Star Formation Newsletter#361 1-15
Yusuke Aso (KASI)

Time-dependent Monte Carlo continuum radiative transfer
A. Bensberg and S. Wolf
Institut für Theoretische Physik und Astrophysik, Christian-Albrechts-Universität zu Kiel, Leibnizstr. 15, 24118 Kiel, Germany
e-mail: abensberg@astrophysik.uni-kiel.deReceived 21 December 2021 / Accepted 21 October 2022

c/Δt

ABSTRACTare briefly described in Sects. 2.3 and 2.4. An overview of different tests of the implemented routines is given in Sect. 3. The
impactof theby
temporalwidth andthe numberof photon
Context. Variability is a characteristic feature of young stellar objects thatis causedvariousstep
underlying
physicalprocesses.
Multipackages
per time step
discussed
Sects.
3.1 and 3.2.
The
epoch observations in the optical and infrared combined with radiative transfer
simulations
are iskey
to studyintheseprocesses
in detail.
cinfluenceof theopticaldepthtransfer.
on the temperature
distribution
Aims. We present an implementation of an algorithm for 3D time-dependent
Monte
Carlo
radiative
It allows one
to simulate
simulationsfor 1D and
2Dandmodelsare examinedin radiation
Sects. 3.3fortemperature distributions as well as images and spectral energy distributionsof the scatteredlight
thermalreemission
and 4.1, respectively. In Sect. 4.2, we discuss the impact of avariable illuminating and heating sources embedded in dust distributions,luminosity
such as circumstellar
disks and dust shells on time scales up
outburst of a central star embedded in a circumstellar
to weeks.
dust disk on the resulting dust temperature distribution, as well as
Methods. We extended the publicly available 3D Monte Carlo radiative ontransfercode POLARIS
efficientmethods
for theFinally,
simulathe scattered
light and withthermalreemission
radiation.
tion of temperature distributions, scattering, and thermal reemission of dust
distributions
by temporally
variable radiation
a summaryas wellilluminatedas an outlook
on further applications
andsources. The influence of the chosen temporal step width and the numberof photon
packagesofper
step
key parameters
potential
improvements
the time
methodareasprovided
in Sect. 5.for aTtstationary POLARIS

Runge-Kutta-Fehlberg

104
103

15nm
150nm
1.5 m
15 m
150 m

102
101
100
10 1

given configuration is shown by simulating the temperature distribution in a spherical envelope around an embedded central star. The
10 2
effect of the optical depth on the temperature simulation is discussed for the
spherical envelope as well as for a model of a circumstellar
2. Computational method
disk with an embedded star. Finally, we present simulations of an outburst of a star surrounded by a circumstellar disk.
10 3
In thefollowing,
we outlineof time103
101
102
Results. The presented algorithm for time-dependent 3D continuum MonteCarloradiative
transfertheis theoreticala valuablebackgroundbasis for preparatory
T [K]dependent continuum
radiativesources,
transfer.such
After
the
studies as well as for the analysis of continuum observations of the dusty environment
around variable
as discussing
accreting young
A. Bensberg and S. Wolf: Time-dependent
thermalprocesses andwavelength
properties ofrange
the dust
in Sect.
2.1, westellar objects. In particular, the combined study of light echos in the opticaland near-infrared
and thecorresponding
Fig. 1. Cooling time described in Eq. (8) for different temperatures ofdescribe
theimplementationofatime-dependent
methodoftime-dependent thermal reemission observables of variable, for example outbursting sources, becomes possible on all involved spatial grains with different radii consisting of astronomical silicate.
radiative transfer simulations for temperature, dust reemission
scales.

τ

Δt

105

Cc [s]<<tc

A&A 668, A120 (2022)àPOLARIS

and scattering in Sect. 2.2. For this, we extend the publicly available–3DMonte Carlo disks
radiative
transfercode– POLARIS (Reissl
Key words. radiative transfer – methods: numerical – circumstellar matter
protoplanetary– ISM:
cloudset al. 2016).
phstars: variables: generalN

1. Introduction
Variability is known as a characteristic feature of young stellar objects (YSOs) since
3 their first
5 observations (Joy 1945). The
ph time-dependent phenomena can reveal different
study of theseunderlying physical mechanisms such as spots on the stellar surface caused by magnetic fields, variable mass accretion onto thestar, and obscuration by circumstellar material along the lineof sight to the observer (Herbst et al. 1994). Since all of thesemechanisms have an impact on the measured flux of the YSO,
observed light curves of variable YSOs show different shapes
ranging from periodical to linear or curved (Park et al. 2021). The
time scales of the underlying event can be on the order of hours
to years (e.g., Grankin et al. 2007; Cody & Hillenbrand 2014;
Günther et al. 2014). In the case of the extreme accretion events
that can be found in FUor and EXor objects, the outburst can

N =10 , 10 stc

Under the assumption of local thermodynamic equilibrium(LTE), the dust will emit radiation with an emission rate Ė corresponding to its emission spectrum B at temperature T scaled
2.1. Description of heating and coolingwith the absorption coefficient  :in most of these solutions is that of radiative equilibrium.
Z 1
We first focus on the thermal processes and corresponding propMoreover,while
is photon
an implementationof Monte Carloerties of thedust.there
The ithpackage with wavelengthĖ(T ) = 4⇡
 B (T )d .
(6)
radiative
thatsource
uses withgasluminosity
properties
to isapproach
0emitted transfer
by a radiation
L thatemitting timeNph photon
packagesin coolingthe time-interval
energydependent
heatingand(Harriest carries2011),thethere
is currently
Considering the difference between absorption and emission rate
no solution available focusing on the thermal properties ofwithin
the one time step, that is combining Eqs. (4) and (6), we now
L t
dust.
ei, =
.
(1) can estimate the enthalpy un+1 of the dust after n + 1 time steps
N
with step width t:
In thisph paper, we present an implementation of an algo-

Δt=2s
τ

rithm
foritsfulltime-dependent
Monte
Carlo
radiative
Along
path
through the model3D
space,
the dust
absorbs
energy transfer
un+1 = un + (Ȧ Ė) t.
(7)of the photon
in relation
to the wavelength-dependent(MCRT),
whichpackage
includessimulationsof the scattered light,
absorption
coefficient radiation,
 . The totalandabsorption
rate Ȧ in adistribuHowever, this holds only true if we assume the absorption anddustmass
thermalreemission
temperature
region with mean radiation intensity J is then given by:emission rate is constant over the step width t. We can constraintion with a focus on the latter. For this purpose, we extend
Z
temporal step width by considering the time it takes to emit
the publicly1 available 3D Monte Carlo radiative transfer thecodeȦ = 4⇡1
 J d .
(2) the total enthalpy of a dust grain at temperature T , that is the
POLARIS (Reissl
et al. 2016), which is used for the simulationof time tc (T ). For the case of emission without absorption,0
coolingvarious astrophysical objects ranging from galaxies (Pellegrini
we find
Following Lucy (1999), the mean intensity J of a discrete
et al.
2020;ofReissl
al. 2020)
Bokei,globules
numberphoton etpackages
withandenergywithin a(Brauer
wave- et al. u(T )
)=
(8)
2016;
Zielinski
al. 2021)
to length
circumstellar
disksV(Brauerc (Tal.
lengthrange d ettravelinga pathli in a volume
in the tet

1Ė(T ) Fig. 3. Sketch of the time-dependent ray tracing procedure for two dif-

0.04
60

A [AU]

60

0

10 1 we performed simulations of a model for an outburst of
Lastly,40
star with a circumstellar disk. For this, we used the disk mode
10 2
described
in Sect. 4.1 (see also Table 2). For the sake of sim
0.02
plicityand to save computational time, we chose a low optica
20
10 3
depth of ⌧V = 0.01.
In thisandcase
the Time-dependent
dust temperature
thecontin
cell
A&A 668, A120 (2022)
A. Bensberg
S. Wolf:
MonteofCarlo
areafter about 50 000 s, that is roughly the time ligh4
10 converged
_
[m]
0
0.00
65000s
85000s105000s
needs totravelonce through the model space.85000s
The luminosity o
65000s
60
180
0.15
the
central star was then increased
by a factor of 4 for a dura
75 tion of 10 000 s. The total simulation time was set to 120 000
40
20
0.02
to0.10make sure, that all grid 160
cells return to their previous dus
50
temperature after the outburst.
20
25
140 was performed using a tempora
40
The temperature simulation
0.05
5
0.04
0
0 step width of 4 s and 10 photon packages per time step. Th
120 such that the cooling time of th
temporalstep width was chosen
0.00
60
3
25
0
20
40
60
80
100
10 7highest dust
10 6temperature
10 5 is not reached
10 4 within10one
20
expectedstep
A [AU]
100 _ [m]50 Three vertical cuts through the temperature distribution of th
40
disk
model
for
different
time
steps
duringtheoutburst
areshow
Fig. 9. Difference plot of a vertical cut through the dust tempera- 75 Fig. 10. Normalized SED (top) and relative difference (bottom) of the
80 that a heat wave corresponding t
in Fig. 11. It caneasily bedistribution
seenture distribution of the disk model with ⌧ = 10 described in Sect. 4.1. time-dependent
temperatureof the disk model with ⌧ = 1060
in0thesourceluminosity is traveling0 throughthe gri
The20 reference
temperaturewiththe80 stationary
040
60
80 distribution
100
20is calculated
40
60
100
20the outburst
4050in Sect.
60
discussed
4.1 80(dotted).
is calculated
50
50 100 The reference50SED (solid)
A [AU][AU]in radial
direction.distribution
Furthermore,thethermalxmethod.
response times omethod using 108 photon packages.
using
the temperatureof thestationary
theanddust
arescatteredsmall,
thatmodel
is thedust of a cell is heatin
Fig. 11. Vertical cut through the temperature distribution of the outburst disk model after 65 000,
105000
scomparably
described
in Sect.
The
Fig. 85
13.000
Imagesof thelightof the4.2.
outburst
after 65 000, 85 000 and 105 00
temporal step width of the simulation is set to 4 s; 105 photon package are emitted per time step.
or While
cooling
the equilibrium
temperature
withinone
time step
width: 1000s;
108 photon
packagespertime ofstep).
flux isisnormalizedwith respect
to th
theto overall
shape
theTheSED
well
reproduced,
the100
steps. leading to a discrete form of the heat wave.
difference plot reveals a peak difference of about 15% for waveFurthermore,investigatedthetoscatteredstellarlight anthe afterglow in the thermal
reemission
case20iswe
small
compared
lengthsof aroundµm.
This
is closethe shortest
wavelength
dust.
While the550nm - face-on
10 1
1.0
1.0light
10 mdust- face-onas observables of thi
to the one seen in the
curves ofthe
the thermal
scattered radiation
light.
The reareemissionof theatthermal
which
reemission
radiationof thements
dustofforthe
550nm - 25
crystallm - 25son for this is that the relative
cold Inouterparts of thewedisksimulated
are not 10 thesystem.
particular,
thermalreemission
0.9
(see,
Laordiffer&radiDr
given
temperature
distribution
becomes significant.
The
10 2
significantly contributing
to themeasured reemission
radiationfrommeasureme
ation
of
the
N
band
(10
µm)
and
the
scatteredstellarlight
0.8 enceplotsThisin alsoFigs.
8 and
9 show that the dust temperature ofo
flux before or during the outburst.
causesthe emission
0.8
glass and SiO
2 g
the
V
band
(550
nm)
using
the
respective
time-dependent
algo
10 3
theface-on
cells disk
in the
inner
few
AU of
the midplane is underestimated.
peaks of the inclined andto beclosertogetherthan
opticalrithms for ray tracing and scattering described inbetween
Sect.
2.3
an
0.7
in the case of scattering.
dust
tempera0.6 These cells correspond to the cells with the highest mationof the tem
4Sect.
2.4. In thewith
casethethat
the observed
diskTheis inclined,
it of
is
nec
10
comparedto thetheturedeterminedstationary
method.
spectrum_ [m]0.6
essaryto takedifferent oflighttothe observe4.3. 2D disk outburst –thermalHigh
opticaldepththe radiation
therefore
improvreemission
thetraveling
dust withtimes
the time-dependent
0.4
0.15
Finally,
we
0.5
into
account.
To illustrate
thisofeffect,
intensity
maps werte
temperature
distribution
is thus
shifted theto longer
wavelengths
To test the treatment of
backwarming,thatis the heating
ulatingan outbu
derivedforaninclinationof 25of aswith
welltheas fortheface-oncas
regions of the disk by
the thermal
reemission
radiationcomparedtothespectrumcalculated
stationary
method.
0.4
thetemperature
0.10
0.2 repeated(0 ). Sincetheonlyscatteringandsimulations
areasnotsendifferent regions, wethe time-dependent
temperature
However,
thisaffects
thereemission
radiation of
wavelengths
correas welltherm
simulations of the disk modelSect.
4.2 with
anwidth,optical the step width was
0.3
sitiveusedto thetemporalstep
setwithtoThe
1000
culated.
sca
sponding
toin
thethermalreemissionof the few cells
the
depth
of
⌧
=
100
in
the
midplane.
In
this
case,
the
outer
parts
8
0.05
V
were
and a totalof photon
packagesthe
perflux
timeofstepofsimulated
10
0.0 highest
dustnumber
temperature.
Therefore,
mostof werthe0.2of the midplane of thedisk are shieldedfromthe direct emission
20000
40000
60000
80000
100000120000inclinationof
usedscatteringand reemission,respectively.
Inanorder
avoi
40000
80000
120000wavelengths
largerthan
10µm is 100000
showing
not morethanto 10%of the star and can thus 20000
only
be forheatedby
the60000
thermalreemissiont [s]as the time-depe
t [s]
0.00
an
impactcut
theinitialheating process, all simulations
werof the surrounding
dust.difference
A verticalthroughthe temperature
to of
the
reference
SED.
7
performed indep
10 6light (550 nm)
10of5 the outburst10 4
10 3
Fig. 12. Light curve of 10
the scattered stellar
distribution of the
high
depth
disk
modelreemission
after
000
sreached.
startedonce
athermalsteadystate95was
an offset o
Fig.
14. optical
Light
curve
of theradiation(10 µm)Toof rulethe out
distribution.
_ [m]model discussed in Sect. 4.2. The duration of the outburst
is 10 000 s; during the outburst
is shown
inlight
Fig.
16.
It in
canSect.
bethere
is andoutburst
model
discussed
4.2. thatThe
durationof thereemitted
outburst is radiation
thecurvesof seenthe
scatteredcause
The effectofthe total simulation time is 120 000 s. During the outburst, the stellar a warm region in10
the000
midplanearound70 AU
thatis enclosedby the outburst, the of the dust was ss; bythe additional
total
simulationtime
is 120000
s. Duringlight
traveltimes
duetodifferent
positioningof thluminosity isFig.
increasedby a factor ofSED
4. Theflux andis normalizedwith colder
10. Normalized(top)
relative difference(bottom)themeansstellarluminosity
increasedfactorof 4.emisThe flux is normalized of the outburst m
regions.ofThis
it wasisnot
heated by
by athestellarrespect to the
maximumoftheintegrated
flux
of
all
imagesofthedifdetector
inthescattering
andraytracing
treatments,
the
start
o
time-dependent temperature distribution of the disk sionmodel
withoutburst,⌧ =with
10 respect
maximumof the integratedof thebut by tothethethermal
reemissionof the flux
dustof all images of the midplane.
ferent time steps.
time outburst
steps.
the
60 000s in both cases.
discussed in Sect. 4.1 (dotted). The reference SED (solid)
is calculated
above andbelowdifferent
the midplane.
It shouldwas
alsosetbe to
noted,
that theIt should be nwidth and position of the warm
spot
do not corresponddirectly
using the temperature distribution of the stationary method.
The
resultinglight curvefor the scattered lightition isof shown
temperatu
thinnerlayersthethedisk.
This is time of scattering
packages can be found in the upper half of the third scattering to the outburst intothebeopticallyFig.compared12.
A samplecorresponding
imagessmall
to theofof
shortest
cooling
the sys- fromhoursfortothw
caused by the delay
betweentheheatingbythestellaremission
image.
the numberofphotonpackages
time
stepThere
has to arestellarsystems
inclineddisk
canbeasfound
inper
Fig.
13.
three
effecto
While the overall shape of the SED is well reproduced, tems,
the whilst
reemissionof the
dust,
as
well
the density
The light curve of the simulation of the thermal reemission and the thermal beenoughcorrect
absorption
physicsfluxof thi
canforbe
seenscatteringin the andlightcurves:treatment.
First, the theoveralldifference plot reveals a peak difference of about 15% for wave-large thatradiation is shown in Fig. 14. The corresponding images of the structure of the disk.
Since all photon packages in the grid have to be loaded from and While this is ou
higherdisk.
caused provide
by th
µm.
close
to thereemissionlengths
radiationof
arearoundgiven in20
Fig.
15.ThisThe ismain
effectvis- shortest wavelength
stored ona stackinin theeachcasetime ofstep,thetheinclined
access speed
andThissize is methods
different
ofare
thecrucial
scattering
direction.
scatterwhich
thermalradiation
dust for oftheible in the at
shapeof thethelightcurves isreemission
the light travelingtime of5.thethe conclusions
CPUcache asprobabilities
well as the RAM
parameters
for icalTheeffects
canDiscussion
andof the geometricaldepth structure distributionof the disk. Whilethe heating
the totaling
computation
Even
thoughthe treatmentof the sin- is time-dependentof the time.
stellarlight
towards
the observer
more probabls
given temperature
becomessignificant. The differand coolingencetime plots
scales are
the temporal
stepdust
Wetemperaturepresented angle
implementationanalgorithm
fulltimein POLARIS
parallelized,
the computation
images theare inscat
pri
forpackages
the ofinclined
thanisforfortheface-ondisk. Secondly,in short
Figs.compared
8 and 9toshowthat theofphotonwidth, the heating of each layer of the disk in the reemission light dependent 3D Montefor the MCRT
time Carlois thusradiative
not necessarily
decreasing
linearwith the number of tion, they can alr
teredlighttransferoftheoutburst
iscodefirst
visible in the light curve othevisible
cells within theinnerrelativefew AUof theunderestimated.
curve will be
an offset
to thelightmidplane
traveling isPOLARIS.
It includes
time-dependent
methodsforthe simulation
threats.
Temporal step
width andnumberof photon packages per of the thermal reinclined
disk.
This shows
light travel
cells
correspondcells withhighestdust temperatime withinThesethe disk.
Combinedwith tothethedifferent
light thetravelof temperature
distributions
andthermalreemission
aswell theas importance of thetime stepthe
shouldtherefore
be chosen
accordingly.
dependent
tempe
ing
time
the observer.
Since distributions
the lower parthydrodynamics
of the disk isture
the stationary
method.
spectrumof thetimes due tothedetermined
inclination ofwith
the disk,
the light curveof the The
scattering
images.
The
for toward
the simulationoftemThemethod
time-dependent
simulationof thetemperature
inclined diskshowsabroader
peak
than
the
face-onlight
curve.
perature
distribution
focuses
for
the
first
time
on
the
thermalwas successfully
testedisbycloser
simulating
1D observer,and 2D models
withEvenbethoughthe images
to thethe outburst
willsee
thermal reemission radiation of the dust with the time-dependent
It should betemperature
noted that thedistribution
duration of theisoutburst
relativelytoto longer
propertiesof dustdifferent
insteadthere
ofoptical
gas. first.
Itdepthcan beforthe0.1,
simulation(⌧used0.01,
1, 10, 100).
A deviation
temperature
V =is
This
also
visiblein thefirst imageof Fig.distr13
thus shifted
wavelengths
the light travel times within the grid is a crucial factor for the of the radiative transfer
of reference
various astrophysicalfrom thetemperature objects.
distributions calculated with the structure of the uFurthermore,theoutburst
isstretchedout
compared
to
the
face
comparedtothespectrumcalculatedwiththestationary
method.
shape of the emission peak, both in the case of the thermal reeWe discussedstationary
the influenceof the of
newly
introduced
parameapproach
POLARISwas only
seen for the case of afford increasesoncase,because
parts
of
the
disk
with
the
same
radialdistance
However,
this
onlyaffectstheradiationofwavelengths
corremission radiation and for the scattered radiation. There is also an ters temporal step⌧Vwidthand numberof photon
packages
per inconsistency to represent thet
= 10. t The
main reason
for this
is a slightafterglow visible
in bothHowever,the durationof fewtimecells
step Nwas foundthat
the will
temporalwidth hasthehave
different
distances
Thisopticca
the source
appliedopticalandstep
calorimetricpropertiestoofthe
theobserver.
sponding
tolight
the curves.
thermalreemissionof thewith
thewith a highph . Itbetween
highest dust temperature. Therefore, the flux of most of the be seen in the second of the scattering images. The third effect i
A120, page 10 of 14

105000s
180

40

160
140

T [K]

I [AU]

20
0

120
20
100
40
80
60

0

20

40

60

80

100

20

40

A [AU]

60

80

100

20

40

60

80

100

Fig. 11. Vertical cut through the temperature distribution of the outburst disk model after 65 000, 85 000 and 105 000 s described in Sect. 4.2. The
temporal step width of the simulation is set to 4 s; 105 photon package are emitted per time step.

1.0
0.9
0.8

550nm - face-on
550nm - 25

the afterglow in the thermal reemission case is small comparedto the one seen in the light curves of the scattered light. The reason for this is that the relative cold outer parts of the disk are notsignificantly contributing to the measured reemission radiation
flux before or during the outburst. This also causes the emission
peaks of the inclined and face-on disk to be closer together thannear/far sidedT/Ty [AU]

T [K]_85000s再放射

Norm. Flux

65000s60

Relative Diﬀerence

Norm. Flux

A&A 668, A120 (2022)

𝑖 = 0°
𝑖 = 25°

Relative Diﬀerence

I [AU]

4_ [norm.]

10000s_50000s

100

10
4.2.
2D disk outburst
0.04

I [AU]

τ=0.01

80

SED

60

中心星光度が外へ伝わる40

Fig. 9. Difference plot of a vertical cut through the dust temperature distribution of the disk model with ⌧ = 10 described in Sect. 4.1.
The reference temperature distribution is calculated with the stationary
method using 108 photon packages.
A. Bensberg and S. Wolf: Time-dependent Monte Carlo continuum radiative transfer

τ=10Δt=8s, Nph=10520_ [norm.]

stationary0Articlehttps://doi.org/10.1038/s41550-022-01831-8

UV-driven chemistry as a signpost oflate-stage planet formation

Jenny K. Calahan 1 , Edwin A. Bergin 1, Arthur D. Bosman1, Evan A. Rich1,
Sean M. Andrews2, Jennifer B. Bergner 3, L. Ilsedore Cleeves4,
Viviana V. Guzmán5,6, Jane Huang1, John D. Ilee 7, Charles J. Law2,
Romane Le Gal8,9,10, Karin I. Öberg 2, Richard Teague11, Catherine Walsh 7,
David J. Wilner 2 & Ke Zhang12

Received: 19 July 2022

Accepted: 14 October 2022

Published online: 8 December 2022
Check for updates

Received: 19 July 2022

>1 mm

Accepted: 14 October 2022

Published online: 8 December 2022

CH CN, HC3N

Check for3 updates

Jenny K. Calahan 1 , Edwin A. Bergin 1, Arthur D. Bosman1,The
Evanchemical
A. Rich1, reservoir within protoplanetary disks has a direct impact on2
3
4
Sean M. Andrews , Jennifer B. Bergner , L. Ilsedore Cleeves planetary,compositions and the potential for life. A long-lived carbon- and2
Viviana V. Guzmán5,6, Jane Huang1, John D. Ilee 7, Charles J. Law,
nitrogen-rich chemistry at cold temperatures (≤ 50 K) is observed within
7
Romane Le Gal8,9,10, Karin I. Öberg 2, Richard Teague11, Catherine,
coldWalshand evolved
planet-forming disks. This is evidenced by bright emission
2
12David J. Wilner& Ke Zhang<10 µm

誘

CH3CN

J

1–10 MyrK
aged systems that would otherwise
K laddar from small organic radicals in25-50

have frozen out onto grains within 1 Myr. We explain how the chemistry ofa planet-forming
The chemical reservoir within protoplanetary disks has a direct
impact on disk evolves from a cosmic-ray/X-ray-dominated regime
toaultraviolet-dominatedchemical equilibrium. This, in turn, will bringplanetary compositions and the potential for life. A long-lived carbon- andabout awithin
temporal transition in the chemical reservoir from which planets will
nitrogen-rich chemistry at cold temperatures (≤ 50 K) is observed
accrete.
This photochemical dominated gas phase chemistry develops as
cold and evolved planet-forming disks. This is evidenced by bright emission
dust evolves
via growth, settling and drift, and the small grain population isfrom smallotherwise
3 organic radicals in 1–10 Myr aged systems that would
depletedfromhave frozen out onto grains
within 1 Myr. We explain how the chemistry of the disk atmosphere. A higher gas-to-dust mass ratio allows
https://doi.org/10.1038/s41550-022-01831-8
for deeper
penetration of ultraviolet photons is coupled with a carbon-richa planet-forming disk evolves from a cosmic-ray/X-ray-dominatedregime
gas
(C/O
>
1)
to a ultraviolet-dominated chemical equilibrium. This, in turn, will bring to form carbon-bearing radicals and ions. This further results ingas phaseformationof organic molecules, which then would be accreted by
ISM
about a temporal transition in Approximatethechemicalreservoir from whichplanets
will
number density
any
actively
formingplanets present in the evolved disk.
accrete. This photochemical dominated gas phase chemistry develops as

UV

àC/O

C2H, CH CN

ArticleEarly-stage planet formation <~1 Myr

Molecules with ALMA at Planet-forming Scales(MAPS) structures that are thought to be indicative of planet formation . The
compositions
their icy
mantles
1-2 Myr IM
Lup of pebbles
CH3andCN,
HC
3Ndirectly influence thefinal composition of a solid planetary core . After a core becomes
sufficiently
planetsCN,
the gaseous
1-2 Myr AS
209massive,CHHC3material
N fromz/r>0.1
3 start to accrete
reservoir surrounding the pebble-rich midplane of a protoplanetary
disk to form CH
their atmospheres
directlyprobe the>6 Myrladder3CN . It remains difficult toK
chemistry of actively forming planets. The solid cores of giant plan- gas within the planet-forming midplane owing to the high dust densi<10–3 cm–3

UV
Late-stage planet formation >~1 Myr

dust evolves via growth, settling and drift, and the small grain population isoutdepleted fromFreeze
the disk
atmosphere.
A arehighergas-to-dust
massforratio
allowsProtoplanetarydisks
the natal
environments
planets.
Disks
10–1 –10–3 cm–3
Numberfor deeper penetration
ultraviolet
photons
is coupledwitha carbon-richhaveofthreemain components:
a pebble-rich
dusty
midplane(dust grain
density of1> ~1 mm), a gaseousatmosphere
extendingwell above
and radismall
grainsgas (C/O > 1) to formradiuscarbon-bearing
radicalsandions.
This further
resultsin2
ally beyond 10(by–1 a factor–3 of ~2; ref. ) the pebble-rich midplane and a–10cm
gas phase formationsmallof organic
molecules, which then would be accreted by
dust population (radius < 10 µm) that is coupled to the gas. Eachany actively formingcomponent
planets present
in the evolvedof the protoplanetary
diskdisk.
has an impact on shaping the

à

4

5,6

7

–3

>10 cmets must form over a short timescale (~1 Myr) for the eventual planet
4 in atoobtain its
fullaremass
over thecourseof a typicallifetimeof gas. The
structuresthatthought
to beindicativeof planet
formation
3
disk
(3–10
Myr;
ref.
)
and
to
explainthewidelyobservedgap
andcompositions of pebbles and their icy mantles directly influenceringtheties leading to elevated dust optical depths that mask line emission, as
well as cold temperatures that lead to the freezing out of common gas
Protoplanetary disks are the natal environments for planets. Disks
tracers such as CO onto dust grains. However, constraining the chemical
have three main components: a pebble-rich dusty midplane (dust grain
UV flux
1
5,6
radius > ~1 mm), a gaseous atmosphere extending well above and radi- final composition of a solid planetary core . After a core becomes
Planet coreally beyond (by a factor of ~2; ref. 2) the pebble-rich midplane and a sufficiently massive,planets start to accrete material from the gaseous
Freeze out andsmall dust population (radius < 10 µm) that is coupled to the gas.
Eachreservoir
surrounding
the pebble-rich midplane of a protoplanetary
photodesorption
1
3
Dust grainAnn Arbor,
University
oftheir
Michigan,USA. 2Centerfor∣ Harvard
component of the protoplanetary disk has an impact on shaping the disk
to form
atmospheres7. ItMI,
remains
difficult
toAstrophysics
directly probe
the & Smithsonian, Cambridge, MA, USA. University of Chicago Department
4
5
of
the
GeophysicalSciences,
Chicago,USA.
Departmentof dustAstronomy,ig. 1 | A schematic chemistry
highlightingof
thephysicalevolutionof a diskTheand how
small
dustplansettles, UV
photons
make
their
way deeper
into
the IL,
disk,
allowing
actively
formingplanets.
solid cores of
giant
gas
withinthe planet-forming
midplane
owing
tofor
the high
densi- University of Virginia, Charlottesville, VA, USA. Instituto de Astrofísica,
6
7
Ponticia
Universidad
CatólicadeChile,
Macul,
Santiago,
Chile.
NúcleoMileniodeFormación
Planetaria(NPF),
Santiago,
Chile.
School of Physics
hat physical environment
canaffecttheachemistry.
At the top,
showfora the eventual
photodesorption
complex
species
off grains.dust
Now,optical
there is depths
a cycle ofthat
carbonets must
form
over
short timescale(~1weMyr)planet ofties
leading
to elevatedmask line emission, as
8

9

C

CH3CN

I (mJy beam–1 km s–1)

15

g/d

exp tailµm
1 mm
HD163296
Hya CH3CN, HCN, HC3NC
10
C/O=1-2
1 Myr
UV

1
TW

15

CH3CN J=

120–110
121–111
122–112
123–11310300

CH3CN J=

60–50
61–51

10

0
0

100

200

0

100

Radius (au)

200

r z/r

0.40

–2
–3

z/r

100

200

Radius (au)while dashed lines correspond to modelled radial profiles. Our fina
includes an increase in C/O ratio beyond 20 au and a depletion of sm
represents 1 Myr of chemistry. This model can simultaneously fit th
ratios and general morphology of the observed radial profiles of thorganic molecules.

1

HC3N

0

0.30
0.25

–1

0.20
–2

0.15
0.10

–3

0
0

50

100
Radius (au)

150

200

Fig. 3 | The radial and vertical number density distributions of CH3CN and
HC3N. The density distributions were determined by our final HD 163296 modelwith an increase in C/O ratio beyond 20 au and a depletion of small dust. This
Articledistribution produced the radial profiles shown in Fig. 2. A z/r < 0.1 is a good

Base model

15.0

12

12.5

11

Depleted insmall dust

0

50

100
Radius (au)

150

200

approximation for the midplane of the disk, and both CH3CN and H
partially from the midplane. A comparison can be made with Suppl
Fig. 6, which lacks the inclusion of small dust depletion.

https://doi.org/10.1038/s41550-022

15.0

12

5.0

12.5

11

Height (au)

G0 depleted/G0 baseof planet formation and its direct impact
on the chemis
2.0
10.0
ing planet
atmospheres. Within the early stages of plane107.5
pebble growth7.5and accretion into solid1.0planet cores has
7.5
99Discussion
chemistry
onto these cores are influenced primarily by t
5.0
5.0
5.0
0.5
HC3N and CH3CN are built up in the gas from simple carbon and make-up of large pebbles and their icy mantles.
The chem
88
2.5during this stage of formation is dominat
2.5
2.5 (see Methods section ‘Chemical
nitrogen-based volatiles
reactions’ in the midplanerays49. After timescales of orderfor details). A reservoir of these more simple molecules is maintained and cosmic
1 Myr, at the
0.2
7
0
0
7
0
in the gas phase because0of photodesorption
from dust grainsurfaceshave grown and settled
4002040pebble formation,
0 small
20grains40
20(au)(au) midplane, allowing forRadius(au)by an enhanced UV field. The Radiusmain destruction
products fromRadius
CH3CN
UV photons
to penetrate deeper inthatcanwithincycle backto create larger
nitrileincreasepreviously
assumed.
This excess
UV fluxplot
in concert
with an
3N are simple
Fig.and
4 | A HC
comparisonof thevolatiles
MRI active
regiona non-elevatedUV
in mass thatis MRI active.
The right-mostshows a compar
molecules
includingtheirparentmolecule.
Althoughthere areUV field
C/Ostrength
ratio producescarbon-rich
phase
cycle
of Theproenvironment
andan elevated
UVoriginalenvironment.
The
two left-mostplots show
through a a‘regular’disk andgas
a dustdepleted
disk.
log10 (ngas)

higher dust opacity than modelled. A comparison model without a
10.0
10.0
depletion in small dust mass is shown in Supplementary
Fig. 5.
10
log10 (ngas)

UV

0

0.05

0.05

12.5

300

log10(HC3N per cm3)

–1

log10(CH3CN per cm3)

0

0.25

200

0.40

1

0.20

15.0

100

Radius (au)

0.35

CH3CN

0.30

0

0

Radius (au)Fig. 2 | Radial intensity profiles of the observed complex organic moleculesand HCN towards the disk around HD 163296. The molecules shown are CH3CN,
HCN and HC3N (left to right with two J transitions of CH3CN shown in the firsttwo panels). Solid thick lines in the background correspond to the observations
derived from Ilee et al.12 and Guzmán et al.79 (utilizing the results of Law et al.53),

0.35

0

0

0.10

1 Myr

15

200

100

0.15

C/O
C/O>1
X
UV

HC3N 11–10

50

UMIST, Rate12

HC3N 29–28

20

10

5

5

HCN 3–2
HCN 1–0© 2022. The Author(s). Published by the American Astronomical Society.

The JCMT BISTRO-2 Survey: Magnetic Fields of the Massive DR21 Filament
The JCMT BISTRO-2 Survey: Magnetic Fields of the Massive DR21 Filament

3,4
Tao-Chung Ching1,2,93, Di Li2,5,62,5,6, Zhiyuan Ren2 2, Shih-Ping Lai7,8 7,8, David Berry9 9, Kate Pattle10 10,
1,2,93, Keping Qiu
Tao-Chung Ching, Keping Qiu3,4 , Di
, Zhiyuan Ren, Kate,11,12
14,15 , Shih-Ping Lai
8 , David Berry
16,17 Pattle
Ray Furuya 11,12, Derek Ward-Thompson13 Li, Chang Won Lee
13, Doug Johnstone
14,15, Patrick M. Koch
16,17,Ray Furuya, Derek Ward-Thompson, Doug,21Patrick M. Koch8 , Chang
16,17
18
19,20 Johnstone
2,22 Won Lee
8,
Thiem Hoang 16,17, Tetsuo Hasegawa 18, Woojin Kwon 19,20, Pierre Bastien 21, Chakali Eswaraiah 2,22, Jia-Wei Wang 8,
Thiem Hoang, Tetsuo Hasegawa, Woojin16,17
Kwon, Pierre Bastien, Chakali Eswaraiah, Jia-Wei Wang ,Kyoung Hee Kim16,23, Archana Soam24 24, A-Ran Lyo16 16, Junhao Liu9 9,
16,23, Jihye Hwang
Kyoung Hee Kim25,26, Jihye Hwang16,17 18, Archana Soam
, A-Ran, Junhao Liu 14,15,
27 LyoValentin J. M. Le Gouellec 25,26, Doris Arzoumanian 18, Anthony Whitworth 27, James Di Francesco 14,15,
Valentin28,29
J. M. Le Gouellec, Doris Arzoumanian, Anthony 33,94
Whitworth, James,
30
31,32
34 Di Francesco
Frédérick Poidevin 28,29, Tie Liu 30, Simon Coudé 31,32, Mehrnoosh Tahani 33,94, Hong-Li Liu 34, Takashi Onaka35,36
35,36,
Frédérick, Tie
Liu, Simon Coudé, MehrnooshTahani, Hong-Li40Liu, Takashi Onaka,
37 Poidevin
18,36,38
39
37
41,42
Dalei Li , 37
Motohide Tamura 18,36,38, Zhiwei Chen 39, Xindi Tang 37, Florian Kirchschlager 40, Tyler L. Bourke 41,42,
Dalei Li , Motohide
Tamura,Zhiwei
Chen
,
Xindi
Tang
,
FlorianKirchschlager
,
Tyler
L.
Bourke,15
Do-Young Byun16,17
Huei-Ru Vivien Chen7,8 7,8, Wen Ping Chen43 43, Jungyeon Cho44 44, Yunhee Choi16,16
16,17, Mike Chen , 15
Do-Young Byun
, Mike, Wen Ping44Chen, Jungyeon, Yunhee Choi,
45
16 Chen , Huei-Ru Vivien
41Chen2 Cho
46Youngwoo
Choi
, 45
Minho
Choi
, 16
Antonio
Chrysostomou
41, Eun Jung Chung
44, Y. Sophia Dai
2, Pham Ngoc Diep
46,
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Choi
,
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Choi
,
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Chrysostomou
,
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Chung
,
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Sophia
Dai
,
Pham
Ngoc
Diep
,
47
2
7
48
8,49
50
51
Yasuo
DoiDoi
Duan
Laura
M.M.
Fissel47, Yan Duan , Hao-Yuan
2
7, David Eden
48, Lapo Fanciullo
8,49, Jason Fiege , 50
51,
Yasuo
,
Yan
Duan
,
Hao-Yuan
Duan
,
David
Eden
,
Lapo
Fanciullo
,
Jason
Fiege
,
Laura
Fissel
,
50
9
52
42
53
9
27
9, Rachel Friesen
52, Gary Fuller
42, Tim Gledhill
53, Sarah Graves
Erica
Franzmann
Friberg
Greaves
Erica
Franzmann2750, Per, Per
Friberg, Rachel Friesen, Gary Fuller, Tim Gledhill, Sarah Graves9, Jane
,56,57,95
JaneGreaves27, ,
30
16,17
54
55
The Astrophysical Journal, 941:122 (21pp), 2022 December 20
Ching et al. 16,17
54, Martin Houde
55, Charles L. H. HullMatt
Grifﬁn, 27
Qilao
GuGu30, Ilseung
Han
HayashiMatt
Grifﬁn, Qilao, Ilseung
Han , Saeko, SaekoHayashi, Martin, Charles13L. H. Hull56,57,95,16 ,58
58 (21pp),59
16,60Houde
The Astrophysical Journal, 941:122
2022 December 20
al.
58, Shu-ichiro
58, Kazunari
59
16,60, Vera Könyves Ching
13,etJi-hyun
TsuyoshiInoue
Inutsuka
Iwasaki,
Il-GyoJeong
Kang
Tsuyoshi Inoue 16 , Shu-ichiro Inutsuka, Kazunari Iwasaki, Il-Gyo Jeong, Ji-hyun
Kang16, ,13 Table 1
1861,62,63 , Vera Könyves
64,65,66
16, Janik Karoly
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18, Koji Kawabata
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Miju
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Notes.
aJinjinXie,
Meng-Zhe
Yang,
Hsi-Wei
Yen,
HyunjuYoo,
Jinghua
Yuan,
Hyeong-Sik
Yun,
Tetsuya
Zenko
Figure, 3.Jinghua
Histogram of Yuan
position angles(PA) between theYun
magnetic ﬁeld, Tetsuya Zenko, ,Jinjin Xie , Meng-Zhe, Hsi-Wei, Hyeong-Sik
With respectYang
to the pointing
center at (α, δ) Yen= (20 39 , 1.Hyunju
1, 42n21a17´Yoo).
bsegments
ofsubﬁlaments
andthecrestsof
subﬁlamentsinFigure1.
Aposition
2,792,79
86
2
37
79
40
87
86
2
37
79
40
absolute position angle between PA and PA in a range [0n, 90n].
Chuan-Peng
Zhang, Yapeng
Zhang
Zhang
, Jianjun
Zhou
, the
Lei
Zhu
Ilse
Looze
, Philippe
AndréChuan-Peng
Zhang The
, Yapeng
Zhang , Guoyin, GuoyinZhang, Jianjun,magnetic
Lei
Zhu, IlseLooze, Philippe
André87, ,angle
ofZhou0° means thatﬁeld is, paralleltode
the desubﬁlament
crest,
anda position angle of 90° means that the
88 88
89 89
90 90
91magnetic
91 ﬁeld is perpendicular to the 92 92
The Falle
Astrophysical941:122(21pp),
2022
December 20 , andC. C.
DarrenDowellEyres
Falle, Jean-François
RobitailleSven
vanvanLoo
DarrenDowell, Stewart, Stewart
Eyres, Sam, Sam,Journal,
Jean-François
Robitaille
, and
Sven
Loo
subﬁlament
crest.
1 1

DR21,
4 pc,
15000 Msun, d=1.4 kpc, 24
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Stokes I

a

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f

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|B−f|

f

Ching et al.

are inclined
in a northeast–southwest
orientationinHangzhou
theHangzhou
ResearchCenterforﬁelds
Intelligent
ComputingPlatforms,
Zhejiang
Lab,Lab,
311100,Peopleʼs
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chingtaochung@gmail.com
ResearchCenterforIntelligent
ComputingPlatforms,
Zhejiang
311100,Peopleʼs
Republicof China;
0.82chingtaochung@gmail.com± 0.03 (Ngoc et al. 2021), Rosette: 0.49 ± 0.08 (Könyves
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et al.Peopleʼs2021),
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5 5
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Departmentof Astronomy,
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addition,theorientationUniversity
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Figure1,
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F1,
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6 6NAOC-UKZN
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University,
Hsinchu
TaiwanHennemannetHuaal.
(2012).
In Table
1, we30013,
list30013,
thepositions,themodern grain alignment theory (Lazarian &InstituteofAstronomy
DepartmentofPhysics,
National
Tsing
Hua
University,
Hsinchu
TaiwanAccording to
8 8 magneticﬁelds
aroundmassiveofdense
cores primarily
follow),thepositionanglesofpositionanglesofmagnetic
ﬁelds
(PA
B
Academia
SinicaInstitute
Astronomy
andAstrophysics,
No.1,
Sec.3.,
Roosevelt
Road,
Taipei
10617,
TaiwanAcademia
SinicaInstitute of Astronomy and Astrophysics, No.1, Sec.3., Roosevelt Road, Taipei 10617,
TaiwanHoang
2007; Hoang & Lazarian 2016; Hoang et al. 2021), dust
the horizontal
magnetic ﬁelds of the ﬁlament, except for the), and Hilo,the absolute
positionangles(PAaligned
subﬁlaments(PAfPark,
9 9 East|B−f|)
Asian
Observatory,
660 A’ohōkū
N. A’ohōkū
Place,
University
96720,
USA
grains are
by radiative torques, and the grain alignment
East Asian
Observatory,
Place,
University
Hilo,
HI HI
96720,
USA
orientedmagnetic 660
ﬁeldsN.aroundDR21
and PA
segmentsofbetweenPABPark,10
f of the 13 magnetic ﬁeld
10 northeast–southwest
towardthehighestDepartmentof Physics
andAstronomy,
University
CollegeLondon,
Gower
Street,LondonWC1E
6BT,
UK intensity is caused by internal radiation from
DepartmentofPhysics
andAstronomy,
University
CollegeLondon,
Gower
Street,London
WC1E
6BT,
UK
(OH).
The northeast–southwest
orientationofthemagnetic
pixelsof crests
subﬁlaments. The PAf is determined by the ﬁve themassive
centralstar. As a result, the embedded sources of
11 11 Tokushima University, Minami Jousanajima-machi 1-1, Tokushima
770-8502,
Japanﬁelds aroundTokushima
DR21(OH) areconsistent Minami
with the Jousanajima-machi
small-scale
University,
Tokushima
770-8502,
Japanthat are1-1,
closestto the magnetic
ﬁeld segment.
Figure3 shows
the
may increase the alignment efﬁciency in
DR21
and DR21(OH)12 12 Instituteof
LiberalArts
andSciences
Tokushima
University,
Minami
Jousanajima-machi
1-1,
Tokushima
770-8502,Japanmagneticﬁelds
inferredfrominterferometricobservationsof.
The
histogramofthepositionangleshas
histogram
of
PA
|B−f|
Institute of Liberal 13
Arts and Sciences Tokushima University, Minami Jousanajima-machi
1-1, Tokushima the
770-8502,Japan
high-density
regions, producing a shallower α than those13
dust polarization
(Lai
et
al.
2003;
Girartetal.
2013),
andwemoresamples
between
0°
and
45°
than
between
45°
and
90°,
Jeremiah
Horrocks
Institute,
UniversityofCentralLancashire,
Preston
PR1
2HE,
UK
Jeremiah Horrocks Institute, University of Central Lancashire, Preston PR1 2HE, UK found in clouds without embedded sources.
14
that
the distortionof the andmagneticﬁelds aroundindicatingthat themagneticﬁelds tend
toV9E
be paralleltothe crests
14speculate
NRC
Herzberg
Astronomy
Astrophysics,
5071
West
SaanichRoad,
Victoria,
BC
2E7,
CanadaNRC
Herzberg
Astrophysics,
5071
West
Road,
Victoria,
BCthe
V9E
2E7,
CanadaDR21(OH)
could beAstronomy
driventhe andnortheast–southwest
bipolarofSaanichdifferent
fromperpendicularalignment
15 15 Department
of byPhysics
andAstronomy,
University
ofsubﬁlaments,
Victoria,
Victoria,
BC
V8W
2Y2,
CanadaDepartmentofPhysics
andAstronomy,
UniversityofVictoria,
Victoria,
BC
V8W
2Y2,
Canadabetween the magnetic ﬁelds and the DR21 main ﬁlament. The
outﬂows of DR21(OH) (White et al. 2010; Zapata et al. 2012;16 16 Korea Astronomy and Space Science Institute,
3.2. Comparison
Daedeokdae-ro,
Yuseong-gu,
DaejeonRepublicof Korea between POL-2 and SCUPOL Results
FigurePolarization
fractionp asYuseong-gu,
aalignment
functionof Stokes
I intensity.
The34055,
grayparallelbetweenmagneticﬁelds
andsubﬁlaments
Korea
Astronomy
andScience7764.776
Daedeokdae-ro,
Daejeon34055,
Republicof Korea
et al. 2013).
AtSpacethe southernend Institute,of DR21(OH),theﬁeld
17Girart). The black34113,
and red Republic of Korea
dots
represent
the Gajeong-ro,
low-intensity dataYuseong-gu,(I < 0.5 Jy beam−1Daejeon
Figure 2. Comparison of the dust polarization maps between POL-2 and SCUPOL observations. (a)17
The POL-2
polarization
map
of
the
main
ﬁlament.
The
grayscale
Universityof
ScienceandTechnology,
Korea,
217
revealedin
our POL-2
data
is in34113,
agreementwith theFigurecomparison
morphology
isbeamslightly
northwest–southeast
oriented
2(b) compares the polarization maps of our POL-2−1−1
University
ofJyScience
Technology,
217 along
Gajeong-ro,
DaejeonRepublicof
Korea) of the northﬁlament
dots represent
the high-intensityYuseong-gu,
data (I 0.5 Jy beam
180.25, 0.5,
. Theandmagneticﬁeld segments Korea,
represents the Stokes I intensity, and the contours show the Stokes I emission at levels of 0.125,
1, 2, 4, 8, and 16
18
andPlanck
data
that
trace
STokyosubﬁlament
andSCUPOL
National
Observatoryof Japan,National
Institutes
ofHerschelNaturalSciences,
Osawa,
Mitaka,181-8588,
Japandata of Matthews et al. (2009). The
the Astronomical
connecting
bridgebetweenDR21(OH)
andDR21.
The respectively.
data
with
the
are the same as those in Figure 1, but plotted with a uniﬁed length. The triangles mark the positionsof the massive
dense cores in Motteetal. (2007),withand southﬁlament,
The best-ﬁt
modelsofthep–I correlations
ofthe
National
Astronomical
Observatory
ofDR21(OH)
Japan,
National
Institutes
of ofNaturalSciences,
Osawa,
Mitaka,Tokyo
181-8588,
Japan
19
in the north and DR21 in the south, highlighted in white color. The dotted line marks the19boundaryDepartment
between thenorth ﬁlament
and
south
ﬁlament.
Theorangeandcyanmagneticﬁelds
atalarger
scale
(Hu
et
al.
2021).
the
north
and
south
ﬁlaments
are
shown
by
the
black
and
red
lines,
noise
level
in themagneticﬁeldmorphology
along
the
connecting
bridgeisof Earth Science Education, Seoul National University, 1 Gwanak-ro, Gwanak-gu, Seoul 08826, Republicof SCUPOL
Korea Stokes Q and U maps is about↖
I missing flux

α

I

α=0.3

d

The Ncell

We next calculate the gradiusing the Gaussian derivativ(2013). To obtain gradients
Planck magnetic ﬁeld segme
3 × 3 derivative kernel ov
magnetic ﬁeld segments. Be
computed over two FWH
Planck data, obtaining gradi
adequate sampling of gradie
Figures 8 and 9 show thedensity maps and the maps
The majority of POL-2 seg
tend to be parallel (f = 90
N(H2)  1023 cm−2. In the l
JCMT map, the alignmen
gradient segments becomes
magnetic ﬁeld segments an
map appear to be more ran
segments in the JCMT maangle of the gradient is denoise in the column density
2016b). Because the respe
Stokes I map and the Planck
percent of the I map and τ
gradient directions are typica
criterion of p/δp 3 forcorresponding to an uncert
angle (Naghizadeh-Khouei
expect that the errors in f a
We divide the 765 measur
ﬁve N(H2) bins and the 371
data into three N(H) bins toplot the HROs of the JCM
bins each of 15° width. Th(12)Figure 8. Comparison of the magnetic ﬁeld segments and column density gradient segments of the POL-2 data. (a) N(H2) column density map derived from the JCMT
Stokes I emission overlaid with the gradient segments calculated by convolving the column density map with a Gaussian derivative kernel. The contours show the N(H2) at levels of 0.125, 0.25, 0.5, 1, 2, 4, 8, and 16 × 1023 cm−2. The length of the gradient segments is normalized. The gradient segments shown here are thoseoverlaid with the magnetic ﬁeld segments in Figure 2(a). (b) The map of relative orientation angle f between the magnetic ﬁeld segments in Figure 2(a) and the
HRO
HRO
gradient segments in panel (a) here. The contours are the same as panel (a),with the contourat 5 × 1022 cm−2 emphasized to show the transition from no preferential
orientation of
low-density regions to perpendicular orientation of f in high-density regions.
23 f in−2

1/2

2

Y  CHRO [log10 (N (H) cm2)  XHRO].

where T 2A0 and T 2A90 represent the variances of the areas,characterizing the “jitter” of the histograms. The value of ξ isnearly independent of the number of angle bins selected torepresent the histogram if the bin widths are smaller than the
integration range, but the jitter does depend on the number ofangle bins in the histogram. If the jitter is large, σξ is large
compared to |ξ|, and the relative orientation is indeterminate(Planck Collaboration et al. 2016b).
Finally, analyses of HROs characterize the trend of therelative orientation between magnetic ﬁelds and column
density structures of a cloud from its low- to high-density
regions with a linear regression between ξ and atomic gas
column density N(H) (Planck Collaboration et al. 2016b):(11),

A 90 )4(A 0

A 02 T 2A90 )(10)

A 0  A 90,
A 0 A 90
−2

2 2
T A0
4 (A 9022

Ching et al.

k

T 2Y 

22

where A0 is the area under the histogram in the range
0° < f < 22°. 5 and A90 is the area under the histogram in therange 67°. 5 < f < 90° (Soler et al. 2017). An HRO peaking at
f = 0° would have ξ > 0, an HRO peaking at f = 90° would
have ξ < 0, and a ﬂat HRO would have ξ ∼ 0. The uncertainty
in ξ is obtained from

353

Y

⎜

⎟

h
T 2k  hk ⎛1  k ⎞ ,h tot ⎠⎝

variance in the kth histogram bin is given by

The Astrophysical Journal, 941:122 (21pp), 2022 December 20

where hk is the number of samples in the kth bin and htot is thetotal number of samples (Planck Collaboration et al. 2016b).
To evaluate the preferential relative orientation in eachcolumn density bin, the shape of the HRO is quantiﬁed using ahistogram shape parameter ξ, deﬁned as

U353 N (H) 

The Astrophysical Journal, 941:122 (21pp), 2022 December 20(9 )(Planck Collaboration XI et

for the main
adopt column densities of 41.6 × 10 cm
the relative orientation angle f between the magnetic ﬁeld andhandrasekhar& with aspatial
diffuseemission
scaleslargerHoude
etal.
(2009)
suggestthat,
22 if−2the correlation length δ2jTheturbulencecomponentintheangulardispersion
functionisﬁlament
and
∼2
×
10
cm
(note
this
value
is
consistentwiththe tangent to the column density contour is evaluated usinga(4 )Guerra The
et al.2021)(Liu ettheal.
2021).
2j l , and numerical simulations
sed
method. region.
densityderivedfrom
thePlanckdata
inFigure
9)
for
the
%a
is
much
smallerthanthethickness
of
the
cloud
,
the
e observed
Becausethediffuse
for
B
t
j1
2
al. that(2009)
suggest
that,
if
the
correlation
lengthδ fromfromthetheangular
diffuseregion obtained
DR21 Herschel map
B q N ⎞§ B
B1
2 0 2canlencemagnetic
dispersion
ratio
oflamenttoisHoude
strongeretthanof NGC
1333,
G  arctan ⎛,(8)t
2 (l ) 2 (EBW 2(Hennemanntl 2to). be evaluatedet al. 2012)
with %a  0.34 pc to derive the
b
(
)
e
6
B · N ⎠
⎝
%a
is
much
smallerthanthethickness
of
the
cloud
,
the
for
B
is thepolarization
anglethandifference
between(l ) ﬁlamentf%'variationfunction
DR21
hence is larger
the Ncelltin theB02§ in the formnumber densities n of the POL-2 and Planck maps, usingρ dispersion= Journal,
μmHn whereμ = 2.86 is the mean molecular weight (Kirk
where B is the magnetic
orientation inferred from theation
segments
separateddistancel, W is the
beam
/V
δB/B
; δVInlosof
C 1333.
Figurethea polarization
The Astrophysical(21pp), 2022 December 20
Ching et ﬁeld
al.
B0 by
can
be evaluatedfrom
thevsangular
ratio
B
ΔΦ
𝑙 941:122
A,t to5(b),2 et al.22015) and mH is the atomic mass of
et al. 2013;
Pattle2B§polarization
map
and
∇N
is
the
gradient of column density,
1
2
2
2
t

B
§

B
§
Because
Btthoseis by
theofsource, the tis andthevariation
2δB
data
are
preferentiallylargerthan
8 ln 2of(i.e.,
the
FWHM
beam
), perturbationhydrogen.
ments
columndensity1the  cos [%'in(l )]B§0
q0 [1  el 2(E 2W )]functionin theformdivided
used to characterize the column density structures. Although2
We estimate
from the velocity dispersion (σ) of theNdensity
map resultsN(Hderivedfrom Equation(4) provides
a good approximation
cell Bthe
gontrary
the
line
of
sight,
thegeneralof POL-2
data2) column
0 §δVlsrof
the range of arctan function is [90n, 90n], we use a range
tion
is
atoTaylor
expansion
representing
the structure
in
13
+
H
CO
1–0
data
at
an
hdegrees
the gradientsegments
d 13 + angular resolution of 29″, because the2[0n, 90n] for f without loss of generality as suggested in Soler
ensityρ.
Because
the
δB/B
in
the
DCF
equationforevaluatingthemagneticslightly
smaller
than
those
of
of H CO 2jis well-correlated with the dust emission
1 isBthe2)
does
not involve
turbulence, andposNcellt § number of l 2 2(E 2 2Wemission
witha Gaussianderivativeaet(2017), because the relative orientation is independent of
a
l
,
(
4
)

2j (Schneider
1 Doi
cm
cos
%'
(l )]
§
q [1of
] DR21 ﬁlamentin thej=1 et al. 2010). The derived Bpos theal.reference−2skycomponentofmal. 162018;
et. [al.
2020).
Again,onthe
ﬁeldstrengththe2 planeskye as4, et8,cellsand× 1022and thus f is equivalent to −f. The convention of
nt
along
the
line
of
sight
obtainedby
j1strengths
are
∼0.6
mG
in
the
north
ﬁlament,
∼1.0
mG
in
the
Ncell are
B 0 §
kyThemagnetic
zationdegreesﬁeldof the DR21 ﬁlamentd.
gradient
segmentsf
isequivalent
to the |90° − δ| in Koch et al. (2013) that f = 0°
south
ﬁlament,
and
∼0.1
mG
in
the
diffuseregion.
Our
value
of
1 2
2
d
indicatesthatthemagnetic ﬁeld is parallel to the tangent of theeldsegments
in Figure(6.E 2 I ﬂux) %athe POL2W 2in
the
north
ﬁlamentisconsistent
with
the
SCUPOL
resultsof§Btothe missing
measuredlarge
Stokeswherepolarization (7angle
difference
between column density contour%'⎡ (tl ) ⎤is the
13
+
2
j

(
)
N
.
5

QS
E
B
4
V
.
)
(perpendicular to the column density
Bpos = 0.78 mG
derivedusingCOthe DCF1-0
method (Equation (3))
e f between thecell
magnetic
posa a l ,
los
H
(
4
)
2segments3
2j Astrophysical∼δΦ.
ConseBts⎥⎢
polarization
separatedbyadistance
l,
W
is
the
beam
The
Journal,
941:122(21pp),
2022
December
20
0pos
Ching
et
al.
Q
E
2
gradient),
and
f
=
90°
indicates that the magnetic Chingﬁeld iset al.
§Bin Vallée & Fiege (2006) and Bpos = 0.62 mG derived from thein the top panel.
⎣ 0 ⎦ angularj1dispersion
functionanalysisinGirartetal.
(2013).
perpendiculartothetangent
ofthecolumndensitycontour# by"beam divided
8 ln 2 ), the (parallel to the column density gradient).
width(i.e., ξ, the
FWHM
!our
"!
However,valueandis (11),
about six times Table
weaker2 than the
Figure12.
Relativeorientation
parameterdeﬁnedinEquations(10)bulencecomponentintheangulardispersion
function
is
!
Table
2
4.
Analysis
Δ':
,
W:
,
δ:
,
a:
,
~
2.8–3.9mGJCMTin Poidevin
et al. Dispersion(2013)theusing
the structure
To obtain an HRO, the gradients of a column density map
is a Taylor
expansion
representing
structureinAngularFunction
Fit Parameters
Dispersion
Function
Fit Parameters
calculated
for the different summationbins ofand(blue)
data
where
theAngularpolarization
angle
difference
between%'(l ) 2is4.1.2.
!#" ofData!
Dispersion N(H)
Function
ofthe
thePlanck
JCMT(red)andPlanck$%& et al. (2009), mainly owing tofunctionanalysis
Hildebrandandthe magnetic ﬁeld segments of a polarization map are ﬁrst
DR21 ﬁlament. The
JCMT resultsof ξ in thec Serpens
Main regionareNBt § ofl 2 the(3)1 Function
turbulence,andisnumberofngularDispersionb
a B0 that doesb not involve
2 (E 2n 2W 2 ). DataΔ'
2 assuming
Poidevinetal.2§(2013)
acell
value
ofthe
n Figure
that
is 10
times
comparedpixel
pixelN(H)to binproduceNext,
the map λcNn thebyhighest
δVlosa a map of f.Bpos
δaNsegments
δVB
λa2 b 2 (l ) <B>
11.
HROs
for
the
lowest,theintermediate,
and

B

B
§
a

(
)
e
6
cell
apolarization
cell
los
pos
separatedbyadistance
l,
W
is
the
beam
0and HHRO
shown ingray dots. The−1black dashedline andfunctions
the valuesofof theCt HRO
2
2 Figure
−2
−3
−1
−2
−3shows the angular7
dispersion
JCMT(gray,
blue,
and
red,
respectively)ofthePlanckdata.
Thehorizontal
dashedlarger
than
ours.
Our
B
for
the
south
ﬁlament
is
lower
by
a
of
f
is
divided
into
bins
of
columndensities
containingan
pos obtained
cells(pc)along
theof datasightby )line corresponds to the average HRO per angle bin(cm(arcsec)(km s )
(mG)
(arcsec )(cm ) to the(kms )turbulent(mG)
Ncell B0 § correspondof 15° for an N(H) bin. The
linearﬁt of Equation(12)
fortheofline
JCMT
below
factor
three
than
the
values
from
2.5
to
3.1
mG
derivedusing
4.1.1.
Formalism
equal
numberofsegments,
and
an
HRO
is
generatedforeachof the shaded areas for each histogram correspond to the ±1 σ
23
−2
lnDCF
2 ),0.49
width(i.e.,
thebeam
dividedby 8−3thetheand FWHM
Planckdata
towardthe
DR21 ﬁlament.
Becausethemain350 μm polarizationwidths−5. 1.0North
uncertainties
relatedto theto
histogram
operation.
methodfromdataby(Equation
Kirby± 0.17(3.6 ± 0.3)
× 10
6.1:(9))±2.6bin×binning
105the± 0.18column 3.9± 10) × 103.9
examinechange1.0
in f with 0.63
increasing
2Ching2 ) Journal,
.6
±2022
0.3)
×be
10−5∼0.5,
6.1 ± 2.6 N(H)
4.0=×10105 cm POL-2
0.63± 0.18
2
2 (51
med
to
p),
December
20
et
al.
The
Astrophysical941:122(21pp),
2022
December
204.0
A
,
A
0°,
90°
(
E
%a
2
W

B
§

B
§
e
B
is
the
sourceof
perturbation
in
0
90
−3
−5
5 with small uncertainties in column densities
featuresofhorizontal
segments
in
the
north
of
the
ﬁlament−5
5 B0, thePOL-2
t(2009).
The
difference
between
the
two ×
works
is primarilyt
0 (117densities.
For
maps
POL-2
± 21)
× 10
0.18
±
0.02
(3.8
±
0.7)
10
1.4
±
0.4
4.0
×
10
1.0
1.04
± 0.132.4
magnetic
strengthinmolecularsummationaTaylorrepresenting
thestructurein.8 ± 0.7)
×ﬁeld
10
1.4is
± 0.4
4.0 × expansion
10clouds
1.0South1.04± 0.13
2.4

(
)
N
.
5
cell
−1
along
the
−5
4angles, the typical propagated error in f is
3
owed to Kirby(2009)
usingavelocity
dispersion
of
4.2
km
s
−5 line(4) provides
4 approximation
andpolarizationarenotablydifferent
from
those
of
the
radialPOL-2
segments
in
dlds
from
Equation
a
good
of
1.6
±
0.9
2.0
×
10
0.7
0.13
±
0.04
Planck
1.02
±
0.16
0.27
±
0.16
(4.0
±
0.6)
×
10
QlineE toward
2mass
± 0.9
0.13
± 0.04to the 0.9± 0.6) × 10
Tablemass-to-ﬂux
2 2.0 × 10
Table
2 by presenting an HRO with angle 0.9
n.0 observations,the1.6Davis–Chandrasekhar–
ratio0.7λ, whichrefersratio
the4–3in athe DR21 core, which is about
from
theof
HCN
usually less than 10°.
Hence,
B
that
does
not
involveturbulence,
and
N
is
the
numberof0
cell
olutionof
dust
Angular
Dispersion
Function
FitofParameters
AngularFit Parameters
the
south
the
we magnitude
perform theofanalysis
separately
for/Bpos inDavisthe
DCF
equationfor evaluating
the tomagnetic
four times
largerthanwhat we used. From Equations (4) andbins of aDispersion
width largerFunction
than the error
in f, the uncertainty in the
DCF;
1951;
Chandrasekhar&tubemagnetic
ﬂuxﬁlament,themagnetic
ﬂux(CrutcherThe
turbulencecomponentintheangulardispersion
functionis22
Notes.
HRO
isdominatedbythehistogrambinning process. The(5),
isproportional
to
,
1


cos
[
%'
(
l
)]
§
1
%a
q

B
§

B
§
1;
Ostrikeretal.
turbulent
cells
along
the
line
of
sight
obtained
by
1/2
the
POL-2
segments
in
the
north
ﬁlament(Figure
7(a))
and
in
the
t
0
c
a
b
2
2 mostrengthonthe planeofasmethod.
onisthewidelyusedetNcell Then
δVlos13critical
Bpos
λ
Ncell
n
δV
Data
a In units, al.δ2010).
al. 2004).
of itsvalueof (Schneider
2πG+Bt2§ B02§
a 2a skya 2at § B0 §Valuesfrom
single-dishHCO
1–0
observations
etal.B2010).
−2
−3
−1
−2
−3
2
south
ﬁlament(Figure7(b)).
The
numbersofsegment
pairs
reach

21

2

1
(arcsec ) M to7.6
(cm
(mG)
(pc)
(arcsec )(cm )(km
Bt §theory
12
1 estimate
2 2 (E of
2 2W
2) B .
mes
magnetic
[2
NW
(H10%
cm (km][Bsin Nn)2 G] δV. losIn
the
2 b Assume22))) %a2E10
1q
uncertainty
and
to
the
in
luncertainty
ty in that
Bpos. the ratio of turbulence
pos
2
(
b
(
l
)

(
)
e
.
6
l

30
a
Q dataBpos
fromsimplyfortheamaximum
at
l
=
80″
for
the
and
at
−5
5 POL-2
c
−3
−5
5

§
B
t ⎤
⎡variation
0.49 ±to0.17a similar(3.6 ±
0.3) ×of
10
6.1 ±4.0 ×with
10 Bstar
1.0
0.63± 0.18
3.9
lead
level
in2.6
0.49 ± 0.17
(3.6 ± 0.3) × 10
6.1 ± 2.6
4.0 × 10
1
POL-2
(51
± 10) × 10
B.
Obtainedmagnetically(Shuet(025§ﬁeld
al.
1987;
pos
) segments
Ncell.)  formation
NorthB
ofcellthe magnetic
and column density
4 Figure 9. ComparisonN
 the
EVlos
Bpos of
4QS. thedominated−5
5 (7angular−3−55eﬁnements
Planckdata,
implyingthat
thedispersion
functionsare3
2) column
map10
gradient segments 1.04
of POL-2
the Planckdata.
Top panel:2.4
N(H(117
0.18as± 0.02(3.8 ± 0.7) × δB/B
10
1.4 ± /V0.4 &1.0
± 0.13South±density21) ×
0.18 ± 0.02
(3.8 ± 0.7) × 10
1.4 ± 0.4
4.0 × 10
1
⎢Mouschovias
wellin
the velocities,
2Q4.0E× 101999),
cloudsaretheinitiallymagnetically
§⎥
B0;
los
A, Ciolekmap
overlaid with the gradient segments
derived from
Planck τ
mass-to-flux
4
−5
4
⎦δV
⎣−5λ
2 0.27 ±20.16
1.6 ± below
0.9
2.0 ×
10
0.7
0.13
±
0.04
0.9
0.27analyses± 0.16(4.0 ± 0.6) ×fully
10
1.6
±
0.9
2.0
×
10
0
Planck
1.02
±
0.16
(4.0
±
0.6)
×
10
cal
have
sampled
80″
for
the
JCMT
map
and
fully
sampled
calculatedbyconvolving
thecolumndensitymap
withaGaussianderivative
BecauseBt is
the asourceof perturbationin B , the Bt § B0 §th of the magnetic ﬁeld, δBsubcriticalis the2 variation(λ < 1), andto becomestar-forming
2, 4, 8, and 16 × 10 cm0 .
kernel. The contours
show the N(H ) at levels of region,
2
2
2
a
ﬁt
theangularfully
sampleddata
points,theparametersareand hence%aThe
lengthofour
thefunction
gradientand
segments
is isnormalized.
The gradient a
segments
a 2aj aresvelocity
andarecoupled.
Consideringthatand hence
line
B30′
§ of
for
Bsight,posed
a dispersion
structurebelow
thePlanck
map.
Here,
we
focus
on
the
fully
2japproximation
t λ>2
0§
The
turbulencecomponentinthedispersion
derivedfromEquationprovidesgood
of %a and Bt § B0 § are coupled. Considering that our
along
the
shown
are those overlaid
withthe magneticfor
ﬁeld asegments
in Figure6.
magneticsupercriticality
(λ >
1) of
aherecloudis(4)requiredtheNotes.
assumption of %a  0.34 pc of the diffuse region is anpanel:
offunction
relativeorientation
betweenthe magnetic
data
0.34points.
assumptionof of%athepc
ofThe
thePOL-2
regionis TheanmapatheaBottombecauseﬁttingrange
issmall.
The the magnetic
reducedtodiffuse
ﬁrst
order
ifference
betweensampledangularinanglethefinstars
Angular
Dispersion
Function
JCMT
andPlanck
Data
mall.
The
2adispersion
λ~1à
ﬁeld
segments
inDCF
Figure 6 andthegradient
segments
thetop panel.
isthe
Alfvénspeed
at density
the
δB/B
in
the support
equationforevaluating
self-gravity
themagnetic
andform
2ρ. Because
2 to 2overwhelm
vations(Schneider
etal. 2010).
pos Therefore,
Valuesfrom
single-dish
H13CO+ 1–0
observations(Schneider etB
al.2§2010).
underestimated. Therefore,B02§(10)isandalso
underestimation,isalso
underestimated.

B
§

B
§
best
ﬁts2smallerof theangulardispersion
functions
are shown in Figure 12.underestimation,hownin
t in Equations
t
0 is
given
map
as
a
north
ﬁlamentslightly
than
that
in
the
ﬁlament,b south§
B
Relative orientation parameter ξ, deﬁned
(11),
1
2
2
t of 
estimatetheuncertainty
in Bpos. dispersion
Assume10%
uncertainty
in n and δVlos
to ourestimatetheuncertainty
in
B
.
ments
trace
the
plane-of-sky
component
of
2
l
2
2
W
(
E
)
through
gravitational
collapse.
Using
thestatistically
most
ﬁeld
strength
on
the
plane
of
sky
as
pos
re
7
shows
the
angularfunctionsthe
JCMT
calculatedforthedifferent
N(H)
bins
of
the
Planck
(red)
and
JCMT
(blue)
datavalue
of 0.13 mG in the diffuse region could be an upper
our value bofindicating
mG indiffusecould
bethe
upperare(0.13l) (6with) Binpos Tablee region. anparameters
Figure7, andthe inﬁtted
are listed
2. The of the DR21 ﬁlament.
e 2. The functionQ
c
The JCMT results of ξ in the Serpens Main region areQ
structure
thatthe
the
ﬁelds
north
more
ordered
B.
Obtained2magnetic
ariationintheplane-of-sky
magnetic
ﬁeld(Crutcher
et
al.
2004)
and
the
probable
value
of
B

B
and H
shown
in
gray
dots.
The
black
and thevalues of C data.
4
anck
data
towardtheDR21
ﬁlament.
Becausethemain
pos lengths
limit of B ofdashedinlinethe
Plancklimitof Bposthan
in thePlanckdata.
NinBsouth.
correlation
δ of thePOL-2
north segments,2
1 POL-2
cellthe4to limited
correspond to the linear ﬁtpos
Equation (12) for the JCMT data below
20§s,d POL-2
assumed
to
be
those
Owing
angularresolution,the§ 7⎤ 5 ± 1 5, N(H) = 10 cm .
Bt (seebe proportional
to column
thein themeasureddensityobtained
Herschelmap
s5of
horizontal
POL-2 segments
northof theﬁlament from⎡ aresouth
segments,
and BthePlancksegments±to1 5,
QS EVlos.
(7 )pos  42
2
rest of the JCMT N(H) bins are broadly negative, indicating a
5. Discussion2
2
ndfor
thedifferent
highestN(H) bins∼δΦ.
Consezation
angles,
i.e.,
δB/B
⎥
⎢
2B
2.5
acoupled.
o1.4a3B±λ2§3values
1,B 2and,
respectively
(see
Table
2
for
the
17
ably
from
of
the
radialPOL-2
segmentsin
4.1.2),
the
of
the
main
ﬁlament
anddiffuseposSection
2nd
mass-to-ﬂux
ratio
λ,parameterswhich refers alignment
to theratioof the themass inand
a

B
§

B
§
Becauseisthesourceof
perturbation
in
B
,
the

§
B
a10
aj between%astrongpreferenceforperpendicularat2thosethetheparameters
are
and
are
Considering
that
our
and
hence

§
%a
a
ﬁt
the
fully
sampleddata
points,andtheareandFieldhence
Btin §the DR21
B02Filament§ are co
0
0
t
0
j
t
0
2
⎦
⎣
5.1.
The
Role
of
MagneticThehorizontal
dashedline
magnetic
ﬂux
tube
to
the
magnitudeofmagnetic
ﬂux(Crutchersmall-scale
magnetic
ﬁelds
and
the
ridge
of
the
DR21
ﬁlament
4.2.
HistogramofRelative
Orientations
uationcanbewrittenasδ
inparsec).
Except
for
the
POL-2
north
ﬁlament,
the
4.2.
HistogramofRelative
Orientations
ent,
the
th
of
the
ﬁlament,
we
performtheanalysisseparately
for
regionare listed
2. of
Given
the uncertainties
inorder
the a a becauseet(parallel
0.34
assumption
of %ain Tablepc
the diffuse
regionto%a required 0.34pcthe roleofalignment
betweenﬁelds value
and density
parameters
to evaluate,One of the importantof
al. 2004).
In rangeunits ofmagnetic
criticalof 2πG assumption
ﬁtting
range
The
the
ﬁttingisitssmall.
The
reducedfromEquation(4)
providesa goodapproximation
ofisﬁrstan
bin.
The
15°
forderivedanN(His) small.
2 N
22
2 is the 2dimensionlesscorrelation
lengths
are not resolved by the beams. The 7.6 q 1021[N (H2) cm2][B NG]1. In the theory ofof magnetic ﬁelds in star formation
cell Mgradients).

LETTER TO THE EDITOR

2

Observatorio Astronómico Nacional, IGN, Calle Alfonso XII 3, 28014 Madrid, Spain
3
Observatorio de Yebes,Discovery
IGN, Cerro dePalera s/n, 19141
Spain
oflainterstellar
NCYebes,
NH+ :Guadalajara,
Dicyanopolyynes
are indeed
4?

in space
Received 17 November 2022 / Accepted 13 Decemberabundant
2022

M. Agúndez1 , C. Cabezas1 , N. Marcelino2,3 , R. Fuentetaja1 , B. Tercero2,3 ,P. de Vicente3 , and J. Cernicharo1

ABSTRACT

1

Instituto de Física Fundamental, CSIC, Calle Serrano 123, 28006 Madrid, Spaine-mail: marcelino.agundez@csic.es; jose.cernicharo@csic.es
2
Observatorio Astronómico Nacional, IGN, Calle Alfonso XII 3, 28014 Madrid, Spain
3
Observatorio de Yebes, IGN, Cerro de la Palera s/n, 19141 Yebes, Guadalajara, Spain

The previous detection of two species related to the nonpolar molecule cyanogen (NCCN), its protonated form (NCCNH+ ) and one
metastable isomer (CNCN), in cold dense clouds supported the hypothesis that dicyanopolyynes are abundant in space. Here we+
report the first identification
in17spaceof2022NC. December
This cationis the protonated form of NC4 N, which is the second member ofReceived
November/ Accepted2022
4 NH 13the series of dicyanopolyynes after NCCN. The detection was based on the observation of six harmonically related lines within theYebes 40m line survey of TMC-1 QUIJOTE. The six lines canABSTRACT
be fitted to a rotational constant B = 1293.90840 ± 0.00060 MHz and+The
previous
detectionoftwospecies
relatedtothenonpolarmolecule cyanogenits protonatedform (NCCNH) andone+ based on high-level aba centrifugal distortion constant D = 28.59 ± 1.21 Hz. We confidently
assign(NCCN),
this seriesof lines
to NC
4 NH
metastable isomer (CNCN), in cold dense clouds supported the hypothesis+that dicyanopolyynes are abundant in space. Here we
initio calculations, whichreport
supports
the previous
of HC
from
of a series of lines with a rotational
5 NH
the first identificationin spaceidentificationof NC4 NH+ . This cationis theprotonatedform theof NCobservation
4 N, which is the second member of
+
+1.4
the
seriesofdicyanopolyynes
after
NCCN.
The
detection
was
based
on
the
observationofsix
harmonicallyrelatedlines within
thecm 2 , which implies that
constant 2 MHz lower thanthat
derivedhere.
Thecolumndensity
of
NC
NH in
TMC-1
is (1.1± 0.00060
1010
4rotational
0.6 ) ⇥MHz
Yebes
40m
line
surveyof
TMC-1
QUIJOTE.
The
six
lines
can
be
fittedtoaconstant
B=
1293.90840
and+NC4 NH+ is eight times lessabundant
than
NCCNH
species
CNCN,
previously
toward
L483aband tentatively in TMC-1,
a centrifugal
distortion
constant
D = 28.59 .± The
1.21 Hz.
We confidently
assign thisseries of linesreported
to NC4 NH+ basedon high-level
initio calculations, which supports the previous identification of HC5 NH+ from the observation of a series of lines with a rotational
is confirmed in this latter constantsource.
Welower
estimatethat here.
NCCN
anddensity
NC4ofNNC
arepresent
in TMC-1
with2 abundances a few times to one order
+
+1.4
10
2 MHz
than that derived
The column
4 NH in TMC-1 is (1.1 0.6 ) ⇥ 10 cm , which implies that
+
+
NC
NH
is
eight
times
less
abundant
than
NCCNH
.
The
species
CNCN,
previously
reported
towardL483
and tentatively
TMC-1,of magnitude lower than HC4 3 N and HC5 N, respectively. This means that dicyanopolyynes NC(CC)n inCN
are present at a lower levelis confirmed in this latter source. We estimate that NCCN and NC4 N are present in TMC-1 with abundances a few times to one orderthan the corresponding monocyanopolyynes
The thatreactionsof theCN
andwith HNC arise as the mostof magnitude lower than HC3 NHCC
and HC5 N,(CC)respectively.
This means
dicyanopolyynes
NC radicals(CC)n CN arepresent
at aC
lower
n CN.
3 N level
thanthecorresponding
monocyanopolyynes
HCC
(CC)
CN.
The
reactionsoftheradicals
CN
and
C
N
with
HNC
arise
as
the
most
n
3
likely formation pathwayslikely
to formation
NCCNpathways
and NC
cold dense clouds.
4 N into NCCN
and NC N in cold dense clouds.
4

Key words. astrochemistry – line: identification – molecular processes – ISM: molecules – radio lines: ISM
Key words. astrochemistry – line: identification – molecular processes – ISM: molecules – radio lines: ISM

well known that monocyanopolyynes H C⌘C (C⌘C) C⌘N,
Yebes 40m, TMC-1
more commonly known as cyanopolyynes, are abundant in cold
Dicyanopolyynes are molecules with a linear unsaturated skele- dense clouds and circumstellar envelopes around evolved stars,
well known that monocyanopolyynes H C⌘C (C⌘C)n C⌘N,
1. Introduction
NCCCCNH+ NCCNH+
ton of carbon atoms 8that terminates at eachNCCN
edge by a NCCCCN
and thus it is conceivable that dicyanopolyynes are abundant

1. Introduction

cyano group, that is to say N⌘C (C⌘C)n C⌘N. The simplest member of this family, cyanogen (NCCN), has been longknown to be present in the atmosphere of Titan (Kunde et al.
1981; Coustenis et al. 1991; Sylvestre et al. 2020) and has been
recently identified in the coma of comet 67P/Churyumovn The second member of theGerasimenko (Hänni et al. 2021).
series, dicyanoacetylene (NC4 N), has been observed in solidstate, although not in the gas phase, in the atmosphere of Titan(Samuelson et al. 1997; Jolly et al. 2015). It is also interesting to note that NCCP, a species chemically related to NCCN
that results from the substitution of one N atom by one P

Dicyanopolyynes are molecules with a linear unsaturated skeleton of carbon atoms that terminates at each edge by acyano group, that is to say N⌘C (C⌘C) C⌘N. The simplest member of this family, cyanogen (NCCN), has been longknown to be present in the atmosphere of Titan (Kunde et al.
1981; Coustenis et al. 1991; Sylvestre et al. 2020) and has been
recently identified in the coma of comet 67P/Churyumov-nmore commonly known as cyanopolyynes, are abundant in colddense clouds and circumstellar envelopes around evolved stars,
and thus it is conceivable that dicyanopolyynes are abundantas well. However, dicyanopolyynes are nonpolar and thus
cannot be detected at radio wavelengths. In spite of that, inrecent years we have learned that dicyanopolyynes are very
likely abundant in interstellar clouds. The protonated form ofas well. However, dicyanopolyynes are nonpolar and thus
cannot be detected at radio wavelengths. In spite of that, inrecent years we have learned that dicyanopolyynes are very
likely abundant in interstellar clouds. The protonated form of
cyanogen, NCCNH+ , was detected in the cold dense clouds
TMC-1 and L483 (Agúndez et al. 2015). Just as N2 H+ is used
to probe N2 (Linke et al. 1983), the detection of NCCNH+ wasinterpreted as indirect evidence of the presence of NCCN, forwhich a large abundance in the range (1–10) ⇥ 10 8 relative to
H2 was inferred. A second piece of evidence in support of thepresence of cyanogen in interstellar clouds was provided by

Received 12 August 2022 / Accepted 28 December 2022
Number of stars in the Sun’s birth cluster revisited
1,2 and Eiichiro Kokubo2
Sota Arakawa
ABSTRACT
1

Japan Agency for Marine-Earth Science and Technology, 3173-25, Showa-machi, Kanazawa-ku, Yokohama, Kanagawa 236-0001,

26
The Sun is thought to have been
Al-rich and 26 Al-poor calcium–aluminum-rich
Japan formed within a star cluster. The coexistence ofe-mail: arakawas@jamstec.go.jp
26
inclusions indicates that a direct
injectionofObservatory
Al-richmaterials
a nearby
supernova would be expected to occur
National
Astoronomicalof Japan,
2-21-1 Osawa, from
Mitaka, Tokyo
181-8588, core-collapseJapan5in the first 10 years of the existenceof 2022the / Solar
System.
Therefore, at least one core-collapse supernova ought to occur within the
Received 12 August
Accepted 28
December 2022
duration of star formation in the Sun’s birth cluster. Here, we revisit the number of stars in the Sun’s birth cluster from the point of
ABSTRACT
view of the probability of experiencing at least one core-collapse
supernova within the finite duration of star formation within the birth
The Sunof
is thought
havethebeen birth
formed withina star cluster.
of Al-rich andAl-poorthan
calcium–aluminum-richcluster. We find that the number
starstoincluster
mayThebecoexistence
significantly
greaterthat previously considered, depending
inclusions indicates that a direct injection of Al-rich materials from a nearby core-collapse supernova would be expected to occur
in the first 10 years of the existence of the Solar System. Therefore, at least one core-collapse supernova ought to occur within theon the duration of star formation.
2

26

26

26

5

duration of star formation in the Sun’s birth cluster. Here, we revisit the number of stars in the Sun’s birth cluster from the point ofview of the probability of experiencing at least one core-collapse supernova within the finite duration of star formation within the birthcluster. We find that the number of stars in the birth cluster may be significantly greater than that previously considered, dependingon the duration of star formation.

Key words. meteorites, meteors, meteoroids – supernovae: general – stars: Wolf-Rayet – open clusters and associations: general

26Al
1. Introduction

Key words. meteorites, meteors, meteoroids – supernovae: general – stars: Wolf-Rayet – open clusters and associations: general

10^5

1. Introduction
The majority of stars form within star clusters (e.g.,
Lada & Lada 2003) and our Solar System is also thought tohave been formed within a cluster, based on a number of indications (e.g., Adams 2010). For example, some of trans-Neptunian
objects, including (90 377) Sedna, exhibit a substantial heliocentric eccentricity and perihelion distance (e.g., Brown et al.
2004) and the most plausible scenario for their formation isa close encounter with a passing star in the Sun’s birth cluster (e.g., Morbidelli & Levison 2004; Brasser et al. 2006). Parts
of presolar grains, whose isotopic compositions are significantly anomalous by comparison with the average solar composition, originate from core-collapse supernovae (CCSNe; e.g.,
Nittler & Ciesla 2016), and the chromium isotopic heterogeneity found in various chondritic meteorites might be evidence of
the injection of dust grains from a CCSN (e.g., Sugiura & Fujiya
2014; Fukai & Arakawa 2021).
An injection of dust grains into the Solar System froma nearby CCSN is also a classical scenario for the originof short-lived radionuclides (SLRs) found in meteorites (e.g.,
Takigawa et al. 2008; Ouellette et al. 2010). Such SLRs areradioactive nuclides that were present in the early Solar Sys-

The majority of stars form within star clusters (e.g.,
Lada & Lada 2003) and our Solar System is also thought tohave been formed within a cluster, based on a number of indications (e.g., Adams 2010). For example, some of trans-Neptunian
objects, including (90 377) Sedna, exhibit a substantial heliocentric eccentricity and perihelion distance (e.g., Brown et al.
2004) and the most plausible scenario for their formation isa close encounter with a passing star in the Sun’s birth cluster (e.g., Morbidelli & Levison 2004; Brasser et al. 2006). Parts
of presolar grains, whose isotopic compositions are significantly anomalous by comparison with the average solar com-

2002). Although the majority of primitive (i.e., unmelted) CAIs
are uniformly enriched in Al (e.g., MacPherson et al. 2012),
some unusual CAIs show the order-of-magnitude lower initialabundance of Al. Platy hibonite crystals (PLACs) are classified into those unusual CAIs and they are regarded as theoldest CAIs based on their high condensation temperature andlarge nucleosynthetic anomalies (e.g., Sahijpal & Goswami 1998;
Kööp et al. 2016), although their origins are still under debate(e.g., Larsen et al. 2020; Desch et al. 2022). In addition, CAIs
with fractionation and unidentified nuclear e↵ects (FUN CAIs)
also have low and varied initial abundances of 26 Al, and they arethought to have been formed prior to normal CAIs (e.g., Park et al.
2017). The coexistence of Al-rich and 26 Al-poor CAIs would

2002). Although the majority of primitive (i.e., unmelted) CAIs
are uniformly enriched in 26 Al (e.g., MacPherson26
et al. 2012),
some unusual CAIs show the order-of-magnitude lower initialabundance of 26 Al. Platy hibonite crystals (PLACs) are classified into those unusual CAIs and they are regarded as theoldest CAIs based on their 26
high condensation temperature andlarge nucleosynthetic anomalies (e.g., Sahijpal & Goswami 1998;
Kööp et al. 2016), although their origins are still under debate(e.g., Larsen et al. 2020; Desch et al. 2022). In addition, CAIs
with fractionation and unidentified nuclear e↵ects (FUN CAIs)
also have low and varied initial abundances of 26 Al, and they arethought to have been formed prior to normal CAIs (e.g., Park et al.
2017). The coexistence of 26 Al-rich and 26 Al-poor CAIs would
serve as evidence of a direct injection of 26 Al-enriched dust
grains into the solar nebula (or the solar molecular cloud core)in the epoch of CAI formation (e.g., Sahijpal & Goswami 1998;
Holst et al. 2013). As the duration of CAI formation (i.e., theformation age spread of CAIs) is a few 105 years or less (e.g.,Thrane et al. 2006; Connelly et al. 2012; Kawasaki et al. 2020),
the direct injection event would have occurred in the first 105 yearsof Solar System formation. We note that the initial abundanceof 26 Al recorded in CAIs would not be altered26
even if additional injection of 26 Al-enriched materials into the solar nebula

ABSTRACT

Ryan D. Boyden1 and Josh Eisner1
Steward
University
Arizona,
N. Cherry Ave,
AZ, 85719,
USA
The stellar cluster environmentisObservatory,
expected
to ofplay
a933
centralroleTucson,in theevolutionof circumstellar disks. We
August 23, 2022;
2, 2022;
Accepted December
21, 2022)UV and X-ray radiation fields,
use thermochemical modeling (Received
to constrain
the Revised
dustDecember
and gasmasses,disk sizes,
viewing geometries, and central stellar masses of 20ABSTRACTClass II disks in the Orion Nebula Cluster (ONC). We fit+ disks. We
stellar cluster
environment
expected to play aCO
centralcircumstellara large grid of disk The
modelsto 350GHz iscontinuum,
J role= in3the evolution
2, andof HCO
J = 4 3 ALMA observations
use thermochemical modeling to constrain the dust and gas masses, disk sizes, UV and X-ray radiation fields,of each target, andviewing
we introduce
procedure
forofmodelinginterferometricobservations
geometries, andacentral
stellar masses
20 Class II disks
in the Orion Nebula Cluster(ONC). We fit of gas disks detected in+grid of molecular
disk models to 350
GHz continuum,
CO J = 3We2, andHCO
J =the4 3ONCALMA observations
absorption againstaalargebrightcloud
background.
find
that
disks are massive and compact,of each target, and we introduce a procedure for modeling interferometric observations of gas disks detected in3with typical radii <100 AU,masses10background.
M ,Weandgas-to-dust
100.
The ISM-like gas-to-dust
absorption
againstgas
a bright
molecular cloud
find that
the ONC disks ratios
are massive andcompact,with typical radii < 100 AU, gas masses 10 3 M , and gas-to-dust ratios 100. The ISM-like gas-to-dust
ratios derived fromratios
ourderived
modelingsuggest that compact, externally-irradiated disks in the ONC are less prone
from our modeling suggest that compact, externally-irradiated disks in the ONC are less prone
to gas-phase CO
depletion
the massive
and extended
gas disks
thatdisksare commonly
nearby lowto gas-phase CO depletion
thanthethanmassiveandextendedgasthat foundare incommonly
found in nearby lowmass star-forming regions. The presence of massive gas disks indicates that external photoevaporation may
mass star-forming haveregions.
The
presence
disksunclear
indicatesthatcluster
external
only recently
begun
operating inof
themassive
ONC, thoughgas
it remains
whether othermembers photoevaporation may
older andoperating
more evaporated
ones in our
sample. Finally,
we compare
our dynamically-derived
have only recentlyare
begun
inthanthetheONC,
thoughit remains
unclearwhether other cluster members
stellar masses with the stellar masses predicted from evolutionary models and find excellent agreement. Our
are older and more
evaporated
ones ofindynamical
our sample.
Finally,
we rangecomparedynamically-derived
study
has significantly thanincreasedthe
the number
mass measurementsin the mass
 0.5 M our,demonstrating that the ONC is an ideal region for obtaining large samples of dynamical mass measurements
stellar masses withtowards
the low-massstellarM-dwarfs.
masses predicted from evolutionary models and find excellent agreement. Our
study has significantly increased the number of dynamical mass measurements in the mass range  0.5 M ,Keywords: open clusters and associations: individual (Orion) — planetary systems — protoplanetarydemonstrating that the ONC disks
is anidealpre-main
regionfor obtaining large samples of dynamical mass measurements— stars:
sequence
towards low-mass M-dwarfs.1. INTRODUCTION
A key property of young stellar clusters is the presence1

arXiv:2212.12325v1 [astro-ph.EP] 22 Dec 2022

12.12325v1 [astro-ph.EP] 22 Dec 2022

(Received August 23, 2022; Revised December 2, 2022; Accepted December 21, 2022)

Chemical modeling of Orion Nebula Cluster disks: evidence for massive, compact gas
disks with ISM-like gas-to-dust ratios

Protoplanetary disks form in a variety of environments, and the natal environment has the potential toshape the evolution of disks and planetary systems (fora recent review, see Parker 2020). The most common
environments of star formation are young stellar clusters(e.g., Lada et al. 1993; Lada & Lada 2003; Krumholz
et al. 2019). Understanding how disks evolve in clustered star-formation environments is therefore essential
for interpreting the demographics of disks and exoplanets (e.g., Ndugu et al. 2018; Winter et al. 2020b),
and examining how our own solar system fits into theof massive O- and B-type stars. Strong UV fields frommassive stars launch photoevaporative winds from thesurfaces of disks surrounding lower-mass cluster members through a process known as external photevaporation (e.g., Johnstone et al.
1998; Störzer & Hollenbach 1999). In regions of high stellar density andradiation field strength, external photoevaporation acts
as the dominant mechanism of disk dispersal (Scally &Clarke 2001; Winter et al. 2018; Concha-Ramírez etal. 2019), and can eﬀectively deplete the disk of planetforming material within < 1 Myr of cluster evolution

ClassandII associations: individual (Orion)
ALMA,
CO 3-2, systems
HCO+ 4-3
Keywords: open clusters— planetary— protoplanetary— stars:
pre-main sequencer<100 au, Mgas>10^disks
3 Msun,g/d>100
CO deplete
1. INTRODUCTION

Protoplanetary disks form in a variety of environments, and the natal environment has the potential to

A key property of young stellar clusters is the presenceof massive O- and B-type stars. Strong UV fields frommassive stars launch photoevaporative winds from the3

Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA
Received 2022 Julyrevised2022 November
accepted
November
29;Starpublished
2023 February 14
The 14;Star
Formation
History 29;of theMilky2022
Way’s
Nuclear
Cluster
Zhuo Chen1 , Tuan Do1 , Andrea M. Ghez1 , Matthew W. Hosek, Jr.1,4 , Anja Feldmeier-Krause2 , Devin S. Chu1 ,Abstract
Rory O. Bentley1 , Jessica
R. Lu3 , and Mark R. Morris1
1

Department of Physics and Astronomy, University of California, Los Angeles, CA 90095-1547, USA; zhuochen@astro.ucla.edu

Max Planck Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany
We report the ﬁrst star formation history
study of the Milky Ways nuclear star cluster (NSC), which includes
Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA
2022 Julysample
14; revised 2022
accepted 2022 November
29; published 2023 February
14
observational constraints fromReceived
a largeofNovember
stellar29; metallicity
measurements.
These
metallicity measurements
were obtained from recent surveys from Gemini and Abstract
the Very Large Telescope of 770 late-type stars within the
We report the ﬁrst
star formation historyalong
study ofwith
the Milky
Ways nuclear star(NSC), which includes derived temperatures,
central 1.5 pc. These metallicity
measurements,
photometry
andcluster
spectroscopicallyobservational constraints from a large sample of stellar metallicity measurements. These metallicity measurements
are forward modeled witha Bayesian
metallicity
measurements
were obtained
from recentinference
surveys from approach.
Gemini and the Including
Very Large Telescopeof 770 late-typestars within the improves the overall
central 1.5 pc. Thesered
metallicity
measurements,
along previouslywith photometry andspectroscopicallyderived temperatures,ﬁt quality, as the low-temperature
giantsthat were
difﬁcult
to constrain
are now accounted for, andare forward modeled with a Bayesian inference approach. Including metallicity measurements improves the overall
the best ﬁt favors a two-component
model. The
dominant
component± 3%for,ofandthe mass, is metal-rich
ﬁt quality, as the low-temperature
red giants
that were previously
difﬁcult tocontains
constrain are 93%
now accounted
3 model. The dominant component contains 93% ± 3% of the mass, is metal-richthe best ﬁt favors a two-component
hasH] anage of 52ageGyr,which is ∼3 Gyr younger than earlier studies with ﬁxed (solar)([M H] _ 0.45), and ([M_ 0.45), and has anof 523 Gyr, which is ∼3 Gyr younger than earlier studies with ﬁxed (solar)
metallicity;
younger age challenges
coevolutionary models
models in which
the NSC andblacksupermassive
holes
metallicity; this younger
age thischallenges
coevolutionary
in whichthesupermassive
NSC andblack holes
formed simultaneously at early times. The minor population component has low metallicity ([M H] _ 1.1) andformed simultaneouslycontains
at early
times.
The
minor
componenthas(0.1–5
lowGyrmetallicity∼7% ofthe stellar
mass.
The agepopulationof the minor component
is uncertainold). Using the ([M H] _  1.1) andestimatedparameters,
infer thefollowing
NSC stellar
remnant
population (withis∼18%
uncertainty): (0.1–5
1.5 × 105 Gyr old). Using thecontains ∼7% of theneutron
stellarmass.weTheageof theminor
component
uncertainstars, 2.5 × 105 stellar-mass black holes (BHs), and 2.2 × 104 BH–BH binaries. These predictions result in
estimated parameters, 2–4
wetimesinferfollowing
NSC
stellarremnant
population∼18%
uncertainty): 1.5 × 105
fewertheneutronstars comparedto earlier
predictions
that assume
solar metallicity,(withintroducing
a possible
5 path to understand the so-called “missing-pulsar problem”. Finally, we
present updated predictions for theneutron stars, 2.5 × 10new
stellar-mass black holes (BHs), and 2.2 × 104 BH–BH
binaries. These predictions result in
BH–BH merger rates (0.01–3 Gpc−3yr−1).
2–4 times fewer neutronstars
comparedearlier
predictions
thatassumesolar
UniﬁedAstronomy
Thesaurus toconcepts:
Galactic
center (565); Star
formation(1569);
Galaxymetallicity,formation (595); introducing a possible
Milky
Way
formation(1053);
Infrared
spectroscopy(2285);
Adaptive
optics(2281);
Adaptiveinterferometry(19)
new path to understand the so-called “missing-pulsar problem”. Finally, we present updated
predictions for the−1
Supporting material: −3
machine-readable
tables
BH–BH merger rates (0.01–3 Gpc yr ).
Uniﬁed Astronomy Thesaurus 1.concepts:
StarPrevious
formation(1569);
Galaxy(595);
Introduction Galactic center (565);
understood.
studies havesuggestedthat the formation
star
formation
history
of
the
NSC
is
complex.
The
stellarMilky Way formation(1053);
Infrared
spectroscopy(2285);
Adaptive optics (2281); Adaptive interferometry (19)The innermost
regionof most galaxies
is occupied byapopulation of the NSC is composed of cool, evolved giants;
spectacularly
dense
and
massiveassemblyofstars,
which
Gemini,Supporting
VLT, 770material:,machine-readable
1.5tablespc,
and hot, young main-sequence/post-main-sequence stars.
forms a nuclear
star cluster (NSC). The
star formation in thisBlum et al. (2003) combined spectroscopic and photometric23

region is believed to be affected by the central supermassive
black hole (SMBH), but the physical mechanisms behind it are
not entirely known. The center of the Milky Way galaxy is hostthe closest example of a SMBH (4.2 × 106Me; e.g., Ghez
1.ettoIntroduction
al. 2008; Gillessen et al. 2009; Do et al. 2019), embedded in
an NSC (∼2.5 × 107Me; e.g., Launhardt et al. 2002; Schödel
The innermost regionet ofmostgalaxiesaal. 2014).
Given
its proximity,isthe occupied
Milky Way NSCby
providesuniqueopportunity
to resolve thepopulationand tospectacularly dense andastudy
massiveassembly
ofstellarstars,
which
phenomena and physical processes that may beforms a nuclear star cluster(NSC).
formation in thishappeningin otherThegalacticstar
nuclei.

93%
5 asymptotic
Gyr giant branchobservations of the 79 most luminous(AGB) and supergiant stars in the central 5 pc. They
missing-pulsar
constructed
the Hertzsprung–Russell (H-R) diagram from COfeaturesin H- andhave
K-band suggested that the starandH2O molecular absorption
understood.
Previous
studies
spectra, and claimed that ∼75% of stars formed more thanhistoryof theis complex. The stellar
5 formation
Gyr ago. Maness etal. (2007) reported
the ﬁrstNSC
study using
adaptive
optics(AO)
observationsof329
giantsinthecentralpopulation of the NSC is composed1 of cool, evolved giants;
pc, including helium-burning red clump stars, red giants, and
AGB
These stars
with longer-lived
evolutionary phases
andstars.hot,
youngmain-sequence/post-main-sequencestars.
are better understood by theoretical evolutionary models. They

Received 2022 September 22; revised 2022 November 30; accepted 2022 December 5; published 2023 February 16
Making the Solar System

Abstract
We model the early stages of planet formation in the solar system, including continual planetesimal formation, andAbstract
planetesimal and pebble accretion onto planetary embryosin an evolving disk driven by a disk wind. The aim is to
We model the early stages of planet formation in the solar system, including continual planetesimal formation, andconstrain aspects of planet
formation
large
uncertainties
key
characteristicsof the solar
planetesimaland pebble that
accretionhave
onto planetary
embryosin an evolving diskby
drivenmatchingby a disk wind. The
aim istoconstrainaaspectsof planet
formationthat characteristics
have large uncertainties by
key characteristics
the solarsystem. The model produces
good
ﬁt to
these
formatchinga narrow
rangeof ofparameter space. Planetarysystem. The model produces a good ﬁt to these characteristics for a narrow range of parameter space. Planetary
growththe ice line isby
dominatedby pebble
accretion. Planetesimalaccretion is more
important inside
growth beyond the ice line
is beyonddominatedpebbleaccretion.
Planetesimalaccretionis themore important inside theice line. Pebble accretion inside the ice line is slowed by higher temperatures, partial removal of inﬂowing pebbles
ice line. Pebble accretion byinsidetheformation
ice lineis slowed
by highertemperatures,
partialdueremovalof inﬂowing pebbles
planetesimaland pebble
accretion further
out in the disk,
and increased radial velocities
to gasadvection. The terrestrial planets are prevented from accreting much water ice because embryos beyond the ice line
by planetesimal formationreachandpebble accretionoutin the disk,
andradial velocities due to gas
the pebble-isolation
mass before thefurther
ice line enters
the terrestrial-planet
region.
Whenincreased
only pebble accretionis considered, embryos typically remain near their initial mass or grow to the pebble-isolation mass. Adding
advection. The terrestrial planetesimalplanets accretionare prevented
from accreting much water ice because embryos beyond the ice line
allows Mars-sized objects to form inside the ice line, and allows giant-planet cores to forma widerbefore
region beyond
ice line.
In theentersregion occupied
Mercury, pebble Stokes numbers
are small.
This
reach the pebble-isolationover
mass
thethe iceline
the byterrestrial-planet
region.
When
only pebble accretion
delays the formation of embryos and stunts their growth, so that only low-mass planets can form here.
is considered, embryos typicallyremaininitial(1241);
mass
grow
to system
the pebble-isolation
mass. Adding
Uniﬁed Astronomy
Thesaurus nearconcepts:their
Planet formation
Solar orsystem(1528); Solargas giantplanets (1191); Solar system terrestrial planets (797); Solar nebulae (1508); Mercury (planet) (1024);
planetesimal accretion allows
Mars-sized objects to form inside the ice line, and allows giant-planet cores to form
Planetesimals (1259)
over a wider region beyond the ice line. In the region occupied by Mercury, pebble Stokes numbers are small. This
Introduction
causesa collapse
form asteroid-sized
bodiescancalledform here.
delays the formation of embryos 1.andstunts their growth, sothat
only tolow-massplanets
planetesimals (Johansen et al. 2007, 2012).
The recent discovery of several thousand extrasolar
Further growth can take place in two ways. The largest
planetary systems concepts:
has spurred greatPlanet
interest information
how these
Uniﬁed Astronomy Thesaurus(1241);
Solar system (1528); Solar system gas giant
planetesimals, also known as planetary embryos, can grow bysystems formed. While this emphasis on exoplanets is undersweeping
up planetesimals
or pebbles.
The accretionof pebbles (1024);
planets (1191); Solarstandable,
systemterrestrial
nebulae(1508);
Mercury(planet)it is worth
bearing in planets
mind that we(797);
still lack Solarais made more effective by drag with the disk gas, so thatmodel for how our own planetary system formed. This
pebbles lose energy during an encounter with an embryo
Planetesimals (1259)detailed
unresolved problem is made more compelling by the fact thatJohn Chambers

Earth and Planets Laboratory, Carnegie Institution for Science, 5241 Broad Branch Road NW, Washington, DC 20005, USA; jchambers@carnegiescience.edu
Received 2022 September 22; revised 2022 November 30; accepted 2022 December 5; published 2023 February 16

the solar system differs from most known exoplanet systems.
The absence of super-Earth-mass planets on short-period orbits
in the solar system is particularly noteworthy given that suchplanets are common elsewhere.
There are good reasons to focus on the origin of ourplanetary system. For example, it is the only system known tohost a habitable planet. We also know much more about thesolar system than other systems. We have a complete census ofplanetary-mass bodies out to at least 50 au, as well as much
more detailed information about the orbits of minor planet
populations in our system than elsewhere. In addition, the solar
system is the only system for which we have samples that canbe attributed to particular source objects. All this information
provides a strong test for planet formation theories.

1. Introduction

The recent discovery of several thousand extrasolar
planetary systems has spurred great interest in how thesesystems formed. While this emphasis on exoplanets is understandable, it is worth bearing in mind that we still lack a(Ormel & Klahr 2010). Planetesimal accretion is typicallyslower, especially far from the star (Morbidelli et al. 2015). For
this reason, pebble accretion is thought to be the main driver ofgrowth in the outer parts of disks (Lambrechts & Johansen2012), while both mechanisms may be important in the
inner disk.
Pebbles are expected to drift radially inwards over time due
to gas drag (Weidenschilling 1977). This drift provides acontinuous supply of material for embryos to grow by pebble
accretion, but it also removes pebbles that are not accreted
quickly enough. Thus, pebble accretion can form massive
planets by accessing material from a large fraction of the disk,
but the process is relatively inefﬁcient so a massive disk isrequired. When an embryo grows to a certain mass, it will
perturb the local gas pressure proﬁle sufﬁciently to halt thecauses a collapse to form asteroid-sized bodies called
planetesimals (Johansen et al. 2007, 2012).
Further growth can take place in two ways. The largest
planetesimals, also known as planetary embryos, can grow bysweeping up planetesimals or pebbles. The accretion of pebbles
is made more effective by drag with the disk gas, so that

Accepted 2022 November 30. Received 2022 November 29; in original form 2022 May 19

Length-scales and dynamics of Carina’s Western wall
1 Centre for Astrophysics & Relativity, School of Mathematical Sciences, Dublin City University, Glasnevin, Dublin 9, Irelandof Physics and Astronomy,
Rice University,
6100 Main
Street, MS 108,of
Houston,TX 77005, USA
We present a variety of analyses of Department
the turbulent
dynamicsof theboundarya photodissociation
region (PDR) in the Carina
Nebula using high resolution ALMA observations. Using principal component analysis, we suggest that the turbulence in thisAccepted 2022 November 30. Received 2022 November 29; in original form 2022 May 19
molecular cloud is driven at large scales. Analysis of the centroid velocity structure functions indicate that the turbulence isdominated by shocks rather than local(in k-space) transport of energy. We further find that length-scales in the range 0.02–
ABSTRACT
We present
a varietyof analysesof the turbulentof the boundarya photodissociation(PDR) in theemission
Carinaof thiscloud
and thisfindingdynamics
is supported
byofanalysisof theregion
dominant
structure
0.03 pc are important in the dynamics
Nebula using high resolution ALMA observations. Using principal component analysis, we suggest that the turbulence in thislength-scale. These length-scales arewellcloudresolved
thescales.
observational
data velocity
and weconcludethat theapparentofmolecular
is driven byat large
Analysis of the centroid
structurefunctions indicatethat theturbulence importanceisdominatedby shocks
localwell(in k-space)transport
energy. strongly
We further findthat length-scales
the range 0.02–diffusion, we
this range of scales is physical in origin.
Given
thatrather
it isthanalso
withintheofrange
influenced
byinambipolar
0.03 pc are important in the dynamics of this cloud and this finding is supported by analysis of the dominant emission structure
conclude that it is not primarily a productturbulencealone,
but isby more
likelydata
to andbewea conclude
result thatof thetheinterplay
betweengravitylength-scale.of
These
length-scales arewell resolved
the observational
apparent
importanceofthis range of scales is physical in origin. Given that it is also well within the range strongly influenced by ambipolar diffusion, weand turbulence. Finally, through comparisonof these results with previous observations of H2 ofemission
from the Western Wall,
conclude that it is not primarily a product of turbulence alone, but is more likely to be a result
the interplay between gravity
Finally,
with previous
observationsof Hundisturbed
Western Wall,we demonstrate that observations ofanda turbulence.
PDR canbethroughusedcomparison
to probeof these
the resultsinternalstructureof theof a molecular
2 emission from theportion
we demonstrate that observations of a PDR can be used to probe the internal structure of the undisturbed portion of a molecular
cloud.
cloud.
2

PDR

ALMA

PCA

Downloaded from https://academic.oup.com/mnras/article/519/4/5427/6979812 by KASI user on 23 February 202

words: turbulence – methods: statistical – ISM: kinematics and dynamics – ISM: magnetic fields.
Key words: turbulence – methods:Key
statistical– ISM: kinematics and dynamics – ISM: magnetic fields.

0.02-0.03

effects of density (emission and absorption), temperature, line-ofsight projection, and velocity. Projection can be expected to averageout velocity and density variations along the line of sight throughTurbulence in molecular clouds is known to be likely to have an imeffects of density(emissionand Thisabsorption),
temperature,line-ofpact on star formation (e.g. Krumholz & McKee 2005; Kainulainen,
so-called ‘projectionsmoothing’.
happens preferentiallyfor
1 I N T RO D U C T I O NFederrath & Henning 2013; Burkhart et al. 2015; Kainulainen&
variations
that
occur
on
length-scales
shortin
comparison
tothesight projection,and velocity. Projection can be expected to average
Federrath 2017), and even planet formation. Much of this turbulence
thickness of the cloud being observed. As a result, this projection
velocityand density
variations
alongof how
thethelineof sight throughTurbulence in molecular clouds is known
to beto likely
anflows,
im-with the outis thoughtbe drivento
by havelarge-scalepotentialforsmoothing
can yield misleadinginterpretations
turbulentsmaller scale turbulence also excited by stellar feedback. The energy
energy cascade proceeds. In addition, if large-scale magnetic fields
pact on star formation (e.g. Krumholz &budget
McKee
2005; Kainulainen,
so-called ‘projectionsmoothing’. This happens preferentially forof this turbulence must be large as it is both supersonic andare dynamically important, the turbulence can be anisotropic (e.g.
(Mac; Low
et al. 1998; Stone,& GammieCho, Lazarian
2002; Beattie, Federrath& Seta), asFederrath & Henning 2013; Burkhart fast-decaying
et al. 2015
Kainulainen& Ostrikervariations
thatoccur& Vishniac
on length-scales
shortin2020comparison
to the1998) and understanding it should give us insight into the origin ofwould be expected from the theory of Goldreich & Sridhar (1995):
Federrath 2017), and even planet formation.
Muchof thisthicknessthebeing
observed.
a result,
the initial
mass function(e.g.turbulenceNam, Federrath & Krumholz
2021), ofthenthe cloudangle between
the line-of-sightand Asthe large-scale
fieldthis projection
molecular
cloud
formation,andgalactic
dynamics.
The
natureofthebecomesacriticalcomponentindetermining
what
we
actuallyis thought to be driven by large-scale turbulent
flows,cascadewith isthe
potential forsmoothing canyield misleading interpretations of how the turbulentof particular interest, both from a theoreticalobserve. Indeed, while on the largest turbulence scales this may
smaller scale turbulence also excited byperspective
stellar feedback.
energyenergyofcascade
proceeds.
addition,
if large-scale
magnetic fields
and also a moreThepragmatic
one, since the properties
not be anissue, it will alIn
ways
be an issue below
the length-scale atthis cascade are what create the initial conditions from which stars
which the turbulent velocity is of order the Alfvén speed.
budget of this turbulence must be largeandasplanets
it isultimately
both supersonicandare dynamicallyimportant, the turbulence can be anisotropic (e.g.
form.
Notwithstanding these difficulties, in recent years, significant
1 I N T RO D U C T I O N

pc

ambipolar diffusion

Downloaded from https://academic.oup.com/mnras/article/519/4/5427/6979812 by KAS

Turlough P. Downes ,1 ‹ Patrick Hartigan2 and Andrea Isella2

ABSTRACT© 2023. The Author(s). Published by the American Astronomical Society.

78

Zhejiang Lab, Hangzhou, Zhejiang 311121, Peopleʼs Republic of China
Shanghai Astronomical Observatory, Chinese Academy of Sciences, 80 Nandan Road, Shanghai 200030, Peopleʼs Republic of China
New Molecular
Bubble–Outﬂow
Structurein the
B18
Cloud
ReceivedDiscovery
2022 Juneof
28;a revised
2022 November
30; accepted 2022
November
30;Taurus
published2023
February 7
Yan Duan1,2 , Di Li1,3,4 , Paul F. Goldsmith5 , Laurent Pagani6 , Tao-Chung Ching1,7 , Shu Liu1 , Jinjin Xie8 , andChen Wang1

Abstract
Star formation can produce bubbles and outﬂows, as a result of stellar feedback. Outﬂows and bubbles inject
momentum and energy into the surrounding interstellar medium, and so are related to the overall energy balance ofthe molecular cloud. Molecular bubbles can be resolved by higher-resolution radio telescopes to quantify the effectof star formation on molecular clouds. We report here the identiﬁcation of a new molecular bubble with an outﬂow,
and a Herbig–Haro object, HH 319, located at the bubble
center. Multiwavelength data have been utilized to study
Abstract
its spatial structure, energy
injection,
andbubbles
dynamical
timescale.
bubbleOutﬂows
has aandkinetic
energy of 5.8 × 1043 erg
Star formation
can produce
and outﬂows,
as a result of This
stellar feedback.
bubbles inject
momentum and energy into the surrounding interstellar medium, and so are related to the overall energy balance ofwithin the smallest radiusof a cloud.
bubblein Taurus,
0.077
The bubble
formedyr ago. According to thethe molecular
Molecular
bubbles can be
resolvedpc.
by higher-resolution
radio
telescopes ∼70,000
to quantify the effectof starof
formation
on molecular
clouds.Gaia
We reporthere the identiﬁcationof a new binary
molecular bubble
an outﬂow,
proper-motion velocities
protostars
from
EDR3,
the T Tauristarswith(FY
Tau and FZ Tau) at theand a Herbig–Haro object, HH 319, located at the bubble center. Multiwavelength data have been utilized to study
43
southwest edge of theitsbubble
mayenergy
haveinjection,
producedthe outﬂow–bubble
is10anergunusual new structure
spatial structure,and dynamicaltimescale. This bubble hasstructure.
a kinetic energyThisof 5.8 ×withinthesmallest
radiusofabubblein
Taurus,
0.077
pc.
The
bubbleformed∼70,000
yr
ago.
Accordingtothe by V380 Ori, has afound in low- and intermediate-massstar formation regions. Only a bubble in Orion A, driven
proper-motion velocities of protostars from Gaia EDR3, the T Tauri binary stars (FY Tau and FZ Tau) at thesimilar structure. The southwest
bubble–outﬂow
structureprovidesadditionalstructure.
observational
evidence
for the theory of stellaredge of the bubble
may have produced
the outﬂow–bubble
This is an unusual
new structure
found in low- and intermediate-mass star formation regions. Only a bubble in Orion A, driven by V380 Ori, has awind from T Tauri stars.
enhances
our understanding
howobservational
stellar feedback
molecular clouds.
similarItstructure.
The bubble–outﬂow
structure provides ofadditional
evidence for theacts
theoryonof stellar
1National Astronomical Observatories, Chinese Academy of Sciences, Beijing 100101, Peopleʼs Republic of China; dili@nao.cas.cn
2
University of Chinese Academy of Sciences, Beijing 100049, Peopleʼs Republic of China
3
Research Center for Intelligent Computing Platforms, Zhejiang Laboratory, Hangzhou 311100, Peopleʼs Republic of China
4
NAOC–UKZN Computational Astrophysics Centre, University of KwaZulu-Natal, Durban 4000, South Africa
5
Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109, USA
6
LERMA & UMR 8112 du CNRS, Observatoire de Paris, PSL University, Sorbonne Universités, CNRS, F-75014 Paris, France
7
Zhejiang Lab, Hangzhou, Zhejiang 311121, Peopleʼs Republic of China
8
Shanghai Astronomical Observatory, Chinese Academy of Sciences, 80 Nandan Road, Shanghai 200030, Peopleʼs Republic of China
Received 2022 June 28; revised 2022 November 30; accepted 2022 November 30; published 2023 February 7

wind from T Tauri stars. It enhances our understanding of how stellar feedback acts on molecular clouds.

Uniﬁed Astronomy Thesaurus
concepts: Stellar wind bubbles (1635); Stellar winds (1636); Molecular clouds
Uniﬁed Astronomy Thesaurus concepts: Stellar wind bubbles (1635); Stellar winds (1636); Molecular clouds(1072); T Tauri stars (1681);
Stellarfeedback(1602)(1072); T Tauristars (1681);
Stellar feedback(1602)

HH319
bubbleshave a greater impact on molecular clouds thanFY Tau, FZ
Tau

5.8e48 erg, 0.077 pc, 70,0001. Introduction
Star formation is related to the structure and evolution ofIntroduction
T1.Taurigalaxies and giant molecular clouds, the synthesis of chemicalelements, and the star formation rate and is intimately
Star formation is related
to theand evolutionofconnected
with astructure
critical astrophysicalparameter, the initial
mass function(McKeeOstriker 2007).
all younggalaxies and giant molecular
clouds,the &synthesisofAlmost
chemicalstellar objects (YSOs) have undergone periods of substantial
elements, and the star mass
formation
rateand andis stellar
intimately
loss. YSOs’ fast
collimatedwind sweeps
throughthesurrounding
moleculargas.
The
molecular
gas isconnected with a criticalexpelled
astrophysicalparameter, the initialfrom the cavities and expands as irregular lobes andincomplete
shells (Bachiller
These become
visible as
mass function (McKee &Ostriker2007). 1996).
Almostall youngmolecular bubbles and molecular outﬂows. Herbig–Haro (HH)
stellar objects (YSOs) haveperiods
of gas,substantial
objects,undergone
consisting of knotsof ionized
are the opticalmanifestationsof
thismass-loss
process(Reipurth &mass loss. YSOs’ fast and collimated stellar wind sweeps
Bally 2001; McKee & Ostriker 2007).

outﬂows. The dynamical timescales of bubbles are longer thanthose of outﬂows in Taurus (Li et al. 2015). Thus, the energy
injection from bubbles into the molecular clouds could last
longer. The ensemble of bubbles in Taurus have about 110 times
greater mass and 24 times higher energy than the total of theoutﬂows (Li et al. 2015). Compared to supernova remnants,
molecular bubbles could exist for a longer period of time (Arceet al. 2011). Due to their complex morphologies, bubbles are
more difﬁcult to identify than relatively clear outﬂow features(Arce et al. 2011). Therefore, the study of molecular bubbles isboth important and challenging.
Bubbles arise from a largely spherical stellar wind.
Observationally, bubbles are similar to complete or partial
rings. Arce et al. (2011), Li et al. (2015), and Feddersen et al.

bubbles have a greater impact on molecular clouds thanoutﬂows. The dynamical timescales of bubbles are longer thanthose of outﬂows in Taurus (Li et al. 2015). Thus, the energy
injection from bubbles into the molecular clouds could last
longer. The ensemble of bubbles in Taurus have about 110 times
greater mass and 24 times higher energy than the total of theoutﬂows (Li et al. 2015). Compared to supernova remnants,
molecular bubbles could exist for a longer period of time (Arceet al. 2011). Due to their complex morphologies, bubbles are

3Exzellenzcluster ORIGINS, Boltzmannstr. 2, D-85748 Garching, Germany e-mail: franceschi@mpia.de

Constraining the turbulence and the dust disk in IM Lup: onset ofplanetesimal formation
2, 3 , Thomas Henning1 , Anirudh Sharma2
Riccardo Franceschi1 , Tilman Birnstiel
ABSTRACT1

Max-Planck-Institut für Astronomie (MPIA), Königstuhl 17, 69117 Heidelberg, Germany

Context. Observations of 2protoplanetary
information
on planet
formationthereasons for the diversity of planetary
University Observatory,disks
Faculty provideof Physics, Ludwig
Maximilians University,
Scheinerstr.
1, 81679 andMunich,Germany

3
Exzellenzclusterplanet
ORIGINS,formation
Boltzmannstr. 2, isD-85748
e-mail:franceschi@mpia.de
systems. The key to understanding
the Garching,study Germanyof dustevolutionfrom small grains to pebbles. Smaller grains
December
5, 2022(⇠ 1 µm) are well-coupled
to thegas dynamics, and their distribution is significantly extended above the disk midplane. Larger
grains settle much faster and are efficiently formed only in the midplane. By combining near-infrared polarized light and millimeter
ABSTRACT
observations, it is possible to constrain the spatial distribution of both the small and large grains.
Context. Observations of protoplanetary disks provide information on planet formation and the reasons for the diversity of planetary
Aims. We aim to constructdetailed models of the size distribution and vertical/radial structure of the dust particles in protoplanetarysystems. The key to understanding planet formation is the study of dust evolution from small grains to pebbles. Smaller grains(⇠ 1 µm)
are well-coupled
to the gaswe
dynamics,
and their distribution
is significantly extended
abovedistribution
the disk midplane.in
Largerdisks based on observationaldata.
In particular,
are interested
in recovering
the dust
the IM Lup protoplanetary
grains settle much faster and are efficiently formed only in the midplane. By combining near-infrared polarized light and millimeter
disk.
observations, it is possible to constrain the spatial distribution of both the small and large grains.
Aims. We aim
to constructmodelsof the size distribution
and vertical/radial structureof theWedust particles
in protoplanetary
Methods. We create a physicalmodel
for detailed
the dust
distributionof protoplanetarydisks.
then simulate
the radiative transfer of thedisks based on observational data. In particular, we are interested in recovering the dust distribution in the IM Lup protoplanetary
millimeter continuum anddisk.
the near-infrared polarized radiation. Using a Markov chain Monte Carlo method, we compare the derived
Methods. We create a physical model for the dust distribution of protoplanetary disks. We then simulate the radiative transfer of theimages to the observations
available
forandthe
IM Lup polarized
disk toradiation.
constrain
best
modelwefor
IMtheLup
and to recover the verticalmillimeter
continuum
the near-infrared
Using a theMarkovchainphysicalMonte Carlo method,comparederivedto the observations available for the IM Lup disk to constrain the best physical model for IM Lup and to recover the verticalgrain size distribution. images
grain size distribution.
Results. The millimeter andnear-infrared
tightlytightly
constrainthemass
anddistribution
grain size
distributionof our model. We findResults.
The millimeter andemission
near-infrared emission
constrain the
dustdust
mass andgrain sizeof our model.
We findsize segregation in the dust distribution, with millimeter-sized grains in the disk midplane. These grains are efficiently formed insize segregation in the dust
distribution, with millimeter-sized grains in the disk midplane. These grains are efficiently formed in
the disk, possibly by sedimentation-driven coagulation, in accord with the short settling timescales predicted by our model. This alsosuggests a high dust-to-gas ratio
at smaller radii inin
the accord
midplane, possibly
streaming
instabilities
and planetesimalformation by our model. This also
the disk, possibly by sedimentation-driven
coagulation,with triggering
the shortsettling
timescales
predicted
in the inner disk. We obtain a turbulent ↵ parameter of 3 ⇥ 10 3 .
suggests a high dust-to-gas ratio at smaller radii in the midplane, possibly triggering streaming instabilities and planetesimal formation
Key words. protoplanetary disks – accretion disks – planetsand satellites: formation – circumstellar matter – stars: pre-mainin the inner disk. We obtaina turbulent↵ parameterof 3 ⇥ 10 3 .
sequence— radio continuum:
planetary systems

rXiv:2212.01291v1 [astro-ph.EP] 2 Dec 2022

01291v1 [astro-ph.EP] 2 Dec 2022

December 5, 2022

Key words. protoplanetary disks – accretion disks – planets and satellites: formation – circumstellar matter – stars: pre-main
1. Introduction
sequence — radio continuum:
planetary systems

1. Introduction

IM LupProtoplanetary disks around young stars are the nursery of planetary systems. These disks are almost entirely composed of gas
and a small mass fraction of dust. Small dust grains (⇠ 1 µm) are
well-coupled to the gas structure, and therefore extend verticallyover several scale heights, following the gas distribution. Dust
settling in turbulent disks was extensively studied (e.g. Dubrulle
et al. 1995; Schräpler & Henning 2004; Fromang & Nelson
2009; Zsom et al. 2011; Woitke et al. 2016). On the contrary,large millimeter grains settle much faster towards the disk midplane, where the growth process is particularly efficient becauseof the high densities. By constraining the distribution of both

MCMC

g/d

Protoplanetary disks around young stars are the nursery of planetary systems. These disks are almost entirely composed of gas
and a small mass fraction of dust. Small dust grains (⇠ 1 µm) arecently observed a correlation between the CO emitting height
and the disk size in a sample of disks observed with the Atacama Large Millimeter/submillimeter Array (ALMA). Using
the dust emission, Avenhaus et al. (2018) measured the flaring of V4046 Sgr, RXJ 1615, and IM Lup with the Very Large
Telescope (VLT) Spectro-Polarimetric High-contrast Exoplanet
REsearch (SPHERE) instrument. These near-infrared polarized
light images show unprecedented details of the disk’s surface
layer by tracing the scattering of small, micron-sized dust grains.
At these wavelengths, the disk is expected to be optically thickand the observed light is dominated by light scattered by the dust
grains on the disk surface. Among these disks, IM Lup observations are arguably among the most impressive images of a proto-

sedimentation-driven

cently observed a correlation between the CO emitting height
and the disk size in a sample of disks observed with the Atacama Large Millimeter/submillimeter Array (ALMA). Using
the dust emission, Avenhaus et al. (2018) measured the flaring of V4046 Sgr, RXJ 1615, and IM Lup with the Very Large

9 Carnegie Observatories,
A 3D
View
of Orion:
I. Barnard’s
Loop CA 91101, USA
813
Santa
BarbaraSt., Pasadena,
10 Max Planck Institutefor Astrophysics,
Karl-Schwarzschildstraße
1,85748
Germany
Michael M. Foley ,1, ⇤ Alyssa
Goodman ,1 Catherine
Zucker ,2, † John C.
Forbes Garching,,3 Ralf Konietzka
,1, 4
5
5
6
7
5
Cameren Swiggum , João Alves , John Bally , Juan D. Soler , Josefa E. Großschedl ,Shmuel Bialy ,8 Michael Y. Grudić ,9, † Reimar Leike ,10, 4 and Torsten Enßlin 10, 4

ABSTRACT

arXiv:2212.01405v1 [astro-ph.GA] 2 Dec 2022

arXiv:2212.01405v1 [astro-ph.GA] 2

8 Department of Astronomy, University of Maryland, College Park, MD 20742, USA

1 Center for Astrophysics | Harvard & Smithsonian, 60 Garden St., Cambridge, MA, USA 02138
2 Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA

Barnard’s Loop is 3 Centera famousarc of H↵ emission located in the Orion star-forming region. Here,for Computational Astrophysics, Flatiron Institute, 162 5th Avenue, New York, NY, 10010, USA
4 Ludwig-Maximilians-Universität, Geschwister-Scholl Platz 1, 80539 Munich, Germany
we provide evidence of a possible
formation mechanism for Barnard’s Loop and compare our results
5 Department of Astrophysics, University of Vienna, Türkenschanzstrasse 17, 1180 Wien, Austria
with recent work suggestinga majorandfeedback
event occurred in the region around 6 Myr ago. We
6 Department of AstrophysicalPlanetary Sciences, University of Colorado, Boulder, Colorado 80389, USA
di Astrofisica
e Planetologia
Spaziali(IAPS). INAF.
Via Fosso delcoherent,
Cavaliere 100, 00133
Roma,3D
Italy expansion of thepresent a 3D model7 Institutoof thelarge-scale
Orion
region,indicatingradial,8 Department of Astronomy, University of Maryland, College Park, MD 20742, USA
9 Carnegie
OBP-Near/Briceño-1 (OBP-B1)
clusterin theofSt.,a Pasadena,large dust
cavity.
The large-scale gas in the
Observatories,
813 middle
Santa Barbara
CA 91101,
USA
10 Max Planck Institute for Astrophysics, Karl-Schwarzschildstraße 1,85748 Garching, Germany
region also appears to be expanding from a central point, originally proposed to be Orion X. OBP-B1
appears to serve as another possible center, andwe evaluate whether Orion X or OBP-B1 is more
ABSTRACT
Barnard’s
Loop expansion.
is a famous arc We
of H↵findemission
locatedin the cluster
Orion star-formingHere,likely to be the causeof thethatneitherservedregion.
as thesingle expansion
we provide evidence of a possible formation mechanism for Barnard’s Loop and compare our results
center, but ratherwith
a recent
combinationof feedback from both likely propelled the expansion. Recent 3D
work suggesting a major feedback event occurred in the region around 6 Myr ago. We
present
a 3D model of thethe
large-scale
Orion region,
coherent,
radial,which
3D expansionof thedust maps are used
to characterize
3D topology
of indicating
the entire
region,shows
Barnard’s Loop’s
OBP-Near/Briceño-1 (OBP-B1) cluster in the middle of a large dust cavity. The large-scale gas in thecorrespondence witha large dust cavity around the OPB-B1 cluster. The molecular clouds Orion A,
region also appears to be expanding from a central point, originally proposed to be Orion X. OBP-B1
Orion B, and Orion
the shell
ofcenter,this andcavity.
Simpleestimates
gravitationale↵ects fromappears reside
to serve onas anotherpossible
we evaluate
whether
Orion X orofOBP-B1
is morelikely to be the cause of the expansion. We find that neither cluster served as the single expansion
both stars and gas
indicate that the expansion of this asymmetric cavity likely induced anisotropy incenter, but rather a combination of feedback from both likely propelled the expansion. Recent 3D
the kinematics of dust
OBP-B1.
Wetoconclude
fromthe structure of themaps are usedcharacterize thatthe 3D feedback
topology of theentireOBP-B1
region, whichhasshowsa↵ected
Barnard’s Loop’s
correspondencewith amolecularlarge dust cavity
aroundand
the OPB-B1
cluster. played
The molecular
clouds Orion
Orion A, Orion B,
and Orion
cloudsmay havea majorrole A,in the formation
Orion B, and Orion reside on the shell of this cavity. Simple estimates of gravitational e↵ects fromof Barnard’s Loop.
both stars and gas indicate that the expansion of this asymmetric cavity likely induced anisotropy inthe kinematics of OBP-B1. We conclude that feedback from OBP-B1 has a↵ected the structure of theOrion A, Orion B, and Orion molecular clouds and may have played a major role in the formation
Keywords: stellaroffeedback,dust shells, supernovae, star formation
Barnard’s Loop.

OBP-B1

OBP-B1
1. INTRODUCTION
1. INTRODUCTION
OBP-B11.1. Orion Complex and Barnard’s Loop

Keywords: stellar feedback, dust shells, supernovae, star formation

OBP-B1

Orion X
A, B, λOBP-B11.1. Orion Complex and Barnard’s Loop

The Orion complexThe
is Orionthe closestof current
massiveformation,roughly
400-430
complex is site
the closest
site of current
massivestar
star formation,
locatedlocatedroughly 400-430
pc from
the Sun pc from the Sun(Menten et al. 2007; Zucker et al. 2019; Großschedl et al. 2018). Spanning over 200 deg2 on the sky,
2 the region is home
(Menten et al. 2007;toZucker
et al. 2019; Großschedl et al. 2018). Spanning over 200 deg on the sky, the region is home
one of the largest and best-studied associations of O and B stars, the so-called Orion OB1 association (Bally 2008;
to one of the largest and best-studied associations of O and B stars, the so-called Orion OB1 association (Bally 2008;

School of Natural Sciences, Institute for Advanced Study, Princeton, NJ 08544, USA

Two saturated states of the vertical shear
instability
in protoplanetary disks with verticallyReceived 2022 November 2; Accepted 2022 December
6varying cooling times∗

E-mail: fukuhara.y.ab@m.titech.ac.jp

Yuya FUKUHARA,1,∗ Satoshi OKUZUMI,1 and Tomohiro ONO1,2

1
Departmentof Earth and PlanetaryTokyo
Institute of Technology,
2-12-1 Ookayama, Meguro,
Turbulence in protoplanetary disks plays
an importantroleSciences,in dustevolutionand planetesTokyo 152-8551, Japan
imal formation. The vertical shear instability(VSI) is one of the candidate hydrodynamic
2
School of Natural Sciences, Institute for Advanced Study, Princeton, NJ 08544, USA
mechanisms that can generate turbulencein the outer disk regions. The VSI requires∗
E-mail: fukuhara.y.ab@m.titech.ac.jp
rapid gas cooling in addition to vertical
shear.
linearanalysis suggests that
Received 2022
NovemberA2; Accepted2022stability
December 6
the VSI may not operate around the midplane where gas cooling is inefficient. In this
Abstract
study, we investigate the nonlinear outcomeof the VSI in disks with a linearly VSI-stable
Turbulence in protoplanetary disks plays an important role in dust evolution and planetesmidplane region. We perform two-dimensional
globalhydrodynamicalof animal formation. The
verticalshear instability (VSI) issimulations
one of the candidatehydrodynamic
mechanisms that can generate turbulence in the outer disk regions. The VSI requires
axisymmetric disk with vertically varying
cooling times. The vertical cooling time prorapid gas cooling in addition to vertical shear. A linear stability analysis suggests thatfile determines the thicknesses of thelinearlyandunstablethe VSI
may not VSI-stable
operate aroundmidplane
the midplane layer
where gascooling
is inefficient. In thisstudy,
we findinvestigate
nonlinearoutcomethe VSIin disks withstablea linearly VSI-stablelayers above and below the midplane.
Wethatthethethickness
ofof themidplane
midplane region. We perform two-dimensional global hydrodynamical simulations of anlayer determines the vertical structureof VSI-driven
turbulencein thenonlinearsat- time proaxisymmetricdisk with verticallyvarying coolingtimes.
The vertical coolingthe thicknessesof the(i)linearly
VSI-stablecharacterized
midplane layer and unstableurated state. We identify two typesfileofdetermines
final saturated
state:
T states,
layers above and below the midplane. We find that the thickness of the midplane stableby vertical turbulent motion penetrating
into thetheVSI-stablemidplanelayer,and in(ii)thepT
layer determines
vertical structureof VSI-driven
turbulence
nonlinear saturated
state.
We
identify
two
types
of
final
saturatedstate:(i)T
states,states, characterized by turbulent motion confined in the unstable layers. The pT statescharacterized
by vertical turbulent motion penetrating into the VSI-stable midplane layer, and (ii) pT
are realized when the midplane VSI-stablelayer isthan confined
two gasscale
heights.
states, characterized
by thicker
turbulent motionin theunstablelayers. The pT statesarerealized
whenthemidplane
VSI-stablelayer
isthickerthan
twoscale heights.
We also find that the VSI-driven turbulence is largely suppressed at all heights whengas
thealso find that the VSI-driven turbulence is largely suppressed at all heights when the
VSI-unstable region lying above andWe
below
the midplane is thinner than two gas scale
VSI-unstable region lying above and below the midplane is thinner than two gas scaleheights.
Wepredictpresent empiricalformulas thatthe strengthof VSI-driven turbulenceheights. We present empirical formulas
thatthe strengthof predict
VSI-driven
turbulence
as a function of the thicknesses of the unstable and stable layers. These formulas will
as a function of the thicknesses of theunstable and stable layers. These formulas will
be useful for studying how VSI-driven turbulence and dust grains controlling the disk
be useful for studying how VSI-driven
turbulenceanddust grains controlling the disk
cooling
efficiency evolve
simultaneously.
Key words: hydrodynamics — methods: numerical — protoplanetary disks — turbulencecooling efficiency evolve simultaneously.

Downloaded from https://academic.oup.com/pasj/article/75/1/233/6986070 by KASI user on 23 February 2023

Abstract

rom https://academic.oup.com/pasj/article/75/1/233/6986070 by KASI user on 23 February 2023

2

T
pT

∗

E-mail: fujita@p.s.osakafu-u.ac.jp

Received 2022 March 8; Accepted 2022 December 1

Abstract

Shinji FUJITA ,1,2,∗ Atsushi M. ITO,3 Yusuke MIYAMOTO ,4,5
Yasutomo KAWANISHI ,6 Kazufumi TORII,5 Yoshito SHIMAJIRI ,7
Atsushi NISHIMURA ,8 Kazuki TOKUDA ,9 Toshikazu OHNISHI ,2
Hiroyuki KANEKO ,5,10 Tsuyoshi INOUE,11 Shunya TAKEKAWA ,12
Mikito KOHNO ,13,14 Shota UEDA,2 Shimpei NISHIMOTO,2 Ryuki YONEDA,2
Kaoru NISHIKAWA,14 and Daisuke YOSHIDA14
Publications of the Astronomical Society of Japan (2023), Vol. 75, No. 1
1
Institute of Astronomy, Graduate School of Science, The University of Tokyo, 2-21-1 Osawa, Mitaka,∗E-mail: fujita@p.s.osakafu-u.ac.jp

m, CO 2-1
CLUMPFIND

Downloaded from htt

280
Machinelearning has been successfully applied in various field but whether it is a viable
Tokyo 181-0015, Japan
tool for determining the distance to molecular
clouds in the Galaxy is an open question.
2
Department of Physical Science, Graduate School of Science, Osaka Prefecture University, 1-1
In
the Galaxy, the kinematic distance is Gakuen-cho,commonly
employed
to represent
the distance toNaka-ku,
Sakai,
Osaka
599-8531,
Japan"
C The Author(s) 2023. Published by Oxford University Press on behalfof the Astronomical
Societyof Japan.
3
National
Institute twofor Fusion
Science (NIFS),
National Institutesof Natural
Sciences (NINS), 322-6
a molecular
cloud.
However,forthe inner
Galaxy,different
solutions,
i.e., the“Near”
All
rights reserved. For
permissions,please e-mail:journals.permissions@oup.com
Oroshi-cho, Toki, Gifu 509-5292, Japansolution and the “Far” solution, can be4 Department
derivedof simultaneously.
We attemptElectrical and Electronic Engineering,
Facultytoof construct
Engineering, Fukui University of
PPV
Technology,
3-6-1
Gakuen,Fukui,
Fukui
910-8505,Japana two-class (“Near” or “Far”) inference
model using a convolutional neural network
5
National Astronomical Observatory of Japan, National Institutes of Natural Sciences, 2-21-1 Osawa,(CNN), which is a form of deep learningMitaka,
thatTokyo
can181-8588,
captureJapanspatial features generally. In this6
RIKEN
Informationof
R&D
and Galactic
Strategy Headquarters,
2-2-2 Hikaridai, with
Seika-cho,study, we use the CO dataset in the first
quadranttheplane obtained
theSoraku-gun, Kyoto
619-0288,◦◦ Japan◦Nobeyama 45 m radio telescope (l = 62
–10Kyoritsu, |b| <
1 ). Inmodel,
we apply
the Fukuoka,
three-Fukuoka 807-8585,
7
KyushuUniversity,
1-8theJiyugaoka,
Yahatanishi-ku,
Kitakyushu,Japan
dimensional distribution (position–position–velocity)of the 12 CO (J = 1–0) emissions as
8Nobeyama Radio Observatory, National Astronomical Observatory of Japan, National Institutes ofthe main input. To train the model, a datasetwith “Near”
or “Far”
annotation
was384-1305,
createdNatural Sciences,462-2 Nobeyama,
Minamimaki,
Minamisaku, NaganoJapan
9
Department
of
Earth
and
Planetary
Sciences,
Facultyof
Science,
KyushuUniversity,from the H II region catalog of the infrared astronomy satellite WISE. Consequently, weNishi-ku, Fukuoka,
Fukuoka 819-0395, Japan
construct a CNN model with a 76% accuracy
rateon thetraining
Using the Joetsu, Niigata
10Graduate Schoolof Education,
JoetsuUniversity ofdataset.
Education, 1 Yamayashiki-machi,
943-8512, Japanproposed model, we determine the distanceto the molecular clouds identified by the11
Department of Physics, Konan University, 8-9-1 Okamoto, Higashinada-ku, Kobe, Hyogo 658-8501, Japan
CLUMPFIND algorithm. We found that12 Faculty
the massof molecular
clouds3-27-1
witha distanceofof Engineering,
Kanagawa University,
Rokkakubashi,
Kanagawa-ku,
Yokohama,
12
Kanagawa
221-8686,
Japan
a power-law distribution with an index of<8.15 kpc identified in the CO data follows
13
Astronomy Section,
Nagoya City Science Museum, 2-17-1 Sakae, Naka-ku, Nagoya, Aichi 460-0008,
approximately −2.3 in the mass range Japan
M > 103 M# . In addition, the detailed molecular
Departmentof Astrophysics,
University,
Chikusa-ku, Nagoya, Aichi 464-8602, Japangas distribution of the Galaxy, as seen14fromthe
GalacticNagoyaNorth
pole,Furo-cho,
was determined.

Downloaded from https://academic.oup.com/pasj/article/75/1/279/6988000 by KASI user on 23 February 2023

Faculty of Engineering, Kanagawa University, 3-27-1 Rokkakubashi, Kanagawa-ku, Yokohama,
Distance determination of molecular clouds in
Kanagawa 221-8686, Japan
13
Astronomy Section, Nagoya City Science Museum,
2-17-1
Sakae, Naka-ku,
Aichiplane
460-0008,
the first
quadrantof theNagoya,
Galactic
using
Japan
deep learning: I. Method and results
14
Department of Astrophysics, Nagoya University, Furo-cho, Chikusa-ku, Nagoya, Aichi 464-8602, Japan

2023

12

WISE

HII

East Asian Observatory (EAO), 660 N. Aohoku Place, University Park, Hilo, HI 96720, USA
Department of Astronomy, University of Florida, P.O. Box 112055, Gainesville, FL 32611-2055, USA
Received 2022 November
11; Formation
revised 2022ofDecember
accepted2022
7; published2022 December 28
Clustered
Massive7;Stars
withinanDecember
Ionized Rotating
Disk
9

Roberto Galván-Madrid1 , Qizhou Zhang2 , Andrés Izquierdo3,4 , Charles J. Law2 , Thomas Peters5, Eric Keto2,
Hauyu Baobab Liu6,7 , Paul T. P. Ho7,8 , Adam Ginsburg9 , and Carlos Carrasco-González1

Abstract
We present Atacama Large Millimeter/submillimeter Array observations with a 800 au resolution and radiativetransfer modeling of the inner part (r ≈ 6000 au) of the ionized accretion ﬂow around a compact star cluster information at the center of the luminous ultracompact H II region G10.6-0.4. We modeled the ﬂow with an
ionized Keplerian disk with and without radial motions in its outer part, or with an external Ulrich envelope. The
Markov Chain Monte Carlo ﬁts to the data give total stellar masses M from 120 to 200 M , with much smaller
ionized-gas masses Mion-gas = 0.2–0.25 Me. The stellar mass is distributed within the gravitational radius
Abstract
wheretheionizedgas
isbound.
The viewing inclination angle from the face-on
Rg ≈ 1000 to 1500 au,
We present Atacama Large Millimeter/submillimeter Array observations with a 800 au resolution and radiative−1transfer modelingthe inner part at(r ≈radii
6000 au)r ofaccretion ﬂow
a compact
star scluster
to around
vr,0 ≈
8.7 km, inor about the speed oforientation is i = 49°–56°.
Radialof motions>theRionizedg converge
formation at the center of the luminous ultracompact H II region G10.6-0.4. We modeled the ﬂow with ansound of ionized gas, ionizedindicating
thisisradial
marginallyat anmost.
additional
constraints on the
Keplerian thatdisk withand gaswithout
motions in its unboundouter part, or withexternalFrom
Ulrich envelope.
The
to 200 theMe, withmuch smaller
Markov
Chain Monte
Carlo ﬁts to thetotal stellarwe
massesMå from 120 thationizing-photon rate andfar-IR
luminosity
ofdatathegiveregion,
concludestellarcluster consists of a few
ionized-gas masses Mion-gas = 0.2–0.25 Me. The stellar mass is distributed within the gravitational radius
32–60
or onestar inrange
massesaccompanied
a population of lowermassive stars with Mstar
1000
to 1500Mau,
the ionizedgasthisis bound.
The ofviewing
inclination
angle from the by
face-onRg ≈=e,whereorientation is i = 49°–56°. Radial motions at radii r > Rg converge to vr,0 ≈ 8.7 km s−1, or about the speed ofmass stars. Any active
accretion of ionized gas onto the massive (proto)stars is residual. The inferred clustersound of ionized gas, indicating that this gas is marginally unbound at most. From additional constraints on theionizing-photon rateto
andthatfar-IRreported
luminosity ofat
thesimilar
region, we scales
conclude thatcluster consistsof a few
density is very large, comparable
in thethestellar
Galactic
center.
Stellar interactions are
massive stars with Mstar = 32–60 Me, or one star in this range of masses accompanied by a population of lowerlikely to occur withinmass
thestars.
next
years.
Anymillion
active accretionof ionized gas onto the massive (proto)stars is residual. The inferred cluster
1

Instituto de Radioastronomía y Astrofísica, Universidad Nacional Autónoma de México, Morelia, Michoacán 58089, México
2
Center for Astrophysics | Harvard & Smithsonian, 60 Garden St., Cambridge, MA 02138, USA
3
European Southern Observatory, Karl-Schwarzschild-Str. 2, D-85748 Garching bei München, Germany
4
Leiden Observatory, Leiden University, P.O. Box 9513, NL-2300 RA Leiden, The Netherlands
5
Max-Planck-Institut für Astrophysik, Karl-Schwarzschild-Str. 1, D-85748 Garching, Germany
6
Physics Department, National Sun Yat-Sen University, No. 70, Lien-Hai Road, Kaosiung City 80424, Taiwan, R.O.C.
7
Institute of Astronomy and Astrophysics, Academia Sinica, 11F of Astronomy-Mathematics Building, AS/NTU No.1, Sec. 4, Roosevelt Rd, Taipei 10617, Taiwan,
R.O.C.
8
East Asian Observatory (EAO), 660 N. Aohoku Place, University Park, Hilo, HI 96720, USA
9
Department of Astronomy, University of Florida, P.O. Box 112055, Gainesville, FL 32611-2055, USAe
Received 2022 November 11; revised 2022 December 7; accepted 2022 December 7;åpublished 2022 December 28

density is very large, comparable to that reported at similar scales in the Galactic center. Stellar interactions are

Uniﬁed Astronomy Thesaurus
II regions(694)
likely to occurconcepts:
within the nextH
millionyears.

UCHII G10.6-0.4

Uniﬁed Astronomy Thesaurus concepts: H II regions (694)

1. Introduction

1. Introduction

Stars with masses greater than 30 Me are rare among thestellar population
galaxies,
they playa main therole inamong
Stars with
massesgreaterthan 30in49-56
M
e arebutrare
0.2-0.25
Msun
shaping the dynamical evolution of the interstellar medium
stellar population in galaxies,
but theyplay
a mainrole
inthrough feedback
from theirradiation,
ionization,and supernovae (for a review
feedback
effects in starmedium
formation, seeshaping the dynamical evolutionof oftheinterstellar
Krumholz et al. 2014). Since these very massive stars arethrough feedback from theirradiation,
ionization,andsupermostly formed
in dense stellar
clusters, their
violent
death maylead to the formation of black holes or neutron-star binary
novae (for a review of feedback
effects in star formation, seesystems that give rise to gravitational-wave events (e.g.,
Mapelli very2018). Massive
protostars
must areaccreteKrumholz et al. 2014). Giacobbo
Since &these
massivestars
more than 30 Me to become early O-type stars, but they begin
mostly formed in dense stellar
clusters,
violent
maycore hydrogenburningtheirwell before
reachingdeath
their ﬁnal
massesr<6000 au
clustered environments
120-200(Goddi
Msun et al. 2020). In a clusteredscenario, and after the protostars have reached several tens of
Vr of extreme-UV
8.7 km/s
solar masses, the amountphotons is large

clustered environments (Goddi et al. 2020). In a clusteredscenario, and after the protostars have reached several tens ofsolar masses, the amount of extreme-UV photons is large
enough to ionize the otherwise molecular accretion ﬂow(s) andthe clump in which they reside (Keto 2003, 2007; Peters et al.
2010a, 2010b; Klaassen et al. 2018). This raises several
outstanding questions: How does star formation proceed insuch an environment? Is it through an ionized accretion disk?
Does this disk surround a single or a multiple stellar system?
When does active (proto)stellar accretion stop due to theincreasing photoionizing feedback?
We present the ﬁrst spatially resolved observations of an
accretion ﬂow in ionized gas around a dense star cluster inenough to ionize the otherwise molecular accretion ﬂow(s) and
100 in which they reside (Keto 2003, 2007; Peters et al.
the clump
2010a, 2010b; Klaassen et al. 2018). This raises several
outstanding questions: How does star formation proceed insuch an environment? Is it through an ionized accretion disk?