), a gas-to-dust mass ratio R=100, anda dust opacity κdust=25.8 cm2 g−1 (consistent with H2 gas
densities of 108 cm−3 and no ices; Ossenkopf & Henning 1994).
The neutral disk mass is ∼2–4×10−4 Me, with a 15%
uncertainty.
3.2. H21α Maser Emission
In Figure 2, we present the H21α spectrum integrated over
the MonR2-IRS2 continuum source (Figure 1). Consistent with
the H30α and H26α lines measured with the SMA (JiménezSerra et al. 2013), the H21α transition shows two emission

6

Stefan Kraus , Jean-Baptiste LeBouquin , Alexander Kreplin , Claire L. Davies , Edward Hone , John
D. Monnier3 , Tyler Gardner3 , Grant Kennedy4 and Sasha Hinkley1

331 Spin-orbit alignment of the β Pictoris planetarysystem2
School Stefanof Physics
Astronomy,
Universityof Exeter,
UK;d’Grenoble
Kraus1,and
Jean-BapBste
LeBouquin2,
Alexander
Kreplin1,
Claire Universite
L. Davies1, Edward
Hone1, Alpes,
John Grenoble, France;
3#25
4
Dept ofD.Astron.,of Michigan,
Ann Arbor,
USA;Hinkley1
Dept of Physics, Univ. of Warwick, Coventry, UK
Monnier3,Univ.
Tyler Gardner3,
Grant Kennedy4
and Sasha
1

ApJ 897, L8K
E-mail contact:s.kraus at exeter.ac.uk

A crucial diagnostic that can tell us about processes involved in the formation and dynamical evolution of planetarysystems is the angle between the rotation axis of a star and a planet’s orbital angular momentum vector ("spin-orbit"
alignment or "obliquity"). Here we present the first spin-orbit alignment measurement for a wide-separation exoplanetary system, namely on the directly imaged planet β Pictoris b. We use VLTI/GRAVITY spectro-interferometry withan astrometric accuracy of 1 µas (microarcsecond) in the Brγ photospheric absorption line to measure the photocenter
displacement associated with the stellar rotation. Taking inclination constraints from astroseismology into account,we constrain the three-dimensional orientation of the stellar spin axis and find that β Pic b orbits its host star on
a prograde orbit. The angular momentum vectors of the stellar photosphere, the planet, and the outer debris disk
are well aligned with mutual inclinations 3◦ ± 5◦ , which indicates that β Pic b formed in a system without significant
primordial misalignments. Our results demonstrate the potential of infrared interferometry to measure the spin-orbitalignment for wide-separation planetary systems, probing a highly complementary regime to the parameter space
accessible with the Rossiter-McLaughlin eﬀect. If the low obliquity is confirmed by measurements on a larger sampleof wide-separation planets, it would lend support to theories that explain the obliquity in Hot Jupiter systems withdynamical scattering and the Kozai-Lidov mechanism.
Acceptedby ApJL obliquity angle ψThree-dimensional
β Pictoris
https://arxiv.org/pdf/2006.10784• 太陽系 ~< 7 deg でよくそろっている• 距離：19.45 +- 0.05 pc• しかし、transiBng extrasolar planetary systemsでは、sky• 年齢：若い 24 +- 3 Myr
projected obliquity
angle βexplanation
は 0 - 180 deg とばらついている• of質量：1.85
Msun、A6V型
Ionization:a possible
for the diﬀerence
mean disk
sizes in star-forming• 光球の見かけの直径 0.736 +- 0.019 mas
regions
その原因は？
• Projected rotaBon velocity v sin i = 130 km/s
M.
Kuﬀmeier1, B. Zhao2 and P. Caselli2
1. 複数の惑星や伴星による重力多体相互作用の可能性
1
Zentrum
für Astronomie,
Heidelberg
University, Albert Ueberle Str.
2, 69120 Heidelberg, Germany; 2 Max-Planck（e.g., Kozai-Lidov
oscillaBons
が惑星落下を早めると同時に惑星：
Institute for Extraterrestrial Physics, Garching, Germany
obliquity angle を大きくする → Hot Jupiters）
1. β Pic b 公転周期 ~20 yr、軌道長半径 10.6
E-mailcontact:
kueﬀmeieratgmail.com
2. そもそも原始惑星系円盤と原始星のスピン軸がそろってい+- 0.5 au
なかった可能性
2. revealed
β Pic c 公転周期
day、軌道長半径
Surveysof protoplanetary disks in star-forming regions of similar age
significant~1200
variations
in average disk
7 are
Wide-separaBon
planets
でのobliquity
angle
ψ
測定が重要
~2.7
au
mass in some regions. For instance, disks in the Orion Nebular Cluster (ONC) and Corona Australis (CrA)

and λBrγ and wBrγ are the central wavelength and width of theline. This function provides a good approximation for thephotocenter displacement due to rotation in a marginallyresolved photosphere. We ﬁt this function to the astrometry
vectors x(λ) in the complex plane:x (l ) = c · S (l ) ,(3 )where c is a complex number. The phase of c gives the position
angle of the spatial displacement between the blue- andredshifted line wings, i.e., the stellar equator is oriented along
Ws = tan-1(Rc Ic ). We determine the uncertainties throughbootstrapping.
For our data set on β Pic, we derive the sky-projected
position angle of the equator to Ωs=29°±4°, with thesouthwestern side of the photosphere approaching the observer.5 The bottom panel of Figure 2 shows the model function
overplotted on the measured astrometric signals.
There are some small residuals between model and data thatmight indicate a tiny (1–2 μas) offset between the line
photocenter and the continuum photocenter to the west (seethe redshifted wing in the bottom panel of Figure 2). Such an
offset could indicate that the continuum photocenter is slightly
shifted with respect to the center of the stellar photosphere, forinstance due to stellar surface structures, diffuse disk emission,or a point source within the ∼0 1 ﬁeld of view. However, theresiduals have a low signiﬁcance and data with higher signalto-noise and higher angular resolution is needed to investigate
this further. Also, future differential phases measurements at
higher spectral resolution could derive an independent estimate
for the inclination of the photosphere (see the model images ofa inclined rotating photosphere in Figure 8 of Kraus et al. 2012)or detect subtle effects in the velocity ﬁeld, such as related todifferential rotation. However, we do not expect these effects tocontribute to the residuals in our measurements.
4. Discussion

Equation (1) asx (l ) =l -1
B f (l)F.
2p(2 )Applying this equation to our data results in the photocenter
vectors shown in Figure 3. The R.A. and decl. components ofthese vectors are also shown as position–velocity diagrams in
Figure 2 (bottom). The data show that the photocenter in theblueshifted wing of the Brγ line is displaced to the northeast(Figure 3, right). As the line is in absorption, this indicates thatthe blueshifted absorption is strongest on the southwestern sideof the star, and thus the photocenter shifts to the northeast.
Accordingly, the blueshifted (approaching) side of the rotating
photosphere is toward the southwest. The phases predicted by
the model are overplotted on the measured differential phases
in Figure 4.
In order to determine the position angle of the stellar rotation
axis quantitatively, we parameterize the wavelength-dependence of the photocenter displacement between the blue- andredshifted line wing with a S-shaped function
S(λ)=a(λ)·exp(−a(λ)2), where a(λ)=2(λ−λBrγ)/wBrγ,

Together with the inclination estimate of is=89°. 1 derived
from the pulsation frequency spectrum6 (Zwintz et al. 2019),
our measurement of Ωs deﬁnes the full spatial orientation of thestellar rotation axis, enabling a comparison with the spatial
orientation of the planetary orbit and of the large-scaledebris disk.
The primary disk is likely seen nearly exactly edge-on, while
the warped disk is tilted by +4°. 0±0°. 6 in position angle withrespect to the primary disk (Golimowski et al. 2006; Lagrange
et al. 2012) and inclined by 6°±1° with respect to the line ofsight, with the northwestern part of the disk being tipped nearer
to Earth (Ahmic et al. 2009, see Figure 3, left and middle
panels). In order to compute the mutual inclination between the
different planes, we adopt the orbital elements of β Pic b
published by Gravity Collaboration et al. (2020), including an
orbital inclination ip=89°. 04±0°. 03 and the longitude of theascending node Ωp=31°. 88±0°. 05. For the primary disk
+0 .2 (Lagrange et al. 2012) andplane we adopt Wd,p = 29 . 30 .3
id,p=90°. 0±0°. 1 (Ahmic et al. 2009). The warped secondary
disk is oriented along Ωd,s=33°. 3±0°. 6 (Lagrange et al.

8regions
1
2
M.
Kuﬀmeier, B. Zhao
and P.
Caselli2
331
IonizaBon:
a possible
explanaBon
for the diﬀerence of mean disk sizes in star-forming

Zentrum
für Astronomie, Heidelberg University, Albert Ueberle Str. 2, 69120 Heidelberg, Germany; 2 Max-Planck#26
regions1

Institute M.
forKuﬀmeier1,
Extraterrestrial
Physics,
Garching, Germany
B. Zhao2 and
P. Caselli2
A&Ap 639, 86K
E-mail contact:
kueﬀmeier at gmail.com

Surveys of protoplanetary disks in star-forming regions of similar age revealed significant variations in average disk
mass in some regions. For instance, disks in the Orion Nebular Cluster (ONC) and Corona Australis (CrA) areon average smaller than disks observed in Lupus, Taurus, Chamaeleon I, or Ophiuchus. In contrast to previous
models that studied the truncation of disks at a late stage of their evolution, we investigate whether disks may
already be born with systematically smaller disk sizes in more massive star-forming regions as a consequence of higher
ionization rates. Assuming various cosmic-ray ionization rates, we computed the resistivities for ambipolar diﬀusion
and Ohmic dissipation with a chemical network, and performed 2D nonideal magnetohydrodynamical protostellar
collapse simulations. A higher ionization rate leads to stronger magnetic braking, and hence to the formation ofsmaller disks. Accounting for recent findings that protostars act as forges of cosmic rays and considering only mild
attenuation during the collapse phase, we show that a high average cosmic-ray ionization rate in star-forming regions
such as the ONC or CrA can explain the detection of smaller disks in these regions. Our results show that on average,a higher ionization rate leads to the formation of smaller disks. Smaller disks in regions of similar age can therefore
be the consequence of diﬀerent levels of ionization, and may not exclusively be caused by disk truncation throughexternal photoevaporation. We strongly encourage observations that allow measuring the cosmic-ray ionization degrees
in diﬀerent star-forming regions to test our hypothesis.
Accepted by A&A
https://arxiv.org/pdf/2006.00019

18

9

Zentrum für Astronomie, Heidelberg University, Albert Ueberle Str. 2, 69120 Heidelberg, Germany; 2 Institut für
331 Linear
dust
polarizaBon
during the embedded phaseof protostar
formaBon:
SyntheBc
Theoretische
PhysikundAstrophysik, Christian-Albrechts-Universität
zu Kiel,Leibnizstrasse
15, 24118 Kiel,
Germany;
3
4
Harvard-Smithsonian
for Astrophysics,#27
observaBonsCenterof bridge
structures 60 Garden Street, Cambridge, MA 02138, USA; Department ofSpace, Earthand Environment,
Chalmers
University
Technology, 412 96, Gothenburg, Sweden
M. Kuﬀmeier1,
S. Reissl1, S. Wolf2,
I. Stephens3
and H.of
Calcuz4
1

A&Ap 639,
137K
E-mail contact:
kueﬀmeierat gmail.com

Measuring polarization from thermal dust emission can provide important constraints on the magnetic field structurearound embedded protostars. However, interpreting the observations is challenging without models that consistently
account for both the complexity of the turbulent protostellar birth environment and polarization mechanisms. We
aim to provide a better understanding of dust polarization maps of embedded protostars with a focus on bridgelike structures such as the structure observed toward the protostellar multiple system IRAS 16293–2422 by comparing
synthetic polarization maps of thermal reemission with recent observations. We analyzed the magnetic field morphology
and properties associated with the formation of a protostellar multiple based on ideal magnetohydrodynamic 3D
zoom-in simulations carried out with the ramses code. To compare the models with observations, we postprocessed asnapshot of a bridge-like structure that is associated with a forming triple star system with the radiative transfer codepolaris and produced multiwavelength dust polarization maps. The typical density in the most prominent bridgeof our sample is about 10−16 g cm−3 , and the magnetic field strength in the bridge is about 1 to 2 mG. Inside thebridge, the magnetic field structure has an elongated toroidal morphology, and the dust polarization maps trace thecomplex morphology. In contrast, the magnetic field strength associated with the launching of asymmetric bipolar
outflows is significantly more magnetized (∼100 mG). At λ = 1.3 mm, and the orientation of the grains in the bridge
is very similar for the case accounting for radiative alignment torques (RATs) compared to perfect alignment withmagnetic field lines. However, the polarization fraction in the bridge is three times smaller for the RAT scenario thanwhen perfect alignment is assumed. At shorter wavelength (λ < 200µm), however, dust polarization does not trace
the magnetic field because other eﬀects such as self-scattering and dichroic extinction dominate the orientation of thepolarization. Compared to the launching region of protostellar outflows, the magnetic field in bridge-like structures isweak. Synthetic dust polarization maps of ALMA Bands 6 and 7 (1.3 mm and 870 µm, respectively) can be used asa tracer of the complex morphology of elongated toroidal magnetic fields associated with bridges.
Accepted by A&A
https://ui.adsabs.harvard.edu/pdf/2020arXiv200601297K/abstract

The Formation of a Stellar Association in the NGC 7000/IC 5070 Complex:
Results from Kinematic Analysis of Stars and Gas

10

M. Kuffmeier et al.: Toward a complete understanding of dust polarization
amin = 5 nm, amax = 6 m
104

ext
abs

104

sca

= 0.8

(m2 g 1)(m2 g 1)

10

103

3

102

101

101

100
10 1

102

100

101

( m)

102

103

100
10 1

100

101

102

103

( m)

Fig. 2. Left: mass opacities for scattering sca (short-dashed ocher), absorption abs (long-dashed red), and extinction ext = abs + sca (solid blue)
for a dust model with grain sizes between amin = 5 nm and amax = 6 µm, leading to an opacity index of = 0.8. The shaded areas indicate thewavelength regimes where dust polarization may be dominated by dichroic extinction (light blue), scattering (light yellow), and emission (lightred). The shaded areas are estimated under the assumption that the product of dust number density ⇢dust and path length ` gives ⇢dust ⇥ ` ⇡ 1 g cm 2(see Sect. 4.6 and Reissl et al. (2014) for further details). Right: same as the left panel, but with an upper grain size of amax = 100 µm, corresponding
to an opacity index of = 1.3.

Fig. 1. Column density seen along the z-axis of the coordinate system⇡4 kyr prior to the formation of the third protostar (blue star), when the
primary S1 (black star) is ⇡70 kyr old and the secondary(cyan star) et al.:M.S2Kuffmeier
Toward a complete understanding of dust polarizationis ⇡0.49 M and the secondary has a mass of ⇡0.25 M . The takes place predominantly within the densest regions. We thereis ⇡27 kyr old.
blue star illustrates the location at which the third protostar forms fore chose the exact value of the density such that most of thelarge dust grains are located within these dense regions and not in
the surrounding material. Consequently, the disks are more saturated with micrometer-sized grains and nanometer-sized grains
2.2.
Dust
model
turbulence in our box is driven by supernova explosions of typeare sparse. However, micrometer-sized grains are also assumed
II from massive stars (Kuffmeier et al. 2016). At the end of their The dust in the interstellar medium (ISM) is canonically modeled to be sparsely present within the bridge itself. For the dust-to-gas
mass-dependent lifetime, energy is injected into the box at the by a dust mixture of 62.5% silicate and 37.5% graphite following mass ratio we applied the canonical value of ⇢dust /⇢gas =ˆ 1% forlocation of the type II supernova explosion. To account for the a power-law size distribution of N(a) / a q (Mathis et al. 1977; both components (Mathis et al. 1977; Bohlin et al. 1978). The
field of measuring grain sizes is rapidly developing, therefore
thermodynamics in the GMC, we use heating and cooling tables Li & Draine 2001). Within the ISM, the range of grain sizes is we also produced synthetic maps assuming an upper grain size
usually
assumedtobebetweenaminimumandmaximumsizethat are based on models of Gnedin & Hollon (2012), who usedcutoff of amin = 5 nm and amax = 250 nm, where the quantity a of amax = 3 mm for the dense regime. In Appendix A we show
results for synthetic maps of emitted radiation, scattered radiaCLOUDY models (Ferland et al. 1998) (optically thin cooling), as
is the effective radius of a dust grain corresponding to a spheretion, polarization including radiative alignment torques (RATs),
explained in Kuffmeier et al. (2017, 2019). To account for UV of equivalent volume.
and perfect alignment for 53 µm, 214 µm, and 1.3 mm waveshielding of photoelectric heating at higher densities, the temHowever, dust scattering models of circumstellar disks andlength as based on an upper grain size for the dense regime ofperature is tapered down exponentially to T = 10 K for number protostellar sources suggest maximum grain sizes of about amax = 100 µm (Fig. A.1) and amax = 3 mm (Fig. A.2).
3
50
300
µm
(Kataoka
et
al.
2015,
2017;
Hull
et
al.
2017a;
densities n > 200 cm (see also Padoan et al. 2016). As protoIn Fig. 2 we show the resulting opacities  of extinction,
Ueda et al. 2020). While studies based on the dust opacity absorption, and scattering for the dust components with an upper
stellar heating is ignored in our model, most of the gas in the index = log ( ) log ( ) / log
log10 2 , orig- dust radius of a
1
1
10
10
10 1
max = 6 µm and amax = 100 µm. The dust compodensest regions is quasi-isothermal and cold.
inally suggested millimeter- to centimeter-sized grains, these nents have an opacity index of 0.8 and 1.3, respectively, whichFig.
3.
Density
distribution
inthebridgeandaroundforming
protostars.
Thepanelsillustratedensity distribution and the pattern of the strong
For the zoom-in, we select individual stars that have accreted grain sizes are expected to be erroneous becausethe
the studies is typical for the natal environment of young stars (Lommen et al.
magnetic
with thein the negative
direction
the axes.
bridgeshown
the left panels along the x-axis, in theignored thescatteringof
opacity(Lin etTheal. 2020;
Liuis2019;
Zhu in2007,
less than 2 M
by the fields
end ofassociatedthe simulation.
For primaryeach zoom-in2009).
etal.
2019).
However,millimeter-sizedgrainsarebelieved
toplay
middlepanelsalong
they-axis,
and
in
the
right
panelsalong
thez-axis.
Theorientationis
also
illustrated
with the red, green, and blue axes (forsimulation, we run the simulation again from the snapshot priora role inweshieldingfrom densities
line emission
in disks.
the x,ofy,theandz direction).
Toapply
showhigher
the bridge
more clearly,
only plotabove
10 16 g cm 3 and gradually increase the opacity of the gas
to the formation
selectedsink andresolution
A further complication is that the growth of grains is a 2.3. Synthetic simulations with the radiative transfer code
14
3
POLARIS0 to
100%densities> 10
illustration
based
on a cubicalboxsizeof distribution
2048 AU in lengtharound the center, computed asgas restin the regionfrom
where
this
star forforms,
while ⇢
theof theg cm
box is. Thecomplex
process,isand
the redistributionof the(rA + rresolution.
time
later. research that has many unanswered To compare the MHD simulation with observations of dust
evolved at coarser
We
do of
notformation
excise partofofthe
thethird
box,companion
remains a⇡4
fieldkyrof ongoingC ) / 2 at the11
but keep the full domain of (40 pc)3 . In the case of the zoom- questions (Ossenkopf & Henning 1994; Hoang et al. 2019; polarization, we postprocessed the simulation data with the pubin simulation of the protostellar multiple in this paper, we used a McKinnon et al. 2019; Kannan et al. 2020; Vogelsberger et al. licly available radiative transfer code POLARIS (Reissl et al.
4 kyr after this snapshot.

1.3 mmA&A 639, A137 (2020)
M. Kuffmeier et al.: Toward a complete understanding of dust polarization

Fig. 7. Left panel: Stokes I and the polarization result considering dust self-scattering at 1.3 mm only. Right panel: logarithmic ratio of scatteredradiation Iscat to the emitted radiation Iem corresponding to the intensity maps shown in Fig. 6 and on the left side.

Fig. 6. Top row: Stokes I component of the 1.3 mm dust emission considering RAT grain alignment (left) in comparison with perfect alignment(right). White contours show the logarithmic column density, while the length and orientation of the pseudovectors represent the degree and angleof the linear dust polarization. The polarization vectors are rotated by 90 to match the actual magnetic field direction. The area of the bridge
structure is enclosed by purple lines, while the natal structures that harbor the primary and secondary (D1) are shown in yellow and the natal
structure of the tertiary (D2) is plotted in green. The borders follow the contour line of I = 4.24 mJy beam 1 . Bottom row: same as the top row, but
the color-scale now shows the polarized fraction, Pfrac .

89 micronM.demonstrates
Kuffmeier et al.:a completeofunderstandingof dust We
polarization
dust polarization.
plot the PI for a wavelength of 1.3 mm
The right panel in Fig. 7 also
thatToward
the relative
radiation due to self-scattering is highest in the vicinity of the considering the cases of RATs, the combined intensity using
forming protostars. Compared to the radiation that is induced RATs and self-scattering, and perfect alignment. The resulting
by dust emission, radiation from self-scattering is only a minor PI relations for the these regions (marked, e.g., in Fig. 6) arecontributor at 1.3 mm wavelength. Self-scattering is only respon- shown in Fig. 8. We determined the slope ↵ with a least-squares
sible for <1% of the polarization in the bridge, and even in the fit in log-space. When the PI is plotted for the entire map, thevicinity of the protostars, it only contributes ⇠10% at most to the value of ↵ for the combined intensity of the polarization Pfrac
is slightly higher than that of RAT alignment, demonstrating apolarized emission at 1.3 mm.
minor contribution of scattered radiation to the total polarization
at 1.3 mm. In perfect alignment, the PI relation is completely
3.2.3. Polarization-intensity dependenceslost. When we compare the values of the degree to those derived
Numerous observations of cores in the submillimeter regime from real polarization measurements, we have to take the limshow a significant depolarization in the most luminous regions ited dynamic range (concerning the intensity) and the minimum(Henning et al. 2001; Wolf et al. 2003; Gonçalves et al. 2005; reliable polarization degree (typically a few 0.1%) into account.
Brauer et al. 2016). This (anti)correlation is commonly fitted Consequently, the observed degree is dominated by the densest,by the polarization-intensity (PI) relation Pfrac / I ↵ (Henning that is, (sub-)millimeter brightest regions.
et al. 2001). In this section we investigate how the polarization in
In the bridge alone, all three polarization cases barely depend
distinct regions behaves dependent on the different mechanisms on intensity. However, for the case of perfect alignment, theA137, page 7 of 17

Fig. 9. Left: same as the dust emission map in Fig. 6 considering only RAT alignment, but for a wavelength of 89 µm. Right: corresponding ratio of
scattered Iscat to emitted radiation Iem . Red contours indicate the ratio Iscat /Iem = 0.1. The polarization vectors are rotated by a considerable amountin D1, D2, and the bridge compared to the 1.3 mm observations in Fig. 6.

degree of polarization is roughly ten times higher than in theother two cases.
For dense cores, the slope ↵ is typically 0.5 1.5 (Henning
et al. 2001; Matthews & Wilson 2002). We see this behavior also
for the regions D1 and D2. Here, the exception is the case of

Fig. 8. PI relation at a wavelength of 1.3 mm for the entire dust emission map (top left) shown in Fig. 6 and the regions of the bridge (top right),
D1 (bottom left), and D2 (bottom right). We compare the fitted trends of the PI relation Pfrac / I ↵ for the cases of RAT alignment (long-dashedred), RAT plus scattering (solid blue), and perfect grain alignment (short-dashed black). The gray-scaled counts represent the pixels of the maps in
We present these trends more systematically in Fig. 10. Here, Fig. 6 considering only RAT alignment.
we show the range of polarization caused by RATs, perfect alignment, and scattering as a function of wavelength. Yet again, we
evaluated the polarization for the entire map, the bridge, andA137, page 8 of 17
regions D1 and D2 separately. In all regions we see comparable

12

Results from Kinematic Analysis of Stars and Gas
1
1
2
MichaelA. Kuhn, LynneHillenbrand, John M.
Carpenterand Angel
AvelarMenendez
331 TheFormaBonofA.
a Stellar
AssociaBonin theNGC 7000/IC
5070Rodrigo
Complex:
Results
from1

Departmentof Astronomy,
California
Instituteof Technology, Pasadena, CA 91125, USA; 2 Joint ALMA Observatory,#28
KinemaBc
Analysisof Starsand Gas1

Alonso deCórdova
3107Lynne
Vitacura,
Santiago,John
Chile
MichaelA. Kuhn1,
A. Hillenbrand1,
M. Carpenter2 and Angel Rodrigo Avelar Menendez1
ApJ 899, 128K
E-mail contact:mkuhn at astro.caltech.edu

We examine the clustering and kinematics of young stellar objects (YSOs) in the North America/Pelican Nebulae, as
revealed by Gaia astrometry, in relation to the structure and motions of the molecular gas, as indicated in molecularline maps. The Gaia parallaxes and proper motions allow us to significantly refine previously published lists of YSOs,
demonstrating that many of the objects previously thought to form a distributed population turn out to be nonmembers. The members are subdivided into at least 6 spatio-kinematic groups, each of which is associated with its
own molecular cloud component or components. Three of the groups are expanding, with velocity gradients of 0.3–
0.5 km s−1 pc−1 , up to maximum velocities of ∼8 km s−1 away from the groups’ centers. The two known O-type stars
associated with the region, 2MASS J20555125+4352246 and HD 199579, are rapidly escaping one of these groups,
following the same position–velocity relation as the low-mass stars. We calculate that a combination of gas expulsion
and tidal forces from the clumpy distribution of molecular gas could impart the observed velocity gradients within thegroups. However, on a global scale, the relative motions of the groups do not appear either divergent or convergent.
The velocity dispersion of the whole system is consistent with the kinetic energy gained due to gravitational collapseof the complex. Most of the stellar population has ages similar to the free-fall timescales for the natal clouds. Thus,
we suggest the nearly free-fall collapse of a turbulent molecular cloud as the most likely scenario for star formation in19
this complex.

Accepted by ApJ
http://arxiv.org/pdf/2006.08622
北アメリカ／ペリカン星雲• 広がり~3度（~40 pc）、距離 795 ± 25 pc にあるHII領域（W80）、分子雲、YSOの複合体
The Mass-Size Relation and the Constancy of GMC Surface Densities in the Milky Way• primary ionizing source： 2MASS
1CharlesJ20555125+4352246
J. Lada1 and T.M.
Damestar,(Bajamar1
O3.5((f*))+O8: spectrobinary)
Harvard-Smithsonian
Center scopic
for Astrophysics,
60 Garden Street, Cambridge MA 02138 USA• OrionからCygnus Xにつながる一連の星形
E-mail contact: clada at cfa.harvard.edu
成領域の一部We use two existing molecular cloud catalogs derived from the same CO survey and two catalogs derived from13local
dust extinction surveys to investigate the nature of the GMC mass-size relation in the Galaxy. We find that the four

6.1. Relation of Clouds to Stellar Groups
There is spatial correspondence between the stellar groups andthe clouds (Figure 10). GroupsA and F are associated with theminor clouds 1 and 13, respectively. GroupE is embedded in the“Gulf of Mexico,” which consists of a major cloud component(Cloud 9) and several minor cloud components (Clouds 6, 10,
11, and 13) in the 13CO map. Group C lies on a bright rim cloudat the southeastern edge of Cloud 2. The stars in Group B are
distributed in the northern part of Cloud 2, as well as Cloud 3.

The Astrophysical Journal, 899:128 (30pp), 2020 August 20

Kuhn et al.

Figure 16. Plots of stellar kinematics for stars in Group D. Top left: vectors show motions of stars in the (x, y) plane relative to the center of the group. Stars with the mostprecise velocity estimates (<1 km s−1) are shown in black, while stars with less certain velocity estimates are shown in gray. Top right: arrow heads are color-coded by
the direction of motion of a star as indicated by the color wheel. Bulk motions such as expansion, contraction, or rotation can be seen as gradients in color across thegroup. Bottom left: distribution of the outward velocity component vout. The median value is indicated by the magenta line, and the 3σ conﬁdence interval is shown by theshaded area. Bottom right: median value of vout as a function of distance from the center of the region. An error-weighted least-squares regression ﬁt is shown.
(The complete ﬁgure set (14 images) is available.)Dambis 2020). However, Wright et al. (2019) point out that the
expansion may be anisotropic. In Appendix D, we describe alinear model for velocity as a function of position, using amethod that can account for anisotropy and is independent ofchoice of coordinate system. Table 6 reports several of theinteresting properties from the models, including velocity
gradients in the directions of maximum and minimum
expansion, the position angle of anisotropy, and the intrinsic
scatter that remains after the velocity gradient is accounted for.
We ﬁnd statistically signiﬁcant expansion for GroupsC, D,
and E, but no group had statistically signiﬁcant rotation. Group

expansion in both dimensions for GroupB, and strong
statistical evidence for expansion in only one dimension (x)for Groups C and E.
7.3. Velocity Gradients in Expanding Groups
Our results, along with the results from previous Gaia
studies, suggest that stars in expanding young stellar groups
often have velocities that are proportional to the distance fromthe groups’ centers (Kuhn et al. 2019; Román-Zúñiga et al.
2019; Wright et al. 2019; Zamora-Avilés et al. 2019; Melnik &19

The O stars are both among the fastest-moving members ofGroupD. The Bajamar Star has a velocity of 6.6±0.5 km s−1
relative to the center of the group, while HD199579 has avelocity of 6.4±0.4 km s−1. However, neither of thesevelocities is discrepant from the linear model. This suggeststhat whatever phenomenon accelerated these stars to theircurrent velocities is also responsible for the expansion of thegroup as a whole. If we make the simplifying assumption thatthe velocities have been approximately constant, the Bajamar
Star would have been nearest, in projection, to the center ofGroupD ∼1.5 Myr ago, while HD199579 would have beennearest ∼1.8 Myr ago.
7.4. Effects of Spatial Covariances in Gaia Astrometric Errors
To ensure that our measurements of velocity gradients arerobust, we examine how these would be affected by thesystematic correlated errors in astrometry (Lindegren et al. 2018,
their Section 14). They found that the covariance of quasar
proper motion as a function of angular separation decreases from∼4000μas2yr−2 at 0°. 07 separation to a local minimum of∼40μas2yr−2 at 0°. 43 separation. This could induce an artiﬁcial
proper-motion gradient of ∼0.17 masyr−1deg−1, which would
look like a velocity gradient of 0.045 km s−1 pc−1 in a stellar
association at a distance of 795 pc. Given that the velocity
gradients that we detect are signiﬁcantly larger than this value,
they are most likely astrophysical.
8. New Member Candidates from Gaia
In addition to validating previously proposed YSOs in the NAP
region, Gaia astrometry can be used to suggest new member
candidates. In most earlier studies, candidates were identiﬁed basedstrong Hα emission, infrared excess, large amplitude variability).
However, many star-forming regions are dominated by pre–mainsequence stars for which disks are not detected (Broos et al. 2013).
Thus, astrometric selection of candidates can reveal whether any
populations of NAP members have remained undetected inprevious analysis, which, in addition to their possible biases, arealso incomplete in their selection of candidate members.
We start with a sample of all Gaia DR2 sources within 3° of
the Bajamar Star that pass the same quality criteria from
Section 2.2 and lie within the same parallax range and the(m ℓ  , mb ) region identiﬁed in Section 3. However, we apply
further parallax cuts to select only objects that are consistent
within 2 standard deviations of the median parallax of thegroups, i.e., ϖ+2 σϖ 1.06 mas and ϖ−2 σϖ1.24 mas.
These criteria greatly reduce the number of Gaia sources in theregion, from ∼75,000 Gaia sources in the original parallax
range down to ∼10,000 sources.
We further reﬁne the sample to ensure that the selected objects
are likely to be young, using an age limit of 3 Myr that matches
the approximate maximum age found for members in Section 3.
We identify Gaia sources whose absolute G-band magnitude andG−RP color are 2 standard deviations above the 3 Myr Bressan
et al. (2012) isochrone with AV=1 mag (Figure 18, left).
However, remaining contaminants can include main-sequencestars with high extinction, as well as post–main-sequence stars. To
reduce contamination from reddened main-sequence stars, we use
the 2MASS J−H versus H−Ks diagram to estimate extinction
using the reddening laws from Rieke & Lebofsky (1985). Typicalreddening is Δ(J−H)∼0.3 mag (≈2.8 mag in the V band) but
ranges from 0 to 0.5 mag, with a tail out to 1.5 mag. Dereddening
on the near-infrared color–color diagram can lead to degeneracies
between low-mass and intermediate/high-mass stars. In most

15

1.
2.
3.
4.
5.

Gaiaの視差と固有運動のデータにより、~3400 NAP member YSO candidatesから395 Gaia-validated YSOsを厳選。
年齢 < 3 Myr、ほとんどは ~1 Myr かそれ以下。
6つの密なグループ。広く散らばるYSOはメンバーでない。
距離 ~795 pc だが、「メキシコ湾」のグループは ~35 pc 近く、V1057 Cyg を含むグループは ~130 pc 遠い。
YSOの各グループは分子雲の主なコンポーネントに対応。NAP分子雲の各部は（質量が十分あれば）星形成を行っている。
6. 各分子雲コンポーネントの質量：13CO mass と dynamical mass はよく一致（1E2 - 1E4 Msun)、free fall Bme ~< 1.5
Myr
7. YSOグループ間の相対運動はランダムで、complex全体の膨張・収縮はみられない。膨張するHII領域が分子雲を押している領域でもYSOの運動には影響が見られない。
8. 個々のYSOグループ（のいくつか）は急速に（0.3-0.5 km/s/pc）膨張している。膨張は非等方的→クランピーな分子
雲内の潮汐力の影響か？
9. 主たる電離放射源（Bajamar star, O3.5）はグループD起源。その速度は低質量メンバーの膨張運動と合致する。もう一つのO型星（HD 199579, O6）もグループDから放出されたようだ。
10. Gaia catalogより新たに ~1000 member candidates を同定。要 followup
Rapid star formaBon in a freefall Bme• YSOグループと分子雲コンポーネントとの空間的一致• YSOの年齢が分子雲コンポーネントのfreefall Bmeと同程度• 分子雲コンポーネントのガスの運動は重力収縮から期待される速度と一致
If star formaBon accelerates → the majority of stars formed during the last few freefall Bmescales

16