Star Formation Newsletter Seminar
11/18
古屋#356 9-15

Chemical evolution in ices on drifting, planet-forming pebbles

Christian Eistrup, Thomas Henning F Planets and their atmospheres are built from gas and solid material in protoplanetary disks. Recent results suggest that solid material like pebbles may contribute significantly to building up planetaryatmospheres. In order to link observed exoplanet atmospheres and their compositions to their formation histories, it is importantto understand how icy pebbles may change their composition when they drift radially inwards in disks. Our goal is to model
the compositional evolution of ices on pebbles as they drift in disks, and track how their chemical evolution en-route changes
the ice composition relative to the ice composition of the pebbles in the region where they grew from micron-sized grains. A
state-of-the-art chemical kinetics code is utilised for modelling chemical evolution. This code accounts for the time-evolving
sizes of the solids that drift. Chemical evolution is modelled locally for 0.1Myr at two starting radii, with the micron-sized
solids growing into pebbles simultaneously. The pebbles and local gas, isolated as a parcel, is then exposed to changing physicalconditions, intended to mimic the pebbles drifting inwards in the disk midplane, moving to 1 AU on three diﬀerent timescales.
A modelling simplification is that the pebbles are not moved through, and exposed to new gas, but stay in the same chemicalgas surroundings in all models. For ice species with initial abundances relative to hydrogen of > 10 5 , such as H2 O, CO2 ,CH3 OH and NH3 , the abundances change by less than 20timescales (10kyr and 100kyr). For less abundant ice species, and thelongest drift timescale (1Myr), the changes are larger. Pebble drift chemistry generally increases the ice abundances of CO2 ,HCN and SO, at the expense of decreasing abundances of other volatile molecules.

Gap Opening and Inner Disk Structure in the Strongly Accreting Transition Disk of
DM TauLogan Francis, Nienke van der Marel, Doug Johnstone, Eiji Akiyama, Simon Bruderer, Ruobing Dong, Jun
Hashimoto, Hauyu Baobab Liu, Takayuki Muto, Yi Yang F Large inner dust gaps in transition disks are frequently
posited as evidence of giant planets sculpting gas and dust in the disk, or the opening of a gap by photoevaporative winds.
Although the former hypothesis is strongly supported by the observations of planets and deep depletions in gas within the gap
some disks, many T Tauri stars hosting transition disks accrete at rates typical for an undepleted disk, raising the questionof how gap opening occurs in these objects. We thus present an analysis of the structure of the transition disk around the
T Tauri star DM Tau, which is strongly accreting (⇠ 10 8.3 M yr 1 ) and turbulent (↵ = 0.078 ± 0.02). Using the DALI
thermochemical code, we fit disk models to simultaneously reproduce the accretion rate, high level of turbulence, the gas tracedby ALMA band 6 observations of 12 CO, 13 CO, and C18 O J=2–1 lines, and the observed dust emission from the mm continuum
and spectral energy distribution. We find a shallow depletion in gas surface density of ⇠ 10 relative to the outer disk and a gas
rich inner disk is consistent with the observations. The planet mass of < 1 MJup implied by the gap depth is in tension withpredictions for dust trapping in a highly viscous disk, which requires a more massive planet of of ⇠ 10MJup . Photoevaporativemodels including a dead zone can qualitatively reproduce some features of the DM Tau disk, but still struggle to explain thehigh accretion rates and the observed mm continuum flux.

Cosmic-ray ionization rate in protoplanetary disks with sheared magnetic fields

Yuri I. Fujii, Shigeo S. Kimura F We investigate the eﬀects of magnetic field configurations on the ionization rate bycosmic rays in protoplanetary disks. First, we consider cosmic-ray propagation from the interstellar medium (ISM) to theprotoplanetary disks and showed that the cosmic-ray density around the disk should be 4 times lower than the ISM value.
Then, we compute the attenuation of cosmic rays in protoplanetary disks. The magnetic fields in the disk are stretched to theazimuthal directions, and cosmic rays need to detour while propagating to the midplane. We show that the detouring eﬀectively
enhances the column density by about two orders of magnitudes. In the case of the disk around IM lup, this increases theionization rate over an order of magnitude for r & 100 au. On the other hand, for r . 100 au, the cosmic rays are shieldedat the disk midplane while the ionization rate is also enhanced at z ⇠ 2H. Our results are consistent with the recent ALMA
observation that indicates the radial gradient in the cosmic-ray ionization rate. The elevated ionization rate in the outer radiiof disks may activate the magnetorotational instability that was thought to be suppressed due to ambipolar diﬀusion. These
results will have a strong influence on the dynamical and chemical evolutions of protoplanetary disks.

Monitoring accretion rate variability in the Orion Nebula Cluster with the Wendelstein
Wide Field Imager

S. Flaischlen, T. Preibisch, M. Kluge, C. F. Manara, B. Ercolano F The understanding of the accretion process
has a central role in the understanding of star and planet formation. We aim to test how accretion variability influences
previous correlation analyses of the relation between X-ray activity and accretion rates, which is important for understanding
the evolution of circumstellar disks and disk photoevaporation. We monitored accreting stars in the Orion Nebula Cluster from
November 24, 2014, until February 17, 2019, for 42 epochs with the Wendelstein Wide Field Imager in the Sloan Digital Sky
Survey u’g’r’ filters on the 2 m Fraunhofer Telescope on Mount Wendelstein. Mass accretion rates were determined from thevariability on timescales from weeks to 2 years. The variability has likely no significant influence on a correlation analysis of the
X-ray luminosity and the mass accretion rate observed at diﬀerent times when the sample size is large enough. The observed
anticorrelation between the X-ray luminosity and the mass accretion rate predicted by models of photoevaporation-starved
accretion is likely not due to a bias introduced by diﬀerent observing times.

Methanol deuteration in high-mass protostars

M. L. van Gelder, J. Jaspers, P. Nazari, A. Ahmadi, E. F. van Dishoeck, M. T. Beltrán, G. A. Fuller, Á.
Sánchez-Monge, P. Schilke F The deuteration of molecules forming in the ices such as methanol (CH3 OH) is sensitive
to the physical conditions during their formation in dense cold clouds and can be probed through observations of deuterated
methanol in hot cores. Observations with ALMA containing transitions of CH3 OH, CH2 DOH, CHD2 OH, 13 CH3 OH, andCH18
3 OH are investigated. The column densities of CH2 DOH, CHD2 OH, and CH3 OH are determined for all sources, where thecolumn density of CH3 OH is derived from optically thin 13 C and 18 O isotopologues. Consequently, the D/H ratio of methanol
is derived taking statistical eﬀects into account. Singly deuterated methanol (CH2 DOH) is detected toward 25 of the 99 sourcesin our sample of the high-mass protostars. Including upper limits, the (D/H)CH3 OH ratio inferred from NCH2 DOH /NCH3 OH wasderived for 38 of the 99 sources and varies between ⇠ 10 3 10 2 . Including other high-mass hot cores from the literature,
the mean methanol D/H ratio is 1.1 ± 0.7 ⇥ 10 3 . This is more than one order of magnitude lower than what is seen for
low-mass protostellar systems (2.2 ± 1.2 ⇥ 10 2 ). Doubly deuterated methanol (CHD2 OH) is detected toward 11 of the 99
sources. Including upper limits for 15 sources, the (D/H)CH2 DOH ratios derived from NCHD2 OH /NCH2 DOH are more than twoorders of magnitude higher than (D/H)CH3 OH with an average of 2.0 ± 0.8 ⇥ 10 1 which is similar to what is found for low-masssources. Comparison with literature GRAINOBLE models suggests that the high-mass prestellar phases are either warm (> 20
K) or live shorter than the free-fall timescale. In contrast, for low-mass protostars, both a low temperature of < 15 K and aprestellar phase timescale longer than the free-fall timescale are necessary.

Methanol deuteration in high-mass protostars

M. L. van Gelder, J. Jaspers, P. Nazari, A. Ahmadi, E. F. van Dishoeck, M. T. Beltrán, G. A. Fuller, Á.
Sánchez-Monge, P. Schilke F The deuteration of molecules forming in the ices such as methanol (CH3 OH) is sensitive
to the physical conditions during their formation in dense cold clouds and can be probed through observations of deuterated
methanol in hot cores. Observations with ALMA containing transitions of CH3 OH, CH2 DOH, CHD2 OH, 13 CH3 OH, andCH18
3 OH are investigated. The column densities of CH2 DOH, CHD2 OH, and CH3 OH are determined for all sources, where thecolumn density of CH3 OH is derived from optically thin 13 C and 18 O isotopologues. Consequently, the D/H ratio of methanol
is derived taking statistical eﬀects into account. Singly deuterated methanol (CH2 DOH) is detected toward 25 of the 99 sourcesin our sample of the high-mass protostars. Including upper limits, the (D/H)CH3 OH ratio inferred from NCH2 DOH /NCH3 OH wasderived for 38 of the 99 sources and varies between ⇠ 10 3 10 2 . Including other high-mass hot cores from the literature,
the mean methanol D/H ratio is 1.1 ± 0.7 ⇥ 10 3 . This is more than one order of magnitude lower than what is seen for
low-mass protostellar systems (2.2 ± 1.2 ⇥ 10 2 ). Doubly deuterated methanol (CHD2 OH) is detected toward 11 of the 99
sources. Including upper limits for 15 sources, the (D/H)CH2 DOH ratios derived from NCHD2 OH /NCH2 DOH are more than twoorders of magnitude higher than (D/H)CH3 OH with an average of 2.0 ± 0.8 ⇥ 10 1 which is similar to what is found for low-masssources. Comparison with literature GRAINOBLE models suggests that the high-mass prestellar phases are either warm (> 20
K) or live shorter than the free-fall timescale. In contrast, for low-mass protostars, both a low temperature of < 15 K and aprestellar phase timescale longer than the free-fall timescale are necessary.

CO(J = 1-0) Observations toward the Filamentary Cloud in the Galactic Region of
153.60 6 l 6 156.50 and 1.85 6 b 6 3.50

Weihua Guo, Xuepeng Chen, Jiancheng Feng, Li Sun, Shiyu Zhang, Chen Wang, Yang Su, Yan Sun, Qingzeng
Yan, Shaobo Zhang, Xin Zhou, MiaoMiao Zhang, Min Fang, Ji Yang F We present observations of J=1-0 transitionlines of 12 CO, 13 CO, and C18 O towards the Galactic region of 153.60 6 l 6 156.50 and 1.85 6 b 6 3.50 , using the Purple
Mountain Observatory (PMO) 13.7 m millimeter telescope. Based on the data, one main filament and five sub-filaments arefound together as a network structure in the velocity interval of [ 42.5, 30.0] km s 1 . The kinematic distance of this molecular
cloud (MC) is estimated to be ⇠ 4.5 kpc. The median length, width, excitation temperature, line mass of these filaments are⇠ 49 pc, ⇠ 2.9 pc, ⇠ 8.9 K, and ⇠ 39 M pc 1 , respectively. The velocity structures along these filaments exhibit oscillatory
patterns, which are likely caused by the fragmentation or accretion process along these filaments. The maximum accretion rateis estimated to be as high as ⇠ 700 M pc 1 . A total of ⇠ 162 clumps and ⇠ 103 young stellar objects (YSOs) are identified inthis region. Most of the clumps are in gravitationally bound states. Three HII regions (G154.359+2.606, SH2-211, SH2-212)
are found to be located in the apexes of the filaments. Intense star forming activities are found along the entire filamentary
cloud. The observed results may help us to better understand the link between filaments and massive star formation.

Methanol deuteration in high-mass protostars
M. L. van Gelder1 , J. Jaspers1 , P. Nazari1 , A. Ahmadi1 , E. F. van Dishoeck1, 2 , M. T. Beltrán3 ,G. A. Fuller4, 5 , Á. Sánchez-Monge5 , and P. Schilke5
1

Leiden Observatory, Leiden University, PO Box 9513, 2300RA Leiden, The Netherlands
e-mail: vgelder@strw.leidenuniv.nl
2
Max Planck Institut für Extraterrestrische Physik (MPE), Giessenbachstrasse 1, 85748 Garching, Germany
3
INAF-Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Firenze, Italy
4
Jodrell Bank Centre for Astrophysics, Department of Physics and Astronomy, University of Manchester, Oxford Road, Manchester,M13 9PL, UK
5
I. Physikalisches Institut, Universität zu Köln, Zülpicher Str.77, 50937, Köln, Germany
Received XXX; accepted XXX
ABSTRACT
Context. The deuteration of molecules forming in the ices such as methanol (CH3 OH) is sensitive to the physical conditions during

their formation in dense cold clouds and can be probed through observations of deuterated methanol in hot cores.
Aims. The aim is to determine the D/H ratio of methanol for a large sample of 99 high-mass protostars and to link this to the physicalconditions during the formation of methanol in the prestellar phases.
Methods. Observations with the Atacama Large Millimeter/submillimeter Array (ALMA) containing transitions of CH3 OH,
CH2 DOH, CHD2 OH, 13 CH3 OH, and CH18
3 OH are investigated. The column densities of CH2 DOH, CHD2 OH, and CH3 OH are determined for all sources, where the column density of CH3 OH is derived from optically thin 13 C and 18 O isotopologues. Consequently,
the D/H ratio of methanol is derived taking statistical e↵ects into account.
Results. Singly deuterated methanol (CH2 DOH) is detected at the 3 level toward 25 of the 99 sources in our sample of the highmass protostars. Including upper limits, the (D/H)CH3 OH ratio inferred from NCH2 DOH /NCH3 OH was derived for 38 of the 99 sources andvaries between ⇠ 10 3 10 2 . Including other high-mass hot cores from the literature, the mean methanol D/H ratio is 1.1 ± 0.7 ⇥ 10 3 .
This is more than one order of magnitude lower than what is seen for low-mass protostellar systems (2.2 ± 1.2 ⇥ 10 2 ). Doubly
deuterated methanol (CHD2 OH) is detected at the 3 level toward 11 of the 99 sources. Including upper limits for 15 sources, the(D/H)CH2 DOH ratios derived from NCHD2 OH /NCH2 DOH are more than two orders of magnitude higher than (D/H)CH3 OH with an averageof 2.0 ± 0.8 ⇥ 10 1 which is similar to what is found for low-mass sources. Comparison with literature GRAINOBLE models suggeststhat the high-mass prestellar phases are either warm (> 20 K) or live shorter than the free-fall timescale. In contrast, for low-mass
protostars, both a low temperature of < 15 K and a prestellar phase timescale longer than the free-fall timescale are necessary.
Conclusions. The (D/H)CH3 OH ratio drops by more than an order of magnitude between low-mass and high-mass protostars due toeither a higher temperature during the prestellar phases or shorter prestellar phases. However, successive deuteration toward CHD2 OH
seems equally e↵ective between low-mass and high-mass systems.
Key words. astrochemistry – stars: formation – stars: protostars – techniques: interferometric – ISM: molecules• CH2DOHから求めたD/H⽐は, CH3OHから求めた値よりも⼆桁⾼く, low-massで観測されている値に近い.

mation.
The gaseous atomic D/H ratio can be increased in theprestellar phases through the exothermic reaction (Watson 1974;
Aikawa & Herbst 1999; Ceccarelli et al. 2014),++
H3 + HD ⌦ H2 D + H2 +

E,

(1)

where E = 232 K. Since in the cold (. 20 K) prestellar cores+
the backward reaction in Eq. (1) is less efficient, H2 D is enhanced and the atomic D/H ratio in the gas phase can be ef+
fectively increased through dissociative recombination of H2 Dwith free electrons. Moreover, gaseous CO is the main destructor++of H3 and H2 D (Brown & Millar 1989; Roberts et al. 2003) and4
3
thus the heavy CO freeze-out in dense (& 10 cm ) prestellar
cores additionally stimulates the increase of the gaseous atomic
D/H ratio. In turn, the enhanced atomic D/H ratio in the gas cantranslate into a higher D/H ratio of molecules forming in the ices
Article number, page 1 of 17

A&A proofs: manuscript no. ALMAGAL_D_H_methanol_paper
Table C.1. continued.

SourceType(1)(D/H)CH3 OH(D/H)CH2 DOH

Refs

G034-C1(MM9)I20293-WC
I22134-G
I22134-B

HMSC
HMSC
HMSC
HMSC

<8.7(-4)<1.9(-3)<1.7(-3)<6.7(-3)

–
–
–
–

27
27
27
27

Notes. The (D/H)CH3 OH and (D/H)CH2 DOH ratios are either directly taken from the reported literature or computed using Eqs. (2) and (3) using thecolumn densities from the reported literature. A 30% uncertainty was assumed in cases where no uncertainty was reported.
(1)
LMP: low-mass protostar, IMS: intermediate-mass protostar, HMS: high-mass protostar, LMPC: low-mass prestellar core, HMSC: high-mass
starless core.

Fig. 1. Integrated intensity maps of the CH3 OH 80,8 71,6 (Eup = 97 K, left), CH2 DOH 171,16 e0 170,17 e0 (Eup = 336 K, middle), and CHD2 OH
70,1 e1 61,1 e1 (Eup = 74 K, right) lines for 881427. The color scale is shown on top of each image. The image is integrated over [-5,5] km s 1with respect to the Vlsr of source A. The white vertical line in the colorbar indicates the 3 threshold. The source positions based on the continuum
emission are indicated with the white stars. The white ellipse in the lower right of each image depicts the beam size and in the lower left a physicalscale bar is displayed.
A&A proofs: manuscript no. ALMAGAL_D_H_methanol_paper

tra are extracted from the same positions as van Gelder et al.
(2022). In G023.3891+00.1851, the emission of CH2 DOH peakso↵set by about half the beam (⇠ 0.600 ) and therefore the spectrum is extracted from the peak of CH2 DOH 171,16 e0 170,17 e0(Eup = 336 K). For all other sources, the the peak in CH3 OH coincides with the peaks of CH2 DOH and CHD2 OH. For sources
also included by Nazari et al. (2022), our spectral extraction
locations are the same as theirs, which are extracted from thepeak position of the CH3 CN 124 114 integrated intensity maps,
except for 721992 and G023.3891+00.1851 where the CH3 OH
80,8 71,6 and CH2 DOH 171,16 e0 170,17 e0 emission peaks o↵set
from the CH3 CN 124 114 emission by about 100 . It is important
to note that these spectral extraction positions are di↵erent by
up to 100 from van Gelder et al. (2022) who extracted their spectra from the peak continuum pixel for all sources. Therefore, thecolumn densities derived in this work may deviate from theirs.
The reason why our spectra are extracted from the peak pixel of
CH3 OH is to have the highest signal-to-noise in methanol lines
and its isotopologues the extracted spectra.

on the work of Pearson et al. (2012). The line list of CHD2 OH
is taken from Drozdovskaya et al. (2022), which is mostly basedon the the work of Coudert et al. (2021).

Only transitions with Eup
50 K are used to derive thecolumn densities since lines with lower Eup likely include also
emission from non-thermally desorbed methanol and emission
possibly related to outflows. For 13 CH3 OH, the 141,13 132,12(Eup = 254 K) transition gives the only constraint on the column
density for many sources as the other transitions su↵er from severe line blending. Similarly, for CH1870,7
3 OH only the 81,8
(Eup = 86 K) and 141,14 132,12 (Eup = 239 K) transitions provide constraints on the column density as well as some information on the excitation temperature. Furthermore, for CH2 DOH
the 52,4 e1
41,5 e1 (Eup = 59 K) line, as well as several
other lines, have rather low Einstein Aij values (< 10 5 s 1 )and are often blended with other COMs. Moreover, the spectroscopy of the CH2 DOH 181,17 o1 182,17 e0 line is unreliable
and shows large discrepancies in Aij between the JPL catalog entry (Aij = 1.8⇥10 5 s 1 ) and that derived by Coudert et al. (2014,
Aij = 8.9 ⇥ 10 7 s 1 ) and is therefore also excluded from the
2.2. Deriving the column densities
analysis. Consequently, the 171,16 e0 170,17 e0 (Eup = 336 K)
The column densities of all methanol isotopologues are derived transition of CH2 DOH provided the best constraint on the colusing the spectral analysis tool CASSIS2 (Vastel et al. 2015) un- umn density of CH2 DOH. However, although the 51,5 e0 41,4 e0
der the assumption of local thermodynamic equilibrium (LTE). (Eup = 36 K) transition is excluded from the fitting, it can pro13 DOH 17
Fig.line
2. Spectral
fits
of CH
e0 CH
1718
e0 (E
= 336 K,vide
top row)
and CHDone1 61,1 etemperaturerow) asfor
information
the70,1
excitation
ofbottom
CH2 DOH
The
lists ofline
CH
and
OH
are
2CH3 OH,1,16
up taken
2 OH
1 (E up = 74 K,
3 OH,
30,17
3
881427A(left), 881427B
andet881427C(right).
TheEndresdata corrected
for the LTE
Vlsr areshown
in black
the best fit for
K is
ex = 150
the best-fit
model
shouldnotandoverproduce
thisTline.
Lastly,(Mülleral. 2001,
2005;
from
the CDMS
catalog(middle),
in red.
etshown
al. 2016).
These entries include the first three (CH3 OH) and for CHD2 OH, the 70,1 e1 61,1 e1 (Eup = 74 K) transition is the13

18

References. 1: van Gelder et al. (2020); 2: Appendix D; 3: Lee et al. (2019a); 4: Bianchi et al. (2017b); 5: Manigand et al. (2020); 6:
Drozdovskaya et al. (2022); 7: Jørgensen et al. (2018); 8: Taquet et al. (2019); 9: Jacobsen et al. (2019); 10: Yang et al. (2020); 11: MartínDoménech et al. (2019); 12: Martín-Doménech et al. (2021); 13: Chahine et al. (2022); 14: Hsu et al. (2022); 15: Lee et al. (2019b); 16: Ligterink
et al. (2021); 17: Fuente et al. (2014); 18: Ospina-Zamudio et al. (2018); 19: Belloche et al. (2016); 20: Müller et al. (2016); 21: Bøgelund et al.
(2018); 22: Neill et al. (2013); 23: van der Walt et al. (2021); 24: Ambrose et al. (2021); 25: Lattanzi et al. (2020); 26: Bizzocchi et al. (2014); 27:
Fontani et al. (2015).

Appendix D: Doubly deuterated methanol in B1-c, Serpens S68N, and B1-bS
Using the database entry of CHD2 OH provided by Drozdovskaya et al. (2022), transitions from CHD2 OH can also be searched
for in a few COM-rich low-mass protostars. Here, this is done for B1-c, Serpens S68N (hereafter S68N), and B1-bS from the2017.1.01174.S ALMA program. The content of oxygen-bearing COMs for these sources was presented by van Gelder et al. (2020),
but no public database entry was yet available for CHD2 OH at that time.
Only one strong transition of CHD2 OH (71,2 o1 70,1 e0 , Eup = 68 K) is available for these sources which lies on the very edgeof the covered frequency range. For B1-c and S68N, this transition is detected at the 3 level, but given that only half the line isobserved this detection is still tentative. Using the same method for deriving the column densities as used by van Gelder et al. (2020)
and assuming an excitation temperature of 150 K, we derive column densities of 2.0±0.6⇥1016 cm 2 for B1-c, 7.2±2.7⇥1015 cm 2
for S68N, and < 1.7 ⇥ 1015 cm 2 for B1-bS. The FWHM was fixed to the average FWHM of those sources of 3.2km s 1 , 5.5km s 1 ,and 1.0km s 1 , respectively (van Gelder et al. 2020). The resulting fits are shown in Fig. D.1. Using these derived column densities
and those reported for CH3 OH by van Gelder et al. (2020), the (D/H)CH2 DOH ratios for B1-c and S68N (for B1-bS, both CH2 DOH
and CHD2 OH are not detected) are shown in Fig. 4 and agree very well with those derived for other low-mass sources as well aswith the high-mass sources.

Fig. D.1. Spectral line fits of CHD2 OH 71,2 o1 70,1 e0 (Eup = 68 K) line for B1-c (left), Serpens S68N (middle), and B1-bS (right). The data
corrected for the Vlsr are shown in black and the best fit for T ex = 150 K is shown in red.

Gelder et al. (2020). A grid of N and the full width at half
maximum (FWHM) of the line is set and a model spectrum iscomputed for each grid point assuming LTE conditions. The size 3. Resultsof the emitting region is fixed to the size of the beam (see Ap- The derived column densities of all isotopologues are presented
pendix B). Blended lines are excluded from the fitting procedure in Table B.1 for the reported excitation temperature. In Fig. 2, theand similarly broad lines (FWHM 10 km s 1 ) and lines with best-fit models to the CH2 DOH 171,16 e0 170,17 e0 and CHD2 OH
Eup  50 K are excluded in the fit to exclude any emission pos- 70,1 e1 61,1 e1 lines are presented for three hot cores in 881427
sibly related to outflows. The best-fit column density and the 2(see Fig. 1). Toward 25 sources, at least one clean unblended line
uncertainty are computed from the grid for each isotopologue. of CH2 DOH is detected at the 3 level, allowing for the deterThe main contributors to the uncertainty of N are the uncer- mination of the column density. For the remaining 74 sources
tainty on the flux calibration of ALMA (assumed to be 10%) where no (unblended) transitions of CH2 DOH are detected, theand the assumed excitation temperature. However, changing the 3 upper limit is reported. For CHD2 OH, the column density
excitation temperature in the 100-300 K range leads to at most could be determined for 11 sources. In Table B.1, the column
a factor 3 variation in the derived column densities. For several densities of 13 CH3 OH, CH18 OH, and CH3 OH are also reported.
3
1
sources (e.g., 705768), the lines are broad (> 7 km s ) making
The column densities of CH2 DOH are generally between one
automated line fitting complicated. For these sources the column and three orders of magnitude lower than those of CH3 OH, seedensity is estimated using the by eye fitting method of Nazari also Fig. 3. Furthermore, the column densities of CHD2 OH are
et al. (2021). In this case, a 50% uncertainty on the column den- about a factor 3–10 lower than that of CH2 DOH, see Fig. 4. In
sity is assumed. Moreover, the sources 101899, 615590, 865468, order to translate the column density ratios to the D/H ratios,and G345.5043+00.3480 showed line profiles consisting of mul- statistical weighting has to be taken into account since a deuFig. 3. The (D/H)CH3 OH ratio (squares) derived from the N/NCH3 OH for low-mass,intermediate-mass,
andeachhigh-mass
protostellarissystems
CH2 DOH
tiple
components.
The
columndensityofcomponent
deincluding data from both this study and the literature (see Appendix C for references). Only D/H ratios derived from interferometric observations terium atom has a three times higher probability to land in the
rived andreported fromseparately.
are included to minimize e↵ects of beam dilution and to exclude
any contribution
larger scales. Also data for low-mass prestellar cores and CH3 group compared to the OH group. Therefore, the D/H rahigh-mass starless cores (circles) from the literature (including observations with single dish telescopes) are included. Upper limits are presented tios of CH OH and CH DOH can be derived through,3
2
For all sources, the column density of CH3 OH is derived
as arrows.
18
from CH18
OH
and,
when
no
lines
originating
from
CH
3
3 OH
(2)
NCH2 DOH /NCH3 OH = 3(D/H)CH3 OH ,13
1213
were detected, from CH3 OH. The adopted C/ C and N(3)
CHD2 OH /NCH2 DOH = (D/H)CH2 DOH .
16
18
O/ O ratios are dependent on the galactocentric distance andare determined using the relations of Milam et al. (2005) and The derived D/H ratios are also listed in Table B.1.
Wilson & Rood (1994), respectively. In cases where only upper
The resulting (D/H)CH3 OH and (D/H)CH2 DOH ratios are prelimits on the column densities of both 13 CH3 OH and CH18
sented in Figs. 3 and 4, respectively. Including upper limits, a
3 OH
could be derived, the range in NCH3 OH was calculated by set- (limit on the) (D/H)CH3 OH and (D/H)CH2 DOH ratios could be deting the 3 upper limit based on scaling the 3 upper limit of rived for 38 and 26 of the 99 studied sources, respectively. Be13
CH3 OH and the lower limit based on the main isotopologue. sides the ALMAGAL sources, also other classical high-mass hot
Article number, page 4 of 17

Gelder et al. (2020). A grid of N and the full width at half
maximum (FWHM) of the line is set and a model spectrum iscomputed for each grid point assuming LTE conditions. The size 3. Resultsof the emitting region is fixed to the size of the beam (see Ap- The derived column densities of all isotopologues are presented
pendix B). Blended lines are excluded from the fitting procedure in Table B.1 for the reported excitation temperature. In Fig. 2, theand similarly broad lines (FWHM 10 km s 1 ) and lines with best-fit models to the CH2 DOH 171,16 e0 170,17 e0 and CHD2 OH
Eup  50 K are excluded in the fit to exclude any emission pos- 70,1 e1 61,1 e1 lines are presented for three hot cores in 881427
sibly related to outflows. The best-fit column density and the 2(see Fig. 1). Toward 25 sources, at least one clean unblended line
uncertainty are computed from the grid for each isotopologue. of CH2 DOH is detected at the 3 level, allowing for the deterThe main contributors to the uncertainty of N are the uncer- mination of the column density. For the remaining 74 sources
tainty on the flux calibration of ALMA (assumed to be 10%) where no (unblended) transitions of CH2 DOH are detected, theand the assumed excitation temperature. However, changing the 3 upper limit is reported. For CHD2 OH, the column density
excitation temperature in the 100-300 K range leads to at most could be determined for 11 sources. In Table B.1, the column
a factor 3 variation in the derived column densities. For several densities of 13 CH3 OH, CH18 OH, and CH3 OH are also reported.
3
1
sources (e.g., 705768), the lines are broad (> 7 km s ) making
The column densities of CH2 DOH are generally between one
automated line fitting complicated. For these sources the column and three orders of magnitude lower than those of CH3 OH, seedensity is estimated using the by eye fitting method of Nazari also Fig. 3. Furthermore, the column densities of CHD2 OH are
et al. (2021). In this case, a 50% uncertainty on the column den- about a factor 3–10 lower than that of CH2 DOH, see Fig. 4. In
sity is assumed. Moreover, the sources 101899, 615590, 865468, order to translate the column density ratios to the D/H ratios,and G345.5043+00.3480 showed line profiles consisting of mul- statistical weighting has to be taken into account since a deuFig. 4. The (D/H)CH2 DOH ratio derived from the NCHD2 OHtiple/NCH2 DOH
for low-mass The
and high-masssystems
data frombothcomponents.
columnprotostellar
density ofeach including
component
is dethis study and the literature (see Appendix C for references). Upper and lower limits are presented as arrows. Only D/H ratios derived from terium atom has a three times higher probability to land in therived
anddilution
reportedinterferometric observations are included to minimize e↵ectsof beamand toseparately.
exclude any contribution from larger scales. The (D/H)CH2 DOH CH3 group compared to the OH group. Therefore, the D/H raof IRAS 2A and IRAS 4A were derived using older spectroscopic data (Taquet et al. 2019).
For all sources, the column density of CH3 OH is derived tios of CH3 OH and CH2 DOH can be derived through,18
from CH18
OH
and,
when
no
lines
originating
from
CH
3
3 OH
(2)
NCH2 DOH /NCH3 OH = 3(D/H)CH3 OH ,131213
cores such as Sgr B2(N2) (Belloche et al. 2016), NGC
6334I
OH
or
the
derived(D/H)
andprothedetectionofCHwere detected, from CH32OH. The adopted CH
C/3 OHC and N(3)(Bøgelund et al. 2018), Orion KL (Neill et al. 2013),16
and CygXtostellar parameters such as Lbol and envelope mass is present
CHD2 OH /NCH2 DOH = (D/H)CH2 DOH .
18
O/
O ratios
arethedependent
the galactocentric
distance
andavN30 (van der Walt et al. 2021) are included in Figs. 3 and4. Onlyamong
high-masson
sources.
Excluding upper
limits, thedeterminedthe3 OHrelations
of±Milamandsources where NCH3 OH is derived from the 13 C or 18areerageusing(D/H)CHO isotoporatio is 1.1
0.7 ⇥ 10et3 . al.
This(2005)
is almost
two The derived D/H ratios are also listed in Table B.1.
& Roodrespectively.
where
only
The resulting (D/H)CH3 OH and (D/H)CH2 DOH ratios are prelogues are included in Fig. 3 to ensure that NCH3 OH isWilson
orders(1994),
of magnitude
higher thanInthecases
D/H ratioin thelocalupper
ISM ofnot under2
13
18
10
et al.of2006;
et al.
2010),
2 ⇥ 10 5 (Linsky
estimated. The (D/H)CH3 OH ratios lie mostly in the 10
limitson4 the⇠ column
densities
bothProdanović
sented in Figs. 3 and 4, respectively. Including upper limits, a
CH3 OH
and
CHsuggest3 OH
ing e↵ective
deuterationin the coldwas
high-mass
prestellar
phases.
range. Interestingly, all the ALMAGAL sources andcould
Orion beKL derived,
the
range
in
N
calculatedbyset- (limit on the) (D/H)CH3 OH and (D/H)CH2 DOH ratios could be deCHOH
3
2
3
show higher (D/H)CH3 OH ratios (1010 ) than Sgr B2(N2) However, both the range of observed (D/H)CH3 OH values and thetingbetween
the 3 average
upper limit
based
onorder
scaling
the 3 upper
limit
ofis rived for 38 and 26 of the 99 studied sources, respectively. Beand NGC 6334I (10 3 10 4 ). No clear correlation
is
more
than
one
of
magnitude
lower
than
what
13
CH3 OH and the lower limit based on the main isotopologue. sides the ALMAGAL sources, also other classical high-mass hot
Article number, page 4 of 17

Article number, page 5 of 17

inheritanceof methanol
ice from the high-mass p
mass
andhigh-mass
systems.
sources, only ALMAGAL datapoints are shown sinceZamudiono other etthe al.
average(D/H)
forintermediate-mass
protostars
inbe2018;Ligterink
et
al.
2022,
the
D/H
ratio
of
CH
OH
3
show somewhat elevated(D/H)CH2 DOH , but these were derived As evident from Figs. 3 and 4, the methanol D/H ratio derived for3
theprotostellar
phase.
(8.9±7.4⇥10). Athesimilaris2 OHseen(Taquet
forvalues).
theethigh-massinterferometric studies of CHD2 OH in high-mass protostellar
HavinghigherD/Hratios forDOH
issignificantly
higherthanthethatdoubly
derived deuterated
for CH3 OH.isousingspectroscopic
data
oftrendCHD
al. 2019).
NGC 7192tween
FIR2
isolder
taken
from
beam
averaged
How- CHto2
low-mass
prestellar
phaseswhere
the average(D/H)
systems are available. The derived (D/H)CH2 DOH ratios lie mostly 4.and
topologuecomparedwithsinglydeuterated
isotopologue
is no
CHlowDiscussion
3 OH
The
average
(D/H)
ratio
is
about
two
orders
of
magIn
Fig.
5,
the
mean
(D/H)
ratio
is
presented
for
CH
DOH
CH
OH
2
3
ever, interestingly
the±(D/H)
ratio
seems
Fig. 5.
The(D/H)
d
2 very similar beCH3 2andDOH3.4
CHratios
OH ratio
1 average
3
are
5.9
5.1
⇥
10
±
1.9
⇥
10
,
respectively.
Theuniquetomethanol.
For
water,
the
D
O/HDO
are
on
in the 0.1 1 range, with an average of 2.0 ± 0.8 ⇥ 10 1 , which is ratios
2 CH3 OH . In contrast
mass, intermediate-mass, and high-mass protostellar systems. nitude higher (2.0 ± 0.8 ⇥ 10 ) than (D/H)
2
forlow-mass,
intermediate-mass,
and
h
tween
low-mass
protostars
andhigh-mass
protostars
(0.1
1,
see
lower
(D/H)
ratio
in
both
high-mass
hot
cores
and
hightheorderof
10(e.g.,
Coutens
et
al.
2014;
Jensenetalmore than two orders of magnitude higher than the (D/H)
to
(D/H)
,
this
is
in
good
agreement
withtheaverage
of
CH
OH
4.2.
Singlyvsdoublydeuterated
methanolThe
average
(D/H)
ratio
for
high-mass
hot
cores
(1.1
±
CH3 OH 4.1.
CH3 OH
Methanol3 deuterationfrom low to high mass
CH3 OH
1 and
3
low-mass
prestellar
cores
andsucceshigh-m
mass0.7prestellar
phasescompared
toIRAS
theirlower-masscounterwhich
about
an orderof magnitude
higherthan
typi
Fig.single
4). Amongsources,
2A andIRAS
4A 2021),3.0 ± 2.0
⇥ 10 isforlow-mass
protostars,
suggestingthat
ratio. Furthermore, this indicates that about 1/5 of the⇥the
10 low-mass) liesmore than
one orderof magnitude
lower
than the3
2 highM. L.
GelderMethanolpartssuggestsalower
deuteration
efficiency
already
inthesive
deuteration
happens
almostequallye↵ectivein al.:
both
lowcalHDO/HO
ratios(.
10;3Persson
et
al.
2014;
Jensen
etand
alrd
lar
phasesareindicatedwithetcirclesaverageratio
forlow-mass
hotcorinos
(2.2
±
1.2
⇥
10
),
with
deuterated methanol molecules gets successively deuterated
fur2
As
evident
from
Figs.
and
4,van
the
methanolD/H
these
derivedshow somewhat
elevateda (D/H)
In this work,(limit CH
on2 DOHthe), but
(D/H)
could be deCH3were
OH ratiomass andhigh-mass
systems.
masstheprestellar
phasesCH(see
Sect.
4.3).
van’t
Ho↵is
etsignificantlyal. 2022). Thishigherdi↵erence
wasthatattributed
toaverage (D/H)
for intermediate-mass
protostars in be- 2019;
ther toward CH2 DOH in high-mass protostellar systems.
Thissquares.
3 OH
CH
DOH
than
deriveusing older
spectroscopic
data
of
CHD
OH
(Taquetetsources.
al.
2019).Since
rived
formethanol
38 of thestudiedhigh-mass
large
2
399
2trendtween(8.9±7.4⇥10).
A
similar
is
seen
for
the
high-mass
HavingD/Hratiosfor thedoublydeuterated
iso-FuThe
D/H
ratiosderivedinlow-mass
prestellarbetheresultof layeredice(⇠chemistry(Dartoiset2018).
al. 2003;
rived
fromhigher
HDCO/H
3%;
Perssonet al.
Small
is in good agreement with low-mass protostellar systems where
2 CO
The
(D/H)
is
about
o
In Fig.
5, and
the
mean
(D/H)
ratio
isthepresented
forCHlowsamples
of wellboth
low-mass
andhigh-mass
protostellar
systemsCH
low-mass
prestellar
whereforaverage (D/H)
CHphases
topologuecompared
withthesinglydeuterated
isotopologue
istwo
not
cores
agree
with
those
derivedlow-mass
protostars
ruya
et al.average2016),
whereof theratio
water
icecanis formed
invariationsin temperature
inbulkthe2 DOH
10–20
K rangechangethethe3 OH
3 OH
about 1/4 of the CH2 DOH is successively deuterated toward1
3
2
ratios
are
5.9
±
5.1
⇥
10
and
3.4
±
1.9
⇥
10
,
respectively.
Theuniquetomethanol.
For
water,theD
O/HDO
ratiosare 4.3),on) than(D/H)
nitudehigher
(2.0
±
0.8
⇥
10
D/H
ratio
of
ice
mantlespeciessuchasmethanol
(see
Sect.
mass,
intermediate-mass,
andhigh-mass
protostellar
systems.
(see methanolFig. 3). Since
CH3values
warmer translucent cloud phase with a low
D/H
ratio
whereastheOH is formed
through
the hydrogena2
CH
with
D/H
are now
available,a more significant
CHD2 OH (Drozdovskaya et al. 2022).
3
2
lower
(D/H)
ratio
in bothhot cores
and(1.1
the
order
of 10
(e.g.,
Coutens
etsinceal. CHD
2014;
Jensenetwithal.
largeerrorbar
it
based
on
but
this
should
a↵ect
both
CH
and
OH
in
a similarCH
surfacelayersformedinthecoldprestellar
phasesshowhighetion(D/H)of CO
ice(e.g.,
Watanabe&high-mass
Kouchican
2002;
Fuchs
ethighal. ±
3 OH
2 DOH
2is
to
(D/H)
,
this
is
in
good
agreement
The average
ratio
for
high-mass
hot
cores
CH
OH
comparison
over
the
mass
regime
be
made.
In
Fig.
3,
also
CH
OH
3
3
mass
prestellar
phasescomparedtotheirlower-masscounter2021),
which
is
about
an orderofismagnitude
higherthan
typi- (Dway
andshouldtherefore
not leadtothe observed
di↵erence.
1(G034-G2(MM2))
D/H
ratios.
However,methanolthought
to
only
start
forming
2009;
Simonsetal.
2020;
Santosetal.
2022),
thisisastrong
3
one
has
a
high
3.0
± 2.0
⇥ratios
10explanationfor
3low-mass protostars, suggesti
0.7 ⇥ 10 (D/H))indication
lies
more
than
one
order
of
magnitude
lower
than
the
ratiosderivedforboth
low-mass
prestellar
cores
parts
suggestsalower
deuteration
efficiency
alreadyinthehighcal
HDO/H
O
(.
10
; Persson
et
al.
2014;
JensenetWatanal.ofOnepossible
could
be
the
optical
depth
CH
OH
2
3
in
the
cold
prestellar
phaseswhereCOisfrozenout(e.g.,
forinheritanceofmethanol
and
other
COMs
ices
be4. Discussion
2/N 3 OH 2019;
Fig.mass
5. Theaveragephases(D/H)
ratio4.3).
derivedfrom
NCH(Fontani
etThisal.2009)2015).
The
other
two
CH
sive
deuteration
happens
almostequallye↵ectiv
prestellar
(see
Sect.
van’t
Ho↵
etlow-mass
al.Fuchs2022).
di↵erence
was
to
3 OH
2 DOH ), CH
average ratioforlow-mass
hotcorinos
(2.2
±
1.2
⇥
10
with abeCH
In thesourceL1551
IRS5,
theattributed
emissionof
2 DOH.
andhigh-mass
starlesscores
areincluded.
Itisevident
that
the&Kouchi
2002;
et
al.
with
little
to
no
forma
tween
low-mass
prestellar
phasesandprotostars.
On
the
other
for low-mass,intermediate-mass,
andhigh-mass
protostellar
systems be
13ice chemistry (Dartois et al. 2003; FuThemethanolD/Hratiosderivedinlow-mass
prestellar
theresultoflayered
CH
DOH
(as
well
as
CH0548-mm3;OH)
was suggested
to befor
opticallyEC
andFontani
et
al.
20
4
2
4.1. Methanol deuteration from low to high mass the average2
3systems.
mass
andhigh-mass(D/H)
forintermediate-mass
protostars
inbetioninthewarmertranslucent
phases.
Indeed,
also
a
direc
ratio
for
high-mass
starless
cores
hand,
the
average
(D/H)
CH
OH
CH
OH
ratio
is
lower
in
high-mass
hot
cores
(10
10
)
(D/H)
3
3
and cores
low-mass
prestellar
cores
and
high-mass
starless
cores.
The
prestelCH3agree
OH3 wellwith those derived for low-mass protostars ruyaet (Bianchi
al. 2016),et
where
the
bulk
of 2⇥10
the wateris(D/H)
formed
in
3ice
thick
al. 2020).
However,
since
the
ra3
CH
OHthe
3CO,
show
D/H
.
which
agree
we
precursor
of
CH
(5.9
±
5.1
⇥
10
)
seems
to
be
about
a
factoroffivehigherOH
onthesurfaceof
dust
grains,Hthed
2
1
3 higher D/H ratios for the doubly2
tween(8.9±7.4⇥10).low-mass
A similar
trend
is
seen
for
the
high-mass
Havinglar phasesare3).
indicatedwithcirclesandtheprotostellar
phaseswithIn this work, a (limit on the) (D/H)CH3 OH ratio could
be dethanin
theircounterpart
(10
10
).
Intermediate(see
Fig.
Since
CH
warmertranslucent
cloud
phase
with
a
low
D/H
ratio
whereastheOHisformedthroughthehydrogenatios
derivedfromCH
DOH
clearlyshow
lower
values
in
high3
3
2
than
that
for
high-mass
protostars
(1.1
±
0.7
⇥
10
D
).
However,
CO/HDCO
ratio
in
IRAS
16293-2422
pointstowardaisotohigh2squares.
GAL
sources
as
well
as
with
Orion
3 CH
2
and large
low-mass
prestellar
phaseswhere
theaverage(D/H)
topologuecomparedwithsinglydeuterated
rived for 38 of the 99 studied high-mass sources. Since
surfacelayersformedinthecoldprestellar
phasesshow
higher
tion
of
CO
ice
(e.g.,
Watanabe&Kouchi
2002;
Fuchs
et
al.
protostellar
systems compared to their lower-mass coun-K
3 OH ),mass
mass
protostars
show
valuesinbetween
(10
10
but
this
average(D/H)
ratio(seeofHowever,⇠ 25% methanol
compareda much
lower
D/H
ratio
de
of thehigh-mass
starless
has a D/H
3etCHal.
3 OH
D/Hratios.
istothought
to only
forming
Simons2020;
Santoset al. 22022),
thiscores
is a strong
terpartSect. 4.1),
this does
not seemlike astart
viablesolusamples of both low-mass and high-mass protostellar systems the 2009;

B2(N2)For
andwater,most13 ofinratios aresubsample
5.9 ± 5.1 ⇥ 10
and3.4 ± 1.9
⇥three10 ,sources
respectively.
Theuniqueto Sgrmethanol.
thetheD2cores
O/HDO
only
consists
of
(NGC
7192
FIR2,
the coldphaseswhere COis frozen
out2(e.g.,indication for inheritance of methanol and other COMs ices be- intion.
4 WatanVeryprestellar
recently,
spectroscopic
data
fortheCH
DOH
has be-T
with methanol D/H values are now available, a more significant
2
lower
D/H
ratiosat
10level.
Articlenumber,
page
6
of
17
lower (D/H)
ratio
in
both
high-mass
hot
cores
and
hightheorderof
10(e.g.,
Coutens
et
al.
2014;
CH
OH low-mass
large
since
it
is
based
on
three
detectionsofwhich
3errorbar
Cep
E-A,
and
Serpens
SMM1-a,Fuenteetal.
2014;
Ospinaabe&
Kouchi
2002;
Fuchs
et
al.
2009)
with
littletonoformatween
prestellar
phasesandprotostars.
On
the
other
come
available
(Ohno
et
al.
2022),
but
these
do
not
yet
includetheCHmassregimecan from
be made.
also
Fig. 5. comparison
The averageover
(D/H)
ratio
derivedNCH2In/NCH3,3 OH
M. L. van Gelder et al.: Methanol deu
DOHFig.
2
3 OH
inheritanceofmethanol
ice
from
the
mass
prestellar
phasescomparedtotheirlower-masscounter2021),
which
is
about
an
order
of
magnitudehi
one(G034-G2(MM2))
has
a
high
(D/H)
ratio
of
⇠
10
tion
in
the
warmertranslucent
phases.
Indeed,also
foradirectratio
forhigh-mass
starless
cores
hand,
the
average(D/H)
acalculationofthepartitionfunctionandline
properties
such
CH
OH
CH
OH
3
(D/H)
ratiosderivedforboth
low-mass
prestellar
cores
Zamudio
et
al.
2018;
Ligterink
et
al.
2022,
the
D/H
ratio
of
for low-mass, CHintermediate-mass,
andhigh-mass
protostellar
systems
OH
3
3 The other two detections (AFGL 51423size(Fontanietal.
2015).
precursor
of
CH
(5.9
±
5.1
⇥
10
)
seems
to
be
about
a
factoroffivehigherOH
onthesurface
of
dust
H
CO,
the
.
When
assumingthatthesourceisgrains,
equal
to
beam
as
A
3
2the
parts
suggestsalower
deuteration
efficiency
alreadyinthehightotheprotostellar
phase.
ij
cal
HDO/H
O
ratios(.
10;
Persson
et
al.
201a
2
and high-masscores
are included.
is evident
that the NGC 7192 FIR2 is taken from the beam averaged
and low-mass
prestellar starlesscores andhigh-mass
starless It
cores.
The prestelrived
from
HDCO/HCO(⇠
3%;
Perssonetal.
2018).
Small2
values).
How3that for(seehigh-mass
protostars(1.1 ±and0.7all⇥ 10). However,
CO/HDCOratio
in depth
IRAS
16293-2422
pointscan
towarda highthe line
opticalof
themostKconstraining
transition,
EC thanand0548-mm3;
Fontaniet al. 2015)
the upper
limits Dsize,
422variations
intemperature
in
the
10–20
range
changethe wC
mass
prestellar
phases
Sect.
4.3).
ratio
is
lower
in
high-mass
hot
cores
(10
10
)
(D/H)
2019;
van’t
Ho↵
et
al.
2022).
Thisdi↵erencelar phases
areCH
indicatedwithcirclesandtheprotostellar
phases
with
2
3 OH
3
ever,
interestingly
the
(D/H)
ratio
seems
very
similarbetheD/H
average(D/H)CH
ratio
⇠ice
25%
toK)a ismuch
lower
D/HSect.
ratio
de- t
the CH
high-mass
starless
has a D/H
(17
e0 ofof17
e0 ,compared
Eupspecies= 336such⌧methanol< 10. (seeOnlyforsourceOH ofshow. 2⇥10
which
agree
well
with mostof coresthe ALMA1,16
0,17
2 DOH
D/H
ratio
mantleas
4.3),
squares.than in their low-mass counterpart (10 2 10 1 ). Intermediate00
The tween
methanolD/Hratiosderivedinlow-mass
prestellarbetheresultoflayeredice
chemistry(Dartoisesizes
smaller
than
<
0.5
does
CH
DOH
becomemarginally
op2
GAL
sources
as
well
as
with
Orion
KL
(Neilletal.
2013).
Only
3
2
but
this
should
a↵ect
both
CH
b
DOH
and
CHD
OH
in
a
similarlow-mass
protostars
andhigh-mass
protostars
(0.1
1,
see
2 vs doubly 2deuterated memass protostars show values in between (1010 cores), but this
4.2.
SinglyArticle withnumber,thosepage 6 ofderived
17
ticallythick>therefore
0.1) fortheleadmostlinerichofsources.
Also,icethe li
agree
well
for
low-mass
protostars
ruyaandetshouldal.(⌧2016),
where
the
bulk
the
water
way
not
to
the
observeddi↵erence.
Sgr
B2(N2)
and
most
of
the
cores
in
NGC
6334I
show
slightly
sources, IRAS 2A and IRAS
4A
subsample only consists of three sources (NGC 7192 FIR2, Fig. 4). Among the low-mass(D/H)
ratios explanation
where Ncloud
wasderivedfromthepossibly
4
CH
OH
CH OH
One
possible
could
be
the
opticaldepthof rarlower
D/Hratiosatthe10level.
This
thereforealsosuggests
(see
Fig.
3).
Since
CH
warmertranslucent
phase
with
a
low
D/H
OH
is
formedthroughthehydrogena3
13
Cep
E-A,
and
Serpens
SMM1-a,Fuenteetal.
2014;
Ospinalarge errorbar since it is based on three detections of which showAsevident
from
3 and4, thethick
Clow-massisotopologue
areL1551
onFigs.
averagelessemission
thana factor
but these
wereopticallyderivedsomewhatelevated(D/H)
eCH
In thesource
IRS5,
the
ofmet
CH2 DOH , prestellar
2 DOH.
inheritanceofmethanol
ice
from
the
high-massphase
surfacelayersformedinthecoldprestellar
phas
tion
of
CO
ice
(e.g.,
Watanabe&Kouchi
2002;
Fuchs
et
al.
2
13
Zamudio et al. 2018;
et al.
2022, the D/H ratio of⇠CH
32019).
than
theas(D/H)
ratioswhere NtoCHbe
couldthanbeone (G034-G2(MM2))
has aLigterinkhigh (D/H)
DOH (as
well
CHCH
was
suggested
OH significantly
OHoptically2higher
3 OH)
CH3 OH ratio of ⇠ 10CHishigherusing
older
spectroscopic
data
of
CHD
OH
(Taquet
et
al.
2 DOH
to
the
protostellar
phase.
2
18
NGC
7192
FIR2 istaken
fromtwothe detections
beam averaged
values).
Howusing theoptically
O isotopologue(seeCHFig.
E.1).
D/H(Bianchi
ratios.
methanolis thought
2009;
Simons et al. 2020; Santos et al. 2022), this is a strong derived
thick
etHowever,al. 2020).thin
However,since the(D/H)
ra-onli
OHto(Fontani
et al.
2015).
The
other
(AFGL
5142Theaverage
(D/H)
ratio
is T
In
Fig.
5,
the
mean
(D/H)
ratio
is
presented
for
lowCH2 DOH
CH
OH
A
more
realistic
explanationisthatsuccessive
deuterationtios
derivedfromCH
DOH
clearlyshow
lower
valuesinhighever,interestingly
the
(D/H)
ratio
seems
very
similar
be3
2
CH
DOH
2
o
for inheritance of methanol and other COMs ices be- in the cold prestellar phases where CO is frozen
EC and 0548-mm3; Fontani et al. 2015) and all the upperindication
limits1
mass
protostellar
systemscomparedtotheirlower-mass
counc
ofmolecules is nitude
more e↵ective
than (2.0
the first
deuteration.
This) extha
higher±0.8⇥
10intermediate-mass,
andhigh-mass
protostellar
systems.
tween
protostars
andwellhigh-mass
protostars
1, see mass,
4.2.
Singlyvsdoublydeuterated
methanolshow D/H. low-mass
2⇥10 3 which
agreewith mostof the (0.1
ALMAabe&Kouchi
2002;
Fuchs
et
al.
2009)
with
littl
tween low-mass prestellar phases and protostars. On the other planation
terpart (seeSect. 4.1),bythis
doeslaboratory
not seem studies
like a viablesolu- at p
is supported
severalperformed
Fig. 4). Among the low-mass sources, IRAS 2A and IRAS 4A The average (D/H)CH OH ratio for high-mass hot cores
to
(D/H)
,
this
is
in
good
agrea4(1.1±
13
CH3 OH3
GAL sources
as
well
as
with
Orion
KL
(Neilletal.
2013).
Only
tion.
Very
recently,
spectroscopic
data
for
CH
DOH
has
betioninthewarmertranslucent
phases.
Indeed,ratio
forhigh-mass
starless
cores
hand,
the
average(D/H)
of 10–20 K (e.g.,1Nagaoka et2al. 2005, 2007;
CH3 OH
3 and 4, the methanol D/H ratio derived for low temperatures
show somewhat elevated (D/H)CH2 DOH , but these were derived As evident3 from Figs.
3.0±Drozdovskaya
2.0
10 theforlow-mass
protostcome available
(Ohno
et
al. ⇥
2022),
butet
these
do not yet
include
0.7
⇥
10
than
the
3 ) lies more than one order of magnitude lowerHidakat
Sgr B2(N2)
and
most
of
the
cores
in
NGC
6334I
show
slightlyetal.
2009).
al.
(2022)
showedthat
precursor
of
CH
(5.9
±
5.1
⇥
10
)
seems
to
be
about
a
factoroffivehigherOHonsurface
of
dust
grai
CH
DOH
is
significantly
higherthanthatderivedforCHOH.
using older spectroscopic
data
of
CHD
OH
(Taquetetal.
2019).
3
2
3
2
4
a2 ),
calculationof(D/H)the partition
functionand line
properties
suchequtheir
observedratio
asderived
from
CH
DOH
for s
sive
deuteration
happens
almostlower D/HInratios
at
the
10
level.
This
thereforealsosuggests
3
CHOH
2
averageratio
forlow-mass
hotcorinos
(2.2
±
1.2
⇥
10
with
The
average
(D/H)
ratio
is
about
two
ordersofmagFig. 5, the mean (D/H)CH3 OH ratio is presentedthan
for lowthat for high-mass protostars
D
ratio
in
IRAS
16293-2422
points
CH2 DOH (1.1 ± 0.7 ⇥ 10 ). However,assumingthat 16293-2422
thesourcesize
is equal
to thebeam
as
ACO/HDCO
2T
ij . When binary
thelow-mass
IRAS
could
be
well
explained1
inheritancemethanol ice fromand
thehigh-mass
high-massprotostellar
prestellar phasemass
andhigh-mass
systems.
than (D/H). In contrast
nitudehigher(D/H)(2.0 of± CH
0.8⇥high-mass
10for) intermediate-masstheaverage
protostars
inbemass,ofintermediate-mass,
systems.
size,
the
line
optical
depth
of
the
most
constraining
transition,
CH3 OHcores
D
OH
3
the
average
(D/H)
D/H
ratio
of
⇠
25%
compared
to
a
much
lower
the
starless
has
a
OH
and
CD
OH
were
overbythese
experiments
whereasCHD
CH3 OH2
3
2
3 good agreement with the average ofto the protostellar
to (D/H)
(17
e
17
e
,
E
=
336
K)
is
⌧
<
10
. Onlyfor result
source⇠
The averagephase.
(D/H)CH3 OH ratio for high-mass hot cores (1.1 ± tween
1,16
0
0,17
0
up
CH3 OH , this is in(8.9±7.4⇥10).
A
similar
trend
is
seen
for
the
high-mass
HavinghigherD/Hratiosfortproducedbytheexperiments.
The
latter
could
be
the
of
00
sizes smaller than < 0.5 does CH2 DOH become marginally op- w
3.0 ± 2.0 ⇥ 10 1 for low-mass protostars, suggesting that succes- the0.7 ⇥ 10 3 ) lies more than one order of magnitude lower than the andhighatomictopologue
D/H flux of 0.1
used in the laboratory
studiesdeu
inlow-mass
prestellar
phaseswhere
theaverage(D/H)
Articlenumber,
page
6
of
17
comparedwithsinglyCHOH2tically
thick
(⌧
>
0.1)
for
the
most
line
rich
sources.
Also,
the
3
⇠
5
equallye↵ectivein
both
lowaverage ratio for low-mass hot corinos (2.2 ± 1.2 ⇥ 10 ), with sive deuteration happens almost
contrast
to
the
ISM
value
of
⇠
10
,
althoughtheatomicD/H
ra3
2(D/H)The
where Nto
was derived fromthepossiblythedratios
are
5.9
±
5.1
⇥
10
and
3.4
±
1.9
⇥
10
, respectively.
4.2. Singly
vsdoublydeuterated
methanol
CH OH ratios
CH OH
uniquemethanol.
water,
the average (D/H)CH3 OH for intermediate-mass protostars in be- mass and high-mass systems.
tio is enhanced
in cold dense
prestellar
cores. For
Assumingthat the
13
3

3

3

3

3

3

3

3

3

3

3

optically thick C isotopologue are on averageless than a factor2tCH3 OH are found ranging from 0.1–1 with an avt 4.3. (D/H)
Linking
the±methanol
1D/H to the physical conditions
erage
of
2.0
0.8
⇥
10
. The latter is good agreementrduringformation
withresults
on low-mass protostars and suggests that abouti
1/5 singly deuterated methanol molecules gets successivelys Giventhe sensitivityof the methanol
deuteration
process todeuterated
further
independent
of
the
mass
of
the
system.
ebothtemperatureanddensity(i.e.,CO
freeze-out),themeasured
–
Based
on
a
comparison
with
GRAINOBLE
modelsinthelitH
ratios(D/H)
are linked
to thesephysicalproperties
durtoward
high-massproerature,D/H
the lower
CH3 OH ratios- methanol
tostars
suggestphases.
either a To
temperature& 20
in protostellar
the highprestellar
quantifyofthis
forKthesys- ing the) of theare
highprestellar
phasesor thea shortlifetime D/H(. tFFratios
n temsmassstudiedin thiswork,observed
comparedmass
prestellar phases.
This is models
in strongpresented
contrast with- to theastrochemical
gas-grainby theBøgelundsources for which the higher (D/H)CH3 OH ratio cann et al.low-mass(2018)
andTaquetetal.
(2019).
These
works
used
the
only be achieved when the low-mass prestellar phases areo
GRAINOBLE
model(Taquet
et
al.
2012,
2013,
2014)
to
test
both
cold
(<
15
K)
and
long
lived
(
t
).
FF
s
g the e↵ect of the dust and gas temperature T (assumed to beThis workthatthe 2CH
A&Atheproofs:manuscript
no.of
ALMAGAL_D_H_methanol_paper
3 -groupand demonstrates
hydrogen density
nHdeuteration=n(H) +2n(H) on
the resulting
d equal)ofmethanolasmeasured
towardprotostellar
systemscouldbet (D/H)CH3 OH ratio in the ices during the prestellar phases. In thisused to probe the physical conditions (e.g., temperature) of thework,
we phases.
compareresults toin theirand therefore
only
prestellar
Theourdiscrepancy(D/H)results
between
lowCH3 OH
amass
briefand
description
the model
is presented.
high-massof
sources
indicatesthat the physical condiInareGRAINOBLE,
gas-icechemistry
is computedtions
already di↵erentthe
beforethe onset
of star formation.
Ad- in threeditional observationsof multiple
isotopophases:
the bulk ice,
the icedeuterated
surface methanol
layers, andin the gas
d phase,
CD3 OH)initiallyas well as
CH3 OD will
shed
logues following(e.g., CHDthe
2 OH,
approachpresented
by Hasegawa&
further (1993).
light on theefficiencyofmethanoldeuteration
between, Herbst
The model includes both adsorption and desorpf low-mass and high-mass systems. In combination with addition
reactionsandcomputestherateequationsineachphase.
The- tional modeling studies (such as those performed by Bøgelundchemical network used for the gas-phase chemistry is describedin Taquet et al. (2014) and includes both ion-neutral chemistry

smaller than the free-fall timescale are considered. For the lowest density of 104 cm 3 , this e↵ect is most evident since the CO
freeze-out timescale is the highest and therefore the methanol
deuteration is hampered the most.
Overplotted in Fig. 6 are the observed methanol (D/H)CH3 OH
ratios for both low-mass and high-mass protostars. It is evident
that the observed (D/H)CH3 OH ratios suggest a di↵erent temperature during methanol formation or di↵erent prestellar phase
lifetimes for low-mass and high-mass protostars. For high-mass
protostars, a temperature of > 20 K is needed when the density is larger than 105 cm 3 and the timescale of the high-mass
prestellar phase is > tFF . A lower temperature of > 13 K at low
densities of 104 cm 3 can also explain the observed methanol
D/H ratio toward high-mass protostars, but such low densities
in the dense high-mass starless phase are unlikely. Alternatively,
the temperature in the high-mass prestellar phases can be in the
15 < T < 20 K range with a lifetime of . tFF . For prestellar lifetimes much smaller than the free-fall timescale, any temperature
can explain the observed methanol deuteration toward high-mass
protostars.
Contrary to the high-mass protostars, the observed(D/H)CH3 OH ratio for low-mass protostars suggests both atemperature of < 15 K and a prestellar phase duration longer

Fig. 6. The (D/H)CH3 OH ratio in the ices as function of the gas and dust temperature for nH = 104 (dashed), 105 (dotted), and 106 cm 3 (solid) as
modeled by Bøgelund et al. (2018) and Taquet et al. (2019) using the GRAINOBLE model (Taquet et al. 2012, 2013, 2014). The model results
are shown for 0.1 (left), 1 (middle), and 10 (right) times free-fall timescales tFF . The observed average (D/H)CH3 OH ratios are indicated in blue
and orange for the low-mass and high-mass protostars, respectively. In the right panel, the range of possible desorption temperature of CO ice isindicated with the gray bars for binding energies between 863 K and 1307 K for CO ice deposited on non-porous amorphous solid water (Nobleet al. 2012).

Article number, page 7 of 17

ph.SR] 9 Aug 2022

Monitoring accretion rate variability in the Orion Nebula Cluster?
with the Wendelstein Wide Field Imager
1
1
1
2
1
S. Flaischlen , T. Preibisch , M. Kluge , C.F. Manara , and B. Ercolano
1

Universitäts-Sternwarte, Fakultät für Physik, Ludwig-Maximilians-Universität München, Scheinerstr. 1, 81679 München, Germany
2
European Southern Observatory, Karl-Schwarzschild-Str. 2, 85748 Garching bei München, Germany
Received November 9, 2021; accepted July 25, 2022
ABSTRACT
Context. The understanding of the accretion process has a central role in the understanding of star and planet formation.
Aims. We aim to test how accretion variability influences previous correlation analyses of the relation between X-ray activity andaccretion rates, which is important for understanding the evolution of circumstellar disks and disk photoevaporation.
Methods. We monitored accreting stars in the Orion Nebula Cluster from November 24, 2014, until February 17, 2019, for 42 epochswith the Wendelstein Wide Field Imager in the Sloan Digital Sky Survey u0 g0 r0 filters on the 2 m Fraunhofer Telescope on Mount
Wendelstein. Mass accretion rates were determined from the measured ultraviolet excess. The influence of the mass accretion ratevariability on the relation between X-ray luminosities and mass accretion rates was analyzed statistically.
Results. We find a typical interquartile range of ⇠ 0.3 dex for the mass accretion rate variability on timescales from weeks to ⇠ 2years. The variability has likely no significant influence on a correlation analysis of the X-ray luminosity and the mass accretion rate
observed at di↵erent times when the sample size is large enough.
Conclusions. The observed anticorrelation between the X-ray luminosity and the mass accretion rate predicted by models ofphotoevaporation-starved accretion is likely not due to a bias introduced by di↵erent observing times.
Key words. accretion, accretion disks – open clusters and associations: individual (Orion Nebula Cluster) – protoplanetary disks –stars: pre-main sequence – stars: statistics – X-rays: stars

10:00.0

2. Observation
data reduction• 著者らのPaperandI
A&A 648, A121 (2021)
https://doi.org/10.1051/0004-6361/202039746

Astronomy&
Astrophysics© ESO 2021

2.1. Target region

Testing the models of X-ray driven photoevaporation withaccreting stars in the Orion Nebula Cluster?

The ONC is among the nearest regions of ongoing star formation
T. Preibisch , C. F. Manara , and B. Ercolano(Bally 2008), with a distance of ⇠S. Flaischlen
403 , pc
(Kuhn et al. 2019).
Therefore, it serves as a benchmark for star formation and has
been studied in detail over the years. As such, stellar parameters for a large fraction of the young stellar population are available in the literature (Manara et al. 2012). Especially relevant
for this work are the spectral type determinations from Da Rio
et al. (2012) and Hillenbrand et al. (2013). Together with our
photometric observation detailed below, the spectral types allow
an estimate of the accretion rates.
1

1

2

1

-05:20:00.0

Universitäts-Sternwarte, Ludwig-Maximilians-Universität München, Scheinerstraße 1, 81679 München, Germany
e-mail: sflaisch@usm.lmu.de
2
European Southern Observatory, Karl-Schwarzschild-Straße 2, 85748 Garching bei München, Germany

Received 23 October 2020 / Accepted 1 February 2021

ABSTRACT

Context. Recent works highlight the importance of stellar X-rays on the evolution of the circumstellar disks of young stellar objects,

especially for disk photoevaporation.

Aims. A signature of this process may be seen in the so far tentatively observed dependence of stellar accretion rates on X-ray

luminosities. According to models of X-ray driven photoevaporation, stars with higher X-ray luminosities should show lower accretion
rates, on average, in a sample with similar masses and ages.
Methods. To this aim, we have analyzed X-ray properties of young stars in the Orion Nebula Cluster determined with Chandra during
the COUP observation as well as accretion data obtained from the photometric catalog of the HST Treasury Program. With these
data, we have performed a statistical analysis of the relation between X-ray activity and accretion rates using partial linear regression
analysis.
Results. The initial anticorrelation found with a sample of 332 young stars is considerably weaker compared to previous studies.
However, excluding flaring activity or limiting the X-ray luminosity to the soft band (0.5 2.0 keV) leads to a stronger anticorrelation,
which is statistically more significant. Furthermore, we have found a weak positive correlation between the higher component of theplasma temperature gained in the X-ray spectral fitting and the accretion rates, indicating that the hardness of the X-ray spectra may
influence the accretion process.
Conclusions. There is evidence for a weak anticorrelation, as predicted by theoretical models, suggesting that X-ray photoevaporation
modulates the accretion rate through the inner disk at late stages of disk evolution, leading to a phase of photoevaporation-starved
accretion.

30:00.0

Key words. protoplanetary disks – stars: pre-main sequence – planetary systems – stars: statistics – X-rays: stars

2012; Alcalá et al. 2014, 2017). However, the observed spread
1. Introductionof up to two orders of magnitude in accretion rates for stars with
The interaction
betweenthe
youngstellarof
objects(YSOs)
andphotoevaporation
their
S. Flaischlen
et al.: TestingmodelsX-ray
drivenaccretingin theOrion
Nebulasimilar masseswith
dilutesthe Ṁacc stars
M relation(e.g.,
Hartmann
et al. Cluster
circumstellar disks is of fundamental importance for an under2016; Manara et al. 2017; Alcalá et al. 2017). Since the disk matestanding of planet formation and migration (e.g., Ercolano &rial gets dispersed over time, the accretion rate also decreases
Pascucci 2017). During their evolution, many YSOs are sur32-6
with isochronalage (e.g., Hartmann et al. 1998). However, there
rounded by circumstellar disks, from which material is accreted
is also evidence for stars with high accretion rates at old ages
onto the protostar. This process is thought to be controlled by(e.g., Venuti et al. 2019).
stellar magnetic field lines connecting the disk with the proThe dispersal of the disk material is caused by strong disk
tostellar surface, leading to a complicated system of magnetic
winds. Observational evidence and simulations suggest that irra-7“accretion funnels” that channel the accretion flow (Hartmanndiation31
from the central star is an important driving mechanism
et al. 2016). The infalling material generates shocks at the proof these disk winds (Alexander et al. 2014). In particular, models
tostellar photosphere, giving rise to strong optical emission lines
predict X-ray emission to be very efficient in photo-evaporating(in particular H↵) and optical and ultraviolet excess emission,
the circumstellar material (Ercolano et al. 2008a,b, 2009; Owen
-8
from which the accretion luminosity Lacc can be inferred withet al. 2010, 2011, 2012; Picogna et al. 2019; Wölfer et al. 2019),
spectroscopy (e.g., Alcalá et al. 2017) or photometry (e.g., Venuti
30the accretion rate (e.g., Drake et al. 2009) as well asaffectinget al. 2014). Together with the stellar parameters, the accretion
planet formation and migration (e.g., Ercolano & Rosotti 2015;
rate Ṁacc can be obtained from Lacc via Ṁacc = Lacc R/(0.8 G M),
Ercolano & Pascucci 2017; Jennings et al. 2018; Monsch et al.
-9where R and M are the stellar radius and mass, respectively,2019). Compared to main sequence stars, YSOs show highly elewhile the factor 0.8 originates from the assumption that the infall
vated X-ray activity through all evolutionary stages (Feigelson
stems from the magnetically truncated disk with a typical radius
29& Montmerle
1999; Preibisch et al. 2005; Preibisch & Feigelsonof ⇠5 stellar radii (Hartmann et al. 2016).
2005). Observations (e.g., Preibisch et al. 2005; Telleschi et al.
-10
Several observational studies support the expectation that the2007) have shown that the X-ray luminosity of YSOs increases
accretion rate is higher for more massive YSOs (e.g., Muzerollewith increasing mass, analogously to the accretion rates.
et al. 2003; Natta et al. 2006; Telleschi et al. 2007; Manara et al.
The28Chandra Orion Ultradeep Project (COUP) producedthelargest available data set of X-ray emitting YSOs in a star?
-11
The data underlying Figs. 5log(Mand A.1 [M
are only]) available at the CDS forming region (Getman et al. 2005b). The analysis
log(M of[Mthese€€])via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via-1.0 -0.8 -0.6 -0.4 -0.2 0.0 0.2-1.0 -0.8-0.6
data showed that, on average, accreting
YSOs
have-0.4
lower -0.2
X-ray 0.0 0.2
http://cdsarc.u-strasbg.fr/viz-bin/cat/J/A+A/648/A121

2.2. Data reduction and photometric calibration

30.0

05:35:00.0

34:30.0

RA [J2000]

Fig. 1. WWFI u0 g0 r0 color composition of the ONC (image credit:
Wendelstein Observatory). Overlaid are the stars monitored in all threebands. The yellow symbols indicate stars that have a match with the Xray sources listed in the COUP catalog (J/ApJS/160/319/coup, Getman
et al. (2005)). The star symbols further indicate sources that were monitored for at least five epochs and allowed an estimate for the mass accretion rate.

-1-1

36:00.0log(LX [erg s ])We observed the ONC in a multiyear and multicolor monitoring with the 2 m Fraunhofer Telescope at the Wendelstein Observatory located on Mount Wendelstein in Germany (Kosyra
et al. 2014). The telescope hosts the Wendelstein Wide Field Imager (WWFI), a wide field camera consisting of four 4096 px ⇥4109 px charge-coupled device (CCD) detectors arranged in a2 ⇥ 2 mosaic pattern. With 27.60 ⇥ 28.90 , the field of view covers the region of the X-ray data (170 ⇥ 170 ) used in our previous
analysis of the relation between X-ray luminosities and accretion27-12-1.0 based -0.5
0.0
-1.0 Orion-0.5
0.0
rates in F21, which
was
on
data
from
the
Chandralog(M [M ])log(M [M ])Ultradeep Project(COUP) (Getman et al. 2005).
Fig. 1. Correlation of the mass accretion rate and the X-ray luminosity with the stellar mass. Left: mass accretion rate vs. stellar mass for young
0 and0the masses are 0taken from Manara et al. (2012). The dashed line shows a single power-law fit for all objects
stars in the ONC. The accretion data
For the monitoring,theu, gindicates, andr result
filterswerewhichwith masses 2.0 M . The solid line
the fitting
for the mass
range used,
0.2  M  2.0,
and the shaded region is its 1- deviation fromthe data. The cross indicates the typical uncertainty of the values. Right: same but for the X-ray luminosities taken from Getman et al. (2005a). The
are approximately
comparable
fitting parameters
are listed in Table 2. with the photometric system oflog(dMacc/dt [M€ yr ])DEC [J2000]

1

Article published by EDP Sciences€

A121, page 1 of 11

€

,
(1)
The accretion rates can beh!i
estimatedfromthe mas,observedu , correspondsg,
R a mean distance of ues determined by Manara et al. (2012). The resulting color- gaussians accor
¯
=
2.48
±
0.06
which
to
0
with Colorthezero-pointZ⇣ andin Fig. 3 and illustrates how the colors2MASSJ11195652
3.2.
modeling
and rData
photometry
temperatures
known
from
+7
coloredaccordi2.3.
selectionand the⇠e↵ective
color
diagramisshown
403 pc, in agreement with the distance of 403 6 pc determined
ChaIeach
717,starandasV125
with
Iphot,
andwhere R⇤ is the radius of the source
Rthe intensity
the literature with an approach similar to the one used by Da Rio
soli
are
distributedwhenneglecting
extinction.
Furthermore,
Fig.
3
by
Kuhn
et
al.
(2019).
We
excludedthreesourcesaslikelybackUsing withsynthetic(see Casagrande & VandenBerg For
T ⇣,intensity
d photometry
Fthe
each
of theI.
et al.
(2010)and Manaraal. (2012),
whereonly
the lattercalled itof Aacc is the area of the accreting regions
typical
range
of
acc,
In
order
to ensurethat theet catalog
contains
photometry
shows
the
predicted
displacementofthesources
from
the
line
R
F̄
=
,
(3)
stars
to thethatfactor
their distancelarger⇣(2014)
for Ta detailed
description of the topic), the flux density
toFphot,for according
the “2CD”stars,
method.
(R⇤ /R )2 isscalesthe fluxes
accordingto the radius,
non-accretin
accretingthe sourcegroundpositionswereaccording
matched withthecriteria
cat- Thed⇣, photosphere according to Eq. 5. For a given Te↵
describing
the
than
the
one
reported
by
Kuhn
et
al.
(2019)
within
3
and
kept
The
methodassumesthatthepositionofthesourcesinthescaling
the distancecanFbemagnitudes m⇣ for the ⇣ filter
refer toto
an theareaexpectedsimsources
are disp
acc, related
alog of accreting stars from Manara et al. (2012), which are where it is assumed that Fphot, andvalue, a⇣trackin fortheucolor-color
0
0
0
0
0
0
0diagram is determined by vary2(V1118 Ori) and 2(wheretheonesourcewithoutalistedparallaxvaluegr
versusugcolor-color
diagramdepends
solely
on
three
rele
tion. EDR3
They conve
stands,g,andror anyofotherfilter) viafilter. est Gaia
describes
thefractionof the function
complete down to the hydrogen burning limit. This step reduced ilar to ⇡R , and ⌘ B Aacc /(⇡R⇤ )whereisthepassthroughtherespective
T
⇣,
ing
⌘.
The
track
starts
for
⌘
=
0
at
the
line
describingthepureparameters:
thee↵ectivetemperature,theaccretion
luminosity,
the distancethea O
assignedtoitthedistance
of
403
pc,
with
the
typicaluncertainty
describedofbyhcoveredby theaccording
accreting tore-Eq. 2 is given by
our sample size to 61. For 60 sources, we found parallaxes in projected area of the photosphereThemagnitude
photosphericemissionof the pureandGaia
the extinction.
colors fromwouldourfallsample.
onsourcewas scaled
accretion
to totatZ⇣ , and converges at the colors(2)
ofaccretion,⇡3 ±9
pc thededucedmgions.
the
Early DataWithout
Release(EDR3)
catalog(Gaia
Collabo⇣ = 2.5 log F̄ ⇣!
a line etdetermined
the photosphereof the stars
at a sources
positionis
accretion spectrum⇣ for increasing⌘⌘. The plotR⇤ also shows linesthe template source
rational. 2016,by
2020).
The mean parallaxof the(⌘) = 2.5ratio
log between
F̄phot, ⇣ + ⌘the· F̄accretion. page(4)Z⇣and
5Articlelogm⇣constantthe e↵ective
temperatureand scatter toalong
the reddening
acc, ⇣
ofluminosity
theet al. (2013, 2017a)!i¯ by= 2.48± 0.06 mas,which correspondsa meandistance of
hset
number,3totalof 12
R
3.2. Color modelingAccretion
shifts withthe position
towards
bluer+7 values to theluminosity, Lacc /Ltot . The calculation of these lines is detailed in With Fphot, , w⇠vector.
403 pc,in agreement
the distanceof 403pc
determined
6
colorsof eta purelyaccretingsource. three sources as likely backdomly displac
appendix
E. &colors
byKuhn
al. (2019).
We excluded
passthroughfunctioThe
respective
are independent of R⇤ :Using
synthetic
photometry(seeCasagrande
VandenBerg
Thestars
accretion
luminosity
the extinction
can be obtained
suminggroundaccordingto theandcriteriathat their distanceis larger (2014) for a detailed description of the
WWFI
setup Gauss
(Kosy
=
11000
K,
⌧
=
The
model
with
the
parametersT
topic),
theflux
densityslab
300!
simultaneously
from themagnitudes
and thee↵ val15 F̄ phot,3 ⇣ + ⌘ · F̄ acc, ⇣uncertainty
thanthe one reported
by observed
Kuhn et al.
(2019) within
3 Tandkept can be related to the expected magnitudes
to-noiseratio, theco
te
5.00,
and
n
=
10
cm
lead
to
tracksdescribing
the
m
for
the
⇣
filter⇣e
ues when the photospheric colors and the colors of the accretionm⇣ (⌘) mµ (⌘) = 2.5 log(5)Z ⇣ + Zµ .
the one source without a listed parallax value (V1118 Ori) and (where ⇣ stands for u0 , g0 , and r0 orof ourapproacany other filter) via colors
box-width
extinction-corrected
best.
magni-a typicalF̄phot,· F̄acc,correspondingµ + ⌘Theµsource are known. In the following, we describe our implementaassigned to it the distance of 403 pc, with the typical uncertaintystars
in our sam
to magnitude
tudes are macc, u0 = 10.70 mag, macc, g0 = 10.76 mag andfluxes
tion of the approach in more detail and explain how we obtained(2) weof⇡ ±9 pc deduced from our sample.
m⇣ = 2.5 log F̄⇣ Z⇣ ,In the0 following
subsections,howluminosity
we obtainedamounts
thethe time
Since
the span
relat
m
=
10.60
mag,
while
theexplain
accretion
the required photospheric colors.
acc, rand the accretion
flux
photosphericfluxesdescribed
by F̄thesephot, ⇣ parameters
monitoring.
and to
to
log(L
/L
)
=
1.04.
We
use
for
our
subse-complicatedacc,
0
Article number, page 3 of 12F̄acc, ⇣ .analysis.
cific shape, we uti
quent
The 1 un3.1. Flux density modelingin order to appro
sumed
to be
th
tudes
and
the
e↵ec
The flux density of the sources in question is assumed to be a
3.2.1. Photospheric flux
ity
and
the
ext
(2017a),
we
used
3.3.
Accretion
luminosities
andextinctions
fromobservations
composition of the pure photospheric flux density and the accrescribedabove.
sion
withaGaussiation flux density. Without extinction, the flux density of a source
The photospheric flux F̄phot,required
for
Eq.
5
was
deter0 0⇣ 0Weapplied
Eq.
5
to
each
u
g r observationand derived
and AVpyqt_fit.npr_me
tainty is define
located at a distance d can be expressed throughmined from empirically obtained
spectra of young,diskless⌘classusing
T e↵ , thephotosphericflux, andthe accretion
model.
Thetask.values.
thThe 2 With
confidIII
objects
servingasphotospheric
templates.
They
were
ob⇡R2⇤
Aacc
observed
positionin thebycolor-color
diagramis shiftedalong the Inand. stepservedwith
X-Shooter
Manaraetal.
(2013)
andManaratheLacc
final
F = Iphot, · 2 + Iacc, · 2
reddening
vector,it intersects
the ofuniquetrack calculatedddet al. (2017a).
We until
excluded
the spectra
the targets
named values, fromwhich
If therew!2 ⇣from T e↵ . This procedure provides AV as well as ⌘. The yet miss⌘R⇤sects the track(1)=· Fphot, + ⌘ · Facc, ,Articlenumber,
page
4
of
12
ing radius R⇤ can be calculated from the observed magnitudes,Rthe lower accr
00where R⇤ is the radius of the source with the intensity Iphot, andAacc is the area of the accreting regions with the intensity Iacc, .
The factor (R⇤ /R )2 scales the fluxes according to the radius,
where it is assumed that Fphot, and Facc, refer to an area similar to ⇡R 2 , and ⌘ B Aacc /(⇡R2⇤ ) describes the fraction of the
projected area of the photosphere covered by the accreting regions.

=· Fphot, + ⌘ · Facc,

AV , and ⌘ with Eq. 4. From the definition of ⌘, the accretion luminosity follows from!2
R⇤
· Lacc, 0 .
(6)
Lacc = ⌘ ·RIn the next subsection, we explain in more detail how we estimated the uncertainties of AV , R⇤ , and Lacc .

could be foun
the range of t
and the extinc
the results a coful intersection
iterations, welower confiden

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Fig. 3. Color-color diagram of the WWFI sample, dereddened with theextinction values determined by Manara et al. (2012) and plotted as 2D
gaussians according to their photometric uncertainties. The values arecolored according to their T e↵ value. Overlaid are the mean colors foreach star as solid dots. The cross in the upper left corner indicates their
typical range of variability. The thick line shows the photospheric colors
for non-accreting stars. The thin lines represent tracks along which thesources are displaced from the photospheric track for increasing accretion. They converge in a point representing the pure accretion spectrum,described by a hydrogen slab model. The dotted lines indicate constantaccretion to total luminosities ratios, Lacc /Ltot .

-2

3.7

3.6
log(Teff [K])3.5

Fig. 4. Hertzsprung-Russel diagram of the observed sources. Overlaid
are the PARSEC 1.2S isochrones and mass tracks from Bressan et al.
(2012). The dots represent the mean bolometric luminosity for eachstar as a function of T e↵ . The vertical bars indicate the minimum andmaximum range of L⇤ for each star. The cross at the bottom shows
the typical uncertainty. In our analysis, we only regarded sources with
0.2 M  M  2.0 M and 5.5  log(⌧[yr])  7.3 in order to exclude
outliers in the Hertzsprung-Russel diagram.

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for non-accreting stars. The thin lines represent tracks along which the↵12). The dots represent the mean bolometric luminosity for eachsources are displaced from the photospheric track for increasing accre- tion. They converge in a point representing the pure accretion spectrum,r as a function of T e↵ . The vertical bars indicate the minimum ande described by a hydrogen slab model. The dotted lines indicate constant
ximum
range of L⇤ for each star. The cross at the bottom showse accretion
to total luminosities ratios, Lacc /Ltot .

typical uncertainty. In our analysis, we only regarded sources with

V1493 Ori

V1555 Ori

V1475 Ori

V786 Ori

V1477 Ori

V2400 Ori

V1552 Ori

V2080 Ori

V2420 Ori

V1481 Ori

LN Ori
JW 117

V2057 Ori

V488 Ori

MZ Ori

V401 Ori

0.0

-0.5

-1.0

4

a, Fig. 3. Color-color diagram of the WWFI sample, dereddened with the-2- extinction
values determinedby Manara
et al. (2012) and3.5
plotted as 2D
3.7
3.6
The values are- gaussians according to their photometriclog(Teffuncertainties.
[K])s colored according to their T e↵ value. Overlaid are the mean colors for
star
as
solid
dots.
The
cross
in
the
upper
left
cornerindicatestheir
3. 4. each
Hertzsprung-Russel diagram of the observed sources. Overlaid
range of variability. The thick line shows the photospheric colors
ethetypical
PARSEC 1.2S isochrones and mass tracks from Bressan et al.

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Fig. 5. Boxplot showing the accretion rate variability of the targets.
The vertical solid lines indicate the range of the accretion rates and theheight of the boxes the interquartile range. The circles show the accretion rates for each target and epoch. The values are normalized to themedian values and color-coded according to their e↵ective temperature.

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V2400 Ori

V1552 Ori

V2080 Ori

V2420 Ori

V417 Ori
V1531 Ori

V2023 Ori

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JW 117

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log(M
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- TAobserved[days])sources. Overlaid Fig. 5. Boxplot showing the accretion rate
Fig. 4. Hertzsprung-Russel diagram
ofB thevariability
are the PARSEC 1.2S isochrones and mass tracks from Bressan et al. The vertical solid lines indicate the range of the accretion rates and theFig. 6. The
Di↵erencesof thethelogarithmicmass luminosity
accretion for
rates
for two
Logarithmic
mean range.
mass The
accretion
vs. accrelogarithm of the(2012).
dots represent
mean bolometriceachheight Fig.
of the7.boxes
the interquartile
circles rates
show theepochsmassesshown
as andsolid
dots.The
Thevalues
solid arebarsnormalized
indicate themaximum andfunction
theindicate
logarithmic
time di↵erencesstar
as a TfunctionofB Tasfor eachtargetepoch.
to theverticalofbars
the minimum
and tion rates
A and Te↵ . aTheminimum
accretion
rates forsource. temperature.
The dotted line shows
T B T A forrange
eachofmeasurement.
The Thevaluesareatplotted
as squares,
medianthe
valuesand color-coded
accordingto eachtheir e↵ective
maximumL⇤ for each star.
cross
the bottom
showswheretheuncertainty.
In ouranalysis, color
we only
regardedwith to its the result of the regression obtained with the Bayesian LINMIX_ERR
for typicaleach distinct
objecta di↵erent
waschosensources
according
0.2
M  temperature.
M  2.0 M and5.5  log(⌧[yr]) 7.3 inorder ofto excludee↵ective
Overlaid
are the meanvaluesthe logarith- method. The blue cross shows the typical uncertainty of the values, theoutliersin theHertzsprung-Russel
diagram.
to photosphericIn addition,
we limited the sample acred bar theemission.
typical variability
range.
mic mass
accretion
rate di↵erences
for di↵erent time bins as magenta
0.0

columns, together with their estimated 1 uncertainties.

cording to 0.2 M  M  2.0 M and 5.5  log(⌧[yr])  7.3 inlog(ádMacc/dtñ) - log(dMacc,fit/dt(M, t))

tween
LX and M
the linear
regressionrelation:of the residuals would
lar been
massassigned(aswasforshown
numerous
works,
e.g.,
Alcalá
al. light curves ofwindows
withiniteration
significant
flaring
activity
inet theacc ,found
LINMIX_ERR.
We
thefollowing
has
every
step.
Each
time,
we
storedby
Kellyno(2007)
that takesand
uncertainties
in both variables
intoslopeac- tween
aindicate
positivecorrelationin agreementwith theof 1.07
LX ±andandṀaccthe, thelinear regressionofthetheresiduals
would
correlation.
(2017)
previous
section)and
X-ray
luminosities
asdenoted “characthe
COUPsources.
The
corrected
valuesare1.5
thelinearcorrelation
coefficients
drawn
from
the
LINMIX_ERR
4.2.
Relation
betweenaccretion
rates
andstellar
mass
count
as
well
as
scatter.
Theofresultingrelation
reads
no(e.g.,correlation.
! intrinsic
0.22
found
in ouranalysis332 accreting
sources
in
F21.
Towe
minimizethee↵ectof X-rayvariability,
we triedto cor- ! indicate
well
Preibisch
et
al.
2005;
Telleschietal.
2007),
thevalues.
Additionally,
wecorrected
teristic”
X-ray
luminosities,
L
LX,
posterior
distributions,consistingof
12600
values.
From
theirh Ṁacc i !
char
X,
char
To
minimize
the
e↵ect
of
X-ray
variability,
we
tried
to
cor!
Slope:-0.37± 0.24rect
thehFig.
influenceofflaring
activity
from
the
X-ray
luminosities
logIn
.
=
0.00
±
0.11
+
(
0.22
±
0.39)
·
log
7,
the
regression
resultisdisplayedasadottedline.
Weperformedalinearregression
in
order
to
analyzetherelamust
becorrected
forthecommonmass
dependence
in
order
to
M
Ṁacc i
distribution,
wedetermined
the
mean
and
the
standard
deviation
rect
theinfluenceofflaring
activity
from
the
X-ray
luminosities
thevaluesforthemore
recentGaia
distances
listed
in
Table
L
(M)
(M)
Ṁ
X,
char
=
(
8.05
±
0.15)
+
(1.07
±
0.43)
·
log
,
(8)
log
acc
using
the
“characteristic
countrates”
determined
by
Wolk
et
al.
tion assume
betweenthe (logarithmic)
mean
and the(logavoid
spurious
correlations.
Thisisconclude
done bybyWolkobtainingthelinearr: -0.36 ± 0.23
1 thatWe
the agreement
ofmasses
the regression
resultsfromourtoobtain
r
=
0.74
±
0.25.
We
can
with
95
%
confiM mean
M
yr
using
the
“characteristic
countrates”
determined
et
al.
1.
We
assumedatypicaluncertainty
of
L
of
0.15
dex
de(2005)asthequiescentlevelsof
X-rayemissionoutside
the
time
(10)
X,
char
1.0
arithmic)sample
mass accretion
rates.
To thisaim,
we usedthe fully
M and theLXthetheMtime
relationsisandregression
resultofrtheṀX-rayacc meaning
dence
that
1.00


0.24,
that
probability
smallerwiththeresultsofthelargersampleusedinF21(2005)
asthequiescentlevelsofemissionoutsidewindows
withcorrelation
significant
flaring activity
inthe
the
lightindicating
curves
of
rived
from
the
spectral
fits
used
to
obtainthe
X-rayluminosities
with a linear
coefficientof r = on0.51± 0.19,Bayesian
LINMIX_ERR
approachbased
methoddeveloped
calculatingtheresiduals.
If
there
was
no
intrinsic
relationbewindows
withsignificant
flaring activity
light curves
of
⇠ 95
%
that
the null hypothesis(no inorthe
a positive
correlation)isindicates
that
it
is
not
stronglya↵ectedby
selectione↵ects
and
Also,
the
negativeslope
suggestsananticorrelation,
the
large
the
COUP
sources.
The
corrected
valuesaredenoted“characcorrelationand inuncertainties
agreement withthe slopeof 1.07± acet linear
al. values2005).
Next,
performed
the linear regression
bya positive
Kelly (2007)
that takes
in bothvariables
into
tween
LX (Preibischand TheṀthe, theregressionofwetheresiduals
would
the
COUP
sources.
aredenoted“characacccorrected
“not”
true.
Thus,
anticorrelationissignificant,
even
under
the
uncertainty
of ±0.39
indicates
samplesources
sizeweistoosmall in stars in.theAdditionally,corrected
teristic”
X-ray
luminosities,
Lof
therefore
representativeforthepopulationofaccreting
0.22
we
found
in
our
analysis332
accretingin
F21.
X,thatchar
count as well as intrinsic scatter. The resulting relation reads teristic”. Additionally,
wecorrected
X-raynoluminosities,
Lof
indicate
correlation.
with
LINMIX_ERR
and
found
the
relationX, charconservative
assumption
strongaccretion
variability
contamiordertodraw
conclusions
about
the
correlationoftheresiduals.
thevaluesforthemore
recentGaia
distances
listedin
Table
0.5
In Fig. 7, the regression result is displayed as a dotted line. the values for the more recent Gaia distances listed in Table
the ONC.
the e↵ectof X-ray variability, we tried to cor! athetypical!de- in nationToofminimize
oursample.
• 降着率と星質量は有意に相関する.
!
!
The
linearcorrelation
coefficient
of
r
=
0.19
±
0.32
points1.
We
assumeduncertainty
of
L
of
0.15
dex
We
assumethatagreementoftheregression
results
from
our
X,
char
M
h Ṁacc i
1. Werect
assumeda typical
uncertaintyof LX, charof the0.15X-raydex deLM
theinfluenceofflaring
activity
fromluminosities
X,
char
•
しかし,
論⽂にEq(8)の散布図がない.
=
(
8.05
±
0.15)
+
(1.07
±
0.43)
·
log
,
(8)
log
smallersamplewiththeresultsofthelargersample
used
in
F21
rived
from
the
spectral
fits
used
to
obtainthe
X-rayluminosities
this
direction
as
well.
This
resultiskindofexpected
since
the
log
=
(30.33
±
0.15)
+
(1.54
±
0.45)
·
log
,
(9)
1
rivedusingfrom thespectralfits usedtoobtainthe
X-rayluminosities
M
M
yr
the
“characteristic
countrates”
determined
by
Wolk
et
al.
1
indicates
that
it
is
not
stronglya↵ectedby
selectione↵ectsand(Preibischetal.
2005).
Next,
weperformed
thelinearregression
observed
anticorrelation
inthelargersample
of
F21
is
weak,
as
M
erg
s
4.3. Relation between X-ray activity and accretion (Preibischetal.
2005).
Next,
weperformed
thelinearregression
Summaryandas the quiescentlevels of X-ray emission outside the time
0.0conclusions
therefore
representative
forofthepopulationof accreting stars in with5.(2005)
with
LINMIX_ERR
and
found
the
relationsuggestedbythemodelsX-ray
drivenphotoevaporation.
LINMIX_ERR and found the relation
with a linear correlation coefficient of r = 0.51 ± 0.19, indicating windows
with significant
flaring activity
in the light
curvesof± 0.16, indicating
the Since
ONC.
!
!
it
is
known
that
the
accretion
rates
as
well
as
the
X-ray
The
main
goal
of
this
work
is
to
check
whethertheaccretionwithalinearcorrelation
coefficient
of
r
=
0.65
!
! monitoring
presented
theresultsofamultiyear
photometrica positiveof 1.07 ± We
LX, charcorrelation and in agreement with the slope
M
the
COUP
sources.
The
corrected
valuesaredenoted“characL
M
X,
charluminosities
tendtodecreasewithincreasingage(e.g.,
Preibischlog we found= in(30.33± 0.15)+of(1.54±accreting
0.45)
· log
,in F21.
(9) log of accretion
rates
of0.15)young
stars±in0.45)
the ONC.
TheTheobservations
variability
introduces
a
bias
in
the
detected
anticorrelationbea
significant
positivecorrelation.
slope
isflatter
than
the
one
0.22
our
analysis332
sources=(30.33
±
+
(1.54
·
log
,
(9)
1
1
Mal. 2012), itergs
we corrected
teristic”
X-ray luminosities, LX, char . Additionally,&
Feigelson
2005;
Telleschi
et
al.
2007;
Manara
et
M
erg
s
4.3.
Relation
between
X-ray
activity
andaccretion
carried out with
the WWFI instrument and comprises phoIn Fig. 7, the regression result is displayed as a dotted line. were
-0.5
the values for 0the 0 more recent
Gaia distances listed in Table0
canbe
expected
thatthescatterintherelationcanbereducedtometryin the u , gcoefficient, and r offilters.
Accretion
rates were estiWe
assumethattheagreementof checktheresults
from our
Articlenumber,page
8workof 12is to
with
amain
linearcorrelation
coefficientof rregression= whether
0.65 ± 0.16,
indicating
The
goal
of
this
the
accretionwithalinearcorrelation
r
=
0.65
±
0.16,
indicating1.
We
assumedatypicaluncertainty
of
L
of
0.15
dex
deX,
char
by
correcting
not
only
for
the
mass
but
also
for
the
isochronalmated fromthe observed
displacementofthe thansourcepositionsinsmallersamplewith correlation.
theresultsof detectedthelargersampleusedin F21
avariability
significant
positiveThe
slope
isanticorrelation
flatterthanthebeone
introduces
a
bias
in
the
a
significant
positivecorrelation.
The
slope
is
flatter
the
one
rived
from
the
spectral
fits
used
to
obtainthe
X-rayluminosities
age. Therefore, we fitted planes to the quantities
withastandard
color-color diagram by modeling the colors as a combination
X線強度と星質量は有意に相関する.
indicates that it is not strongly a↵ected by•selectione↵ects and a(Preibisch
et al.-1.0
2005).
Next,photospheric
we performed
the linear regression
Levenberg-Marquardt
regression and regarded
both
the
mean
of
an
empiricallydetermined,
contribution
and anArticlenumber,
page
8
of
12
しかし,
論⽂にEq(9)の散布図がない.
therefore representative for the population• ofaccretingstars in with LINMIX_ERR and found
log(T
[K])
eff
the relation
mass and the mean isochronal age simultaneously. Uncertain- accretion model.
the ONC.
3.7 3.6 3.5
!
!
ties
were
not
taken
into
account
in
this
step.
Then,
we
performed• そこで共通する星質量への依存性をEqs.(8)&(9)から補正したうえで, 降着率とX線強度のThe
results
allowed
us
to
study
how
the
accretion
rate
varies
LX, charMthepartialregression
withtheresiduals
calculated
from
these
relogtimescales= (30.33± 0.15)(1.54± 0.45)· log interquar, (9)onupto a few
years.+We
found
a typical
相関を調べる(この図も論⽂にないが, 先⾏研究と整合するらしい).
1-1.5
M varierg sof ⇠ 0.3
lations
and found:
4.3.
Relation
between
X-ray
activity
andaccretion
tile
range
dex.
We
showedthat
theaccretion
rate• 降着率とX線強度の相関の残差(それぞれの変化分)を調べた. ⇨ Eq.(11) and Fig.8
-1.0
-0.5
0.0
0.5
1.0
!
abilityhaslikelynot introduced
a biasinourprevious
study
offit(M, t))log(L)±- 0.16,
log(L
X,
char
X,char,
The main
goal
of
this
work
is
to
check
whethertheaccretionwithalinearcorrelation
coefficient
of
r
=
0.65
indicatingh Ṁacc i
the relation between X-ray activity and accretion rates (Flais=
logvariability
introduces a bias in the detected anticorrelation be- a significant positive correlation. The slope is flatter than the one
Ṁacc (M, ⌧)chlen et al.Fig.
2021),
where weaccretion
reported arates
weakvs.anticorrelation
8.
Residual
residual X-ray luminosities for the!
between theONC
quantities,
supportingthetheoreticalmodelsof
XLX, char
Article number, page 8 of 12
taking both the mean mass and the mean isochronal age depen0.00 ± 0.09 + ( 0.37 ± 0.24) · log(11) ray driven photoevaporation(Ercolanoetal.
2008b,2009;
Drake
LX, char (M, ⌧)
dence
into
account.
The
dotted
line
shows
the
best
fit
obtainedwithet al. 2009; Owen et al. 2010; Picogna et al. 2019).
LINMIX_ERR. The solid bars indicate the maximum and the minimum
with a linear correlation coefficient of r = 0.36±0.23. A scatter Acknowledgements. We wish to thank the referee for helpful suggestions. We
residualaccretion rates for each source.
plot of the residuals together with the regression result is shown thank the LMU
master physics student Benedikt Mayr for his contribution toute of Astronomy and Astrophysics, Academia Sinica, 11F of Astronomy-Mathematics Building, AS/NTU No.1, Sec. 4, Roosevelt
Rd, Taipei 10617, Taiwan, ROC
11Gap
and
Inner
Disk Structurein theShinjuku-ku,
Strongly
AccretingTransition Disk of DM Tau
DivisionOpeningof Liberal Arts,
Kogakuin
University,
1-24-2 Nishi-Shinjuku,Tokyo
163-8677, Japan
onal Astronomical Observatory of Japan (NAOJ), National Institutes of Natural Sciences (NINS), 2-21-1 Osawa, Mitaka, Tokyo
Logan Francis ,1, 2 Nienke van 181-8588,Japander Marel ,3, 1 Doug Johnstone ,4, 1 Eiji Akiyama,5 Simon Bruderer,6
13 Department of Astronomy,
7, 8, 9 7-3-1 Hongo, Bunkyo-ku, Tokyo10113-0033, JapanThe Universityof Tokyo,
Ruobing Dong ,1 Jun
Hashimoto,Hauyu Baobab Liu , Takayuki Muto,11 and Yi Yang 12, 13
1 Department of Physics and Astronomy, University of Victoria

ABSTRACT

3800 Finnerty
Elliotof
Building
Large inner dust gaps in transition disks are frequently
posited Road,
as evidence
giant planets sculpting
Victoria,
BC, V8P 5C2,
Canadagas and dust in the disk, or the opening of a gap by
photoevaporative
winds.
Although the former
2 loganfrancis3@uvic.ca
hypothesis is strongly supported by the observations of planets and deep depletions in gas within the
LeidenUniversity,
P.O. atBox
9513,
2300-RAThe Netherlands
gap some disks, many3T
TauriObservatory,stars hostingLeidentransition
disks accrete
rates
typicalfor anLeiden,
undepleted
Herzberg
Astronomy
and occurs
Astrophysics,
West We
SaanichVictoria,
BC, V9E 2E7, Canada
disk, raising the4 NRC
questionof how
gap opening
in these5071
objects.
thus Rd,
present
an analysisof 5thestructureof the transition
disk aroundthe TArts,
Tauri
star DM Tau,
which is strongly
Divisionof Fundamental
Educationand Liberal
Departmentof Engineering,
Niigataaccreting
Institute of Technology 1719 Fujihashi,
945-1195,
Japan code, we fit0.02). Using Niigata
the DALI
thermochemical(⇠ 10 8.3 M yr 1 ) and turbulent (↵ = 0.078 ±Kashiwazaki,
6 Max-Planck-Institutfür theExtraterrestrische
Giessenbachstrasse
1, gas85748
Garching, Germany
disk models to simultaneously
reproduce
accretion rate, Physik,high levelof turbulence, thetraced7 Astrobiology
by ALMA band
6 observationsof 12National
CO, 13 CO,
and C18of
O Natural
J=2–1 lines,
and theobserved
dust
emission
Center,Institutes
Sciences,
2-21-1
Osawa,
Mitaka,Tokyo 181-8588, Japan
8 Subaru
from the mm continuum
andspectral energy
distribution.
WeObservatory
find a shallow
depletion
in gasTokyo
surfaceTelescope,
National
Astronomicalof Japan,
Mitaka,181-8588, Japan
9densityof ⇠ of
10 Astronomy,
relative to theouterand aGraduate
gas rich innerdisk isfor
consistent
the observations.
CH2D円盤の外側に⽐べて,内側は(少Department
SchoolofdiskScience,
University
Advancedwith
Studies(SOKENDAI),•Mitaka,
Tokyo 181-8588,
Japan
なくとも⼀桁ほど)ガス枯渇が進んで
implied by theAcademia
gap depth
is in tension
predictions for dust trappingThe
planetofmassof < 1 M
10
JupAstrophysics,
Institute
Astronomy
andSinica,11F of with
Astronomy-Mathematics
Building, AS/NTU
No.1, Sec. 4, Roosevelいない.
in a highly viscous disk, which requires a more massive
planetof of Taiwan,⇠ 10MJup. Photoevaporative
Rd, Taipei
10617,
ROC
11 Divisionmodels including
a dead ofzone
can qualitatively
reproduce
some1-24-2
featuresof the DM Tau
disk, but still
LiberalArts, Kogakuin
University,
Nishi-Shinjuku,
Shinjuku-ku,
Tokyo 163-8677, Japan
struggle
to Astronomical
explain the highaccretion ofrates
and(NAOJ),
the observed
mm continuum
12 National
ObservatoryJapan
National
Institutes offlux.
Natural Sciences (NINS), 2-21-1 Osawa, Mitaka, Tokyo
181-8588,Japan
13 Department of Astronomy, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan

Keywords: stars: formation - protoplanetary disks - planetary systems – protoplanetary disks - accretion, accretion disks - stars: variables: T Tauri, Herbig Ae/Be

ABSTRACTTable 2. Image properties

AASTeX v6.3.1 Sample article

Line

Beam FWHMa (00 )Approx. Beam FWHM (au)RMSrobustmultiscale (pix)a

uvtaper (00 )

cont.
12
CO
13
CO
C18 O

0.045 ⇥ 0.032
0.312 ⇥ 0.296
0.301 ⇥ 0.286
0.892 ⇥ 0.840

7
62
60
178

0.018 mJy/beam
2.9 mJy/beam · km/s
2.7 mJy/beam · km/s
6.2 mJy/beam · km/s

0.5
2.0
2.0
2.0

0, 6, 20, 40
-

0.3
0.3
0.3

Note—
a6 The size given is the full width at half maximum(FWHM)of the major and minor axes of the clean beam ellipse.
Francis etal.

7• 観測結果と使⽤した
データのまとめ
AASTeX v6.3.1 Sample article

7b The pixel size is 5 milliarcsec for all images.

In the following Sections, we describe the detailed
structure of the DALI models used to interpret the DM
Tau observations and determine the gas and dust structure within the 21 au cavity.
4.1. Chemical Network and Gas Composition
The underlying chemical network is the same as inprevious DALI studies, and is based upon the reactions in the UMIST2006 database (Woodall et al. 2007)
with some extensions for non gas-phase reactions described by Bruderer et al. (2009) and Bruderer et al.
(2012). Molecular data for collisions and radiative rates
are taken from the LAMDA database (Schöier et al.
2005); the calculation of the reaction rate coefficientsfor CO in particular are described by Yang et al. (2010)
and Jankowski & Szalewicz (2005). The elemental composition of the disk gas used in the DALI chemical net-

work is assumed to be that of the di↵use ISM (Bruderer
et al. 2012), and the gas phase abundances of each element relative to hydrogen used are listed in Table 3.
Although we assume ISM abundances for carbon andoxygen, gas mass measurements of DM Tau from HD
imply that the CO abundance may be lower than the
ISM value, a point which we return to in our discussionof the modelling caveats in Section 6.2.
4.2. Surface density structure
The radial surface density structure for DM Tau isshown in Figure 4, and closely follows the transition disk
structure previously implemented in DALI by Bruderer(2013), with some modifications to account for the inner dust ring. The surface density profile for gas in theouter disk is a radial power law which transitions to an
exponential cut-o↵ at a critical radius rc , as motivatedFigure 3. Spectral energy distribution for DM Tau, redcircles are photometric measurements, black dots are Spitzer
IRS spectra. Top panel: overall SED; Bottom panels: zoom
in on the near to mid infrared portion.

Figure 1. Top panels: 1.3 mm dust continuum images of DM Tau and zoom in on the 21 au dust cavity; lower left panels:
moment 0 maps of each observed CO line; lower right panels: Flux of each CO line integrated over the disk vs the velocity. The
beam for the continuum and moment 0 images is shown in the bottom-left of each panel as a white ellipse.

Figure 2. Deprojected and azimuthally averaged intensity
profiles of the 1.3 mm dust continuum and moment 0 maps
for each observed CO line. The full width at half maximum(FWHM) of the beam is shown by the scalebar.

We assume= 1, which corresponds to a vertically
isothermal disk with an alpha viscosity prescription⌫ = ↵cs H (Pringle 1981) and radial temperature profile
T / r 1/2 (Hartmann et al. 1998). The surface density
for the dust in the outer disk follows the same relation,
Figure 3. Spectral energy distribution for DM Tau, red
but with a constant scaling set by the gas to dust ratio
circles are photometric measurements, black dots are Spitzer
gas/dust . To reproduce the observed dust outer and inner dust rings of DM Tau (Section
we introduce
twopanel: overall SED; Bottom panels: zoom
IRS 2),
spectra.
Top
depletion factors in surface density relative to the outerin on the near to mid infrared portion.
disk, dust1 and dust2 , which begin at radii of rcav = 21
and rinner = 3.5 au respectively. We find that the powerto bethe di↵use ISM (Bruderer
emed
withinthethat
21 auofcavity.
4. MODELLING

Although we assume ISM abundances for ca
work theis assumed
to be thatabundancesof the di↵use ISM
oxygen, gas mass measurements of DM Tau, andgas phaseof (Bruderereach eletal.
2012),
and
the
gas
phase
abundancesofeach
elFigure
3.
Spectralenergy distribution for DM Tau, redurementsofCO
emission
in
disks
can
imply
that
the
CO
abundance
may
be
lower
4.1.
ChemicalNetworkandGas
Composition
iveement
to hydrogenused
arelistedinTable
3.
are photometric measurements, black dots are Spitzer
relative to hydrogen used are listed circles
in
Table
3.
IRS spectra. Top panel: overall SED; Bottom
panels:zoom a point which we return to in ourISMvalue,
tracerofassumethegasabundances
surfaceeaAlthoughassumeISM
abundances
forisfordensity,
carbonandinon thenear
to midThe
underlying
chemical
networkthesame
as infraredin portion.
we
ISM
carbonandofthemodelling
caveatsin
Section6.2.
oxygen,gas
massmeasurements
of
DM
Tau
from
HD
ting
COstudies,emissiongas surface
viousthe
DALI
andtois abased
upon the reac-

mass measurements of DM TauWefrom
assume HD= 1, which corresponds to a vertically
imply that the CO abundance may be lower
than disk
the with an alpha viscosity prescription
isothermalnsinsimple;the UMIST2006
database(Woodallet al.the
2007)
not
the
CO
abundancevariestheCOabundance
may
be
lower
than
⌫ =discussion↵c H (Pringle 1981) and radial temperature profile
ISM value, a point which we return to in our
4.2.
Surfacedensity
structurehsomeextensionsfornongas-phasede(Hartmann
et al. 1998). The surface density
T /reactions
rofthemodelling
caveatsin
Section
6.2.
isk
due
to
photodissociation
inthe
upa point which we return to in our for
discussion
the dust in the outer disk follows the same relation,
ibed by Bruderer et al. (2009) and Bruderer
et scaling
al. set by the gas
Theradialsurface density structure for D
but
with
a
constant
to
dust
ratio
ezeouttowards
the cold6.2.
midplane. The
lling caveats
in Section.Spectral
To reproduce
thedistribution
observedshown
dust
outer
and
in-red
4.2.
Surfacedensity
structure
12).
Molecular
data
forcollisions
andradiative
rates
inFigure4,
andclosely
follows
thetran
Figure3.
energyfor
DM
Tau,
Figure 2. Deprojected and azimuthally averaged intensity
nerdust
rings of DMmeasurements,
Tau (Section 2),
we introduce
two
mical
modellingcode(Brudereretal.
circlesarephotometric
black
dots
are
Spitzer
profiles of the 1.3 mm dust continuum and moment 0 maps
depletion
factors
in surface
densityrelative
topanels:
the outerTheradialsurfacedensity
structure
for
DM
Tau
is
taken
from
the
LAMDA
database(Schöieretal.
structurepreviously
implemented
in
DALI
by
IRS
spectra.
Toppanel:
overall
SED;
Bottomzoom
for each observed CO line. The full width at half maximum
disk,
andmid infrared, whichportion.
begin at radii of r = 21
4.2.
Surfacedensity
structurehese
issuessolvingforthegas
tem(FWHM)
oftheFigurebeam byis shown
by thescalebar.
in
on
the
near
to
shown
in
4,
and
closely
follows
thetransition
disk05); the calculation of the reaction rate
(2013),
with some modifications to account= 3.5 au respectively. We findthat the power
and r coefficients
structurepreviously
implemented
in
DALI
by
Brudererlaw
rise in dust surface density of our model between
abundance
in
a
model
disk,
providedin particular
arestructure
described
by
Yang
et
al.
(2010)
ner
dust
ring.
The
surfacedensity
profileforbysurface
analyticalmodelsof a self-similar
viscously evolving
lCO
density
for
DM
Tau
is
We
assume=1,whichcorresponds
toavertically
21forand the
3.5 auincan provide a good match to the surface(2013),
withsome
modifications
toaccount
accretion disk with a time-independent radial distribuFigure 4. Surface density structure of the gas and dust
brightness
of
the
inner
dust
ring
withouttheneedtoinisothermal
disk
with
an
alpha
viscosityprescriptiond
Jankowski&
Szalewicz(2005).
The
elemental
comouterdisk
isaradial
power
law
which
transit
dust
structure
and
model
for
the
stellar
ner
dust
ring.
The
surfacedensity
profile
for
gas
in
the
model
in
DALI
used
to
simulatetheDM
Tautransition
disk.
& Pringle
1974;
tion of4,viscosityr : (Lynden-Bellgureand⌫ /closely
follows
thetransition
disk
troduce
additional1981)
modelandparameters
for its outer edge⌫=↵c
H
(Pringleradialtemperature
profile The dashed lines show the truncated power law profile for thesHartmannetal.
1998):
ition
theispreviously
disk
gaspowerusedin whichthe
DALI
chemical
netexponential
cut-o↵
ata
criticalradiusr,
as
1/2
outerofdisk
a radiallawtransitions
to
an
and
dust
surfacedensity,
allowingforasimplermodel.
c
gas
and
dust
with
no
depletion.
DALI
has
been
used
for
de(Hartmann et al. 1998). The surface density
Tby/ r Bruderereviously
implemented
in
DALI
For
the gas surface density, corresponding depletion#"
exponential cut-o↵atacriticalradiusr,
asmotivated✓ ◆✓ ◆cfor the dust in the outer disk follows the same relation,
factorsofandinterior to the same radii asrrssurfacedensity
withintransition
disk
h some⌃(r)modifications
toaccountin.
(1) butfor=⌃
expwith atheconstant
scaling set by the gas to dust ratio
4.3. Vertical structure and dust populationsrr
gas/dust . To reproduce the observed dust outer and inThe vertical density structure of the gas disk is aAg.
COobservations(BruderervanThe
surfacedensity
profile 2013;
for
gas
in
the
Deprojected
and azimuthally
averaged
intensity
ner dust rings of DM Tau (Section 2), we introduce two Gaussian with scale height angle h(r) = h (r/r ) ,the 1.3 mm dust continuum and moment 0 maps
depletion
factors
in surface density relative to the outer where is a radial power law index which allows for flar2015,2016),
andmoreover,also
persserveda radial
power
law
which
transitions
toan
CO line. The full width at half maximum
ing of the disk and the physical scale height is H ⇠ rh.
disk,
and
,
which
begin
at
radii
of
r
=
21
dust1
dust2
cav
of the beam is shown by the scalebar.
To approximate vertical dust settling, two populations
adiativetransfer
necessarycompar=3.5
aurespectively.
We
find
that
the
power
and
r
cut-o↵ ata criticalradius rfor,asmotivatedinnercof small (0.005 - 1 µm) and large (0.005 - 1000) µm
law rise in dust surface density of our model between dust grains are included; the small dust grains follow
models continuumof a self-similar viscously
mcaldust
andevolving
SED.
s
1/2

gas/dust

dust1

dust2

cav

inner

2

c

c

gas1

gas2

c

c

c

4.3. Vertical structure and dust populations

The vertical density structure of the gas disk is a
Gaussian with scale height angle h(r) = hc (r/rc ) ,where is a radial power law index which allows for flaring of the disk and the physical scale height is H ⇠ rh.
To approximate vertical dust settling, two populationsof small (0.005 - 1 µm) and large (0.005 - 1000) µm
etdustal. grains are included; the small dust grains follow
the gas with scale height angle h, while the large dust
=0.85 throughout
thethedisk.
The dust
default
set more
to flsconcentrated
grains aretowards
midplaneandcomposition
andscaleopacities
used
for thehave a reducedheightangleof radiativeh, with transfer
fixed toarestandard
ISM composition& Draine
0.2.a The
dust surface
density is (Weingartner
partitioned between
the2001),
and are describedin further
detail byflsBrudererand (1 two populations
by a large
grain fraction(2013).
Aromatic
Hydrocarbons(PAHs) isareby
fls ) for Polycyclic
the small grains.
Thelarge grain fraction
also included in the model, which are largely important
as a source of optical/UV opacity for heating the gas
and an additional pathway for H2 formation (Bruderer
2013). The PAHs are assumed to be entirely in theform of neutral C100 H25 , and their opacities are calculated from the work of Li & Draine (2001) and Draine
et al. (2007) using the methods of Visser et al. (2007).
The PAHs are assumed to be destroyed in regions of thedisk where the local UV flux is stronger than 106 times
the interstellar radiation field. The abundance of PAHs
throughout the disk is set by the parameter fPAH , which
controls the fraction of PAHs in the disk with respect toan ISM PAH fraction of 5% of the dust mass (Draine &Li 2007).

(Woitke et al. 2016). For the observed
accretion
rate ofThe parametersof the fiducial
DALI model fit are pro1
AASTeX
SampleTable
3. v6.3.1
Toshow
the article
accuracy with which
DM Tau of ⇠ 10 8.3 M yr vided, theinaccretion
luminosity
the fiducial model
reproduces
the observations,the outDM
Tau
Parameter
Space
=
0.24.
as a fraction of that from thestar is Table
Lacc3./L⇤
put of the DALI raytracer is compared with the deThis relationship assumes that
all of the gravitational
Parameterazimuthally averaged Value
projected
continuum, CO isotopoenergy liberated by accretionis radiated
away,
which
logues,SED and CO
line
DMfluxes
Tau integrated over the disk
⌃ Figureg cm comparisonin
5. A0.65
directis a reasonable approximation
for magnetosphericac- of the output images
10and moment maps
from the raytracer is provided in Figcretion onto pre-main-sequencestars (Hartmann
et al.
r
0.0343 auure 19 of the appendix.
r
2.5, 3.5, 4.5 au2016).
The radial profileof the dust emission is reasonably
r
21 au
well
described
by thehowever,the model underSimilarly, heating of the disk
from stellar
X-raysisr
124.1
au model,estimates
the 0.03,
emissionin radthe outerh0.06, 0.09
taken into account by addingan X-rayluminosityof disk beyond ⇠ 60
0.2observations of Hashimoto et al. (2021)
au. The ALMA30
1
LX = 10 erg s with a thermal
spectrumatT=X
0.015,
0.15,
1.0
have shown several low contrast dust rings are present at
71.0, 10.0
observed
10 K, which is reasonably consistentthese pointswith
in0.15,thetheouter
disk, which we do not attempt1.0(-50),
2.7(-2)30 as
1focustoreproduce
ouris on the dust cavity and inners 6.0(-10),(Dionatos
X-ray luminosity of LX = 1 3 ⇥ 10 erg1.0(-50),
6.0(-9), 6.0(-8), 6.0(-7), 6.0(-6)
The model
alsoemisagrees well with the obser0.2ofSED
et al. 2019) given the intrinsicdisk.
variability
X-rayvations,although
0.85 the exact shape of the near-infraredf (outer disk)sion. Finally, a background silicate
cosmicfeature
ray0.24
ionization
rate specifically, the flux
is not reproduced,
L
gas,c

2

gas/dust

subl

inner
cav
c

c

gas1
gas2

dust1
dust2

ls

T⇤ ⇠ 7 microns
3580isKoverestimated, suggesting some deatlog(Ṁ )of the inner-9.3,disk-8.3, structure-7.3 [log(M are
yr 1missing.
)]tails
Hashimoto
30
1LX
10 erg setal. (2021) have
performed radiative transfer modeling
TX
107 K
of
the same SED
and show that the dip at ⇠ 7 microns⇣cr
5 ⇥ 1017 s 1
issensitive to 1.0(-50),
the dust1.0(-5),
size 1.0(-3)
distribution between 1 andfPAH
3Gasau,and can be reproduced by a model with a depleAbundance
H1
tion
in the surface
density of micron sized dust grains
He 10 4 (their7.59(-2)
Figure 5). We have also experimentedofC
1.35(-4)
withvarying theamount of small dust grains from theN
2.14(-5)sublimation radius to 3.5 au and similarly find there isO
2.88(-4)
aMgsmall e↵ect 4.17(-6)
on the near to mid infrared shape, however,we were 7.94(-6)
unable to reproduce the exact observed
Si
SED
shape, whichS
1.91(-6)may require more careful modelling
Fe the inner ring
4.27(-6)ofgeometry that is beyond the scope ofNote—Numbers
this paper. written as a(b) indicate a ⇥ 10b . Where a range ofvalues is indicated, parameters denoted bold are the fiducial model
parameters. Parameters in red are the non-fiducial values for the caseof no inner disk.

relevant
dust denare show
midplanthe ⌧ =The sam
5 au offull disk
The gwith heation ofthe disk
molecul
CO is pit can su
sphere oinwarddecrease
The low
deeper ilayers indust temlayers odensity,
the loweis frozen
the COdecreaseeach ofkey radi
CO freeregion.
increase
temperasharp dr
optically
Withinrelevantelling,aformalmodel
fitting
based ontemperature 3580 K and luminosity 0.24 Ldisk.
and
CO
lines
to
the
The parameters of the fiducial DALI modelfit
are produst den
Stellar ofmodel
the stellar4.4.
properties
DMand
Tauradiation(Section field2). UV heat- 4.5. Model grid
AASTeXv6.3.1
Samplearticlevided in Table 3. To descent)show the accuracy
with whichisimpractical.
Ou
are show
projected
imagesaThecentralstar
isapproximatedasablackbodywiththefiducialmodel
reproducestheobservations,theouting from accretion is included by adding the contribuTable 3. DM 9
Tau Parameter Space
midplan
AASTeXv6.3.1
Samplearticle
The
grid
used
for
the
calculationconsists
of
200
cells
in
stead
to imuthally
obtainm4
put of the DALI raytracer
is comparedwith the ade-fiducial
on 4.6. Raytracing and Model Images
temperature
3580 K blackbody
and luminosity
averagedtionof an additionalwith 0.24
Tacc L= 10based
K andthe ⌧ =
Parameterazimuthally averaged Value
projected
continuum, CO isotopothe stellarluminosity
properties of DM Tau(Section
2).
UV
heataccretion
Theof
sam
the radial direction and 80 cells in the vertical direction.
reproducesthefeaturesmanneras
theobs
logues,SED and CO line
over the disk
DMfluxes
Tau integrated
Using
the
converged
outputofthethermochemical
5 au ofing from accretionis included by adding the contribu- 4.7. Model fitting approach⌃212present
throughout
0.65
g
cm
ingas,cFigure 5. A direct comparisonof the output images
To
resolve
the
parts
of
the
disk
with
sharp
transitions
dust
continuum,
the
CO,
4
full
disk
pare
the
model
SED
GM
Ṁ
tion of an additional blackbody ⇤with Tacc =modelling,
10 K and10
gas/dust
the DALI code performs andraytracing
simumoment maps fromthe raytracer is provided in Fig,
(2)
Lacc = ⇡The grsubl , and 0.0343 au 0 maps, and SED, with a,rinand dust
surfacedensity(i.e.,
rR⇤ gas Due
paaccretion luminosity
subl
inner
ure 19DALIof the appendix.
to
the
large
computational
cost
of
the
modobserved
SED
ovewith he2.5, 3.5, 4.5 To
aulations to produce model images andrinnerline
cubes.
The radial
profileof the dust emission is reasonably
2
r
21
au
ation
of
)
the
cell
size
grows
exponentiallystarting
from
these
r
cav
tures
within
the
21
au
dus
cav
and
R
the
where G is the gravitationalGM
constant,
M
⇤
⇤
by
the
photometry
gradient
elling,aformalmodel
fitting
approach(e.g.
well
described 124.1
by is
the
model, however, the model under⇤ Ṁrc
au permatchtheALMA
observations,theraytracing
the disk
=
⇡
,
(2)
L
i
acc
mass and radius of the star, andtheestimates
the 0.03,
emissionin radthe outerdisk beyond ⇠ 60 if the gas
=
10
,
where
=
0.01
au
radii,
i.e,stellar
hc is the0.06, 0.09
wishto
investigate
R⇤ Ṁdescent)cellaccreminminmoleculnear
tomid
infrareisimpractical.
Ourmodellingstrategyisinformedim-with the model disk at an inclination
ofobservations
35.2 of Hashimoto et al. (2021)
0.2
au. The ALMA
tion rate. Modelling the UV heating is particularly
CO is pminimum
radialcellwidth,iistheradialcell
index,
and
cial
model
withinthecavity0.015,
0.15,
1.0
gas1 shown several low contrast dust rings are present at
haveR⇤somethewhere Gofis 200
the gravitationalM⇤ and
In
order
to
obtai
portantfordiskchemistry,
but issubjecttoonsists
cells
in constant,stead
toobtainafiducial
model
which
approximately
it can su
and distance of 144 pc (see Sections these2gas2 and3)witha0.15, 1.0, 10.0
points in the outer disk, which we do not attempt
mass and radius
thepaucity
star, andṀ thestellarUV
accresphereouncertainty
due toofthe
ofdetermines
highqualityob-cell growth rate. The 200 radial
the
cells
with
the
stellaraccretionrat
1.0(-50), 2.7(-2)
dust1
simulations
using
to reproduce
as ourfocus is on the dust cavity and innerheservations,verticaldirection.
reproducesthefeaturesoftheobservations:
the
mm
pixel
size
of
5
milliarcsec.
The
model
SED
was
createdtion
rate. ModellingtheUV
heatingisparticularly
iminward
1.0(-50), 6.0(-10), 6.0(-9), 6.0(-8), 6.0(-7), 6.0(-6)
dust2
and di↵erent approaches using detailed steldisk.
The model
SED also
agreeswell with the modelobseraredistributed
in
this
way
to
provide
50
cells
from
0.01
in
the
fiducial
param
0.2
12
13
18
decrease
raytracingofthecportant
for
the
disk
chemistry,butissubjecttosomethlarsharptransitions
models havebeen used in other astrochemical
codes
dust bycontinuum,
CO,
and
C
O
momentraytracing
the diskfrom CO,0.09 to3000µm
andcalcuvations,although
0.85 the exact shape of the near-infraredfls (outer disk)
The lowc
to
3
au,
50
from
3.1
to
21
au,
and
100
from
21
au
to
uncertainty
due
to
the
paucityofhighqualityUV
obfect
of
the
inner
disk
and
(Woitke et al. 2016). For the observed accretion rate ofsilicatefeature0.24
is not reproduced, specifically,
the andflux SED. I
L
imagesdeeper i
latingF
overtheentireraytraced
image.
Sampling
of
e.,DM
r
,
r
,
and
0
maps,
and
SED,
with
a
particular
focus
on
the
strucsubl
inner
8.3
1
T
3580
K
servations,anddi↵erent
approaches
usingdetailedstel⇤at ⇠ 7 microns is overestimated, suggesting some deM yr ,thethe accretion
Tau of ⇠ 10edge ofluminosity
the calculation grid at 500 au. Vertically,80
layersininemission.
1
usedto
estimate
log(
Ṁ
)
-9.3,
-8.3,
-7.3
[log(Myr)]
tails approximatelyof the inner30diskwestructure
are missing. Hashimotolaramodels
haveused in
codes
the
SED
was
performed
atwavelengths
y startingfromwithin
the
21
au
dust
cavity.
Additionally,
=
0.24.
as
fractionof beenthatthesefromtheotherstar astrochemical
istures
Lacc /L1
dust tem⇤LX
10 erg settheal. (2021)
have
performedTo
radiative
transfermodeling
cells
arearrangedtofollowthescaleheightofgascompare
the
DALI
mo
7
(Woitkeetal.
2016).
For
theobserved
accretion
rate
of
,
after
which
and
r
layerso
T
10KThis relationship assumes that all of the gravitationalXccorresponding
tothephotometricdata,
exceptfor
theofthe ofsametheSED
andshow that the dip at ⇠ 7 microns=
0.01
au8.3isMthe
wishtoinvestigate
ifgas
surfacedensityfidu171
min⇣cr
5 ⇥ 10s 1
density,
DM
Tau
of
⇠
10
yr
,
the
accretion
luminosity
energy liberated by accretion disk.
is radiated away, which
and
CO
lines
to
the
observatissensitive to 1.0(-50),
the dust1.0(-5),
size 1.0(-3)
distributiondensity
between 1 and
via
⌃
,
f
gas,c
PAH
the
lowe
5-20
µm
range,
where
steps
of
0.33
µm
were
used
for
/L⇤ = 0.24.
as aa reasonablefractionof approximationthat andfrom the stariscial
Laccmodel
adialcell index,
within the cavity of the disk is 3Gas
commensurateisfor magnetosphericacau,
and can be reproduced by a model with a depleAbundanceis frozen
projected
images
and
mome
,
un
tors
This
relationshipassumes
that
all
of
the
gravitational
dust1
dust2
H
1
tion
in theModelsurfacedensityof micron sized dust grainscretiononto
pre-main-sequencestars(Hartmannetal.
comparison
withtheSpitzer/IRS
spectrum.
conthe CO
The
200
radial
cells
with
the
stellaraccretion
rate,
and
to
explorevariations
4He
7.59(-2)energy liberated by accretion is radiated away,
4.6.which
Raytracing and Model Images of 10 (their Figure 5).imuthally
We have also
experimented radial
averagedin2016).
approximately
repr
decreaseC
1.35(-4)
tinuumimageswerecreatedbyraytracingatawavewithvaryingtheamountof small dust grains from theisareasonableapproximationformagnetosphericacde 50
cells heatingfrom of0.01
the X-rays
fiducial
model parameters to understandtheefeach ofSimilarly,
the disk frominstellarisN
2.14(-5)mannerastheobservations
of
DM
Tau
displaysublimation radius to 3.5 au and similarly
find
there
is
Using
the
converged
outputofthethermochemical
cretiononto
pre-main-sequencestars(Hartmannetal.
key radilengthof
1.3mm,
followedbyconvolution
with
the
obO
2.88(-4)
taken
intoaccountbyaddingan
X-rayluminosity
of
100
from
21 1 au toaMgsmallone↵ecttheon theCOnear to mid infrared shape, howfect of the inner disk and cavity structure
4.17(-6)
pare
the
model
SED
produc
CO free302016).
ingalowabundan
LX = 10 erg s with a thermal
spectrumatT=modelling,
the
DALI
code
performs
raytracing
simuX continuum beam and deprojection
ever,
weby
weretheunableto reproduce the exact observed
servedsame
Si
7.94(-6)
region.
7
Similarly,
heating
of
the
disk
from
stellar
X-raysis0010au.
Vertically,80
emission.
SED
shape, which
require more careful
K, which is reasonably consistent with the observedS
1.91(-6)may observed
SEDmodellingover theincrease
tot
PAH
fractionf
PAHlationstoproducemodel
imagesandline
cubes.
To
procedure
Sectiontaken luminosity
into accountby =addingan10X-rayluminosityof as described for the observations
30
1
Fe the innerin
4.27(-6)ofringgeometry that is beyond the scope of
1
3
⇥
erg
s
(Dionatos
X-ray
of
L
X
temperaleLXheightgas
Tocomparethe
DALI
modelsforthemmcontinumm
30 of the1bythephotometry
and
the
S
b
ultimately
reNote—Numbers
written as a(b) indicate a ⇥ 10 . sion,
Where a rangeof
thispaper.
=
10
erg
s
with
a
thermalspectrum
at
T
=
match
the
ALMA
observations,theraytracingisperX
et al. 2019) given the intrinsic variability of 3.
X-raySimilarly,
emismodel
line
cubes
were
createdassuming
sharp drvaluesisindicated,parameters
denoted
bold
are
the
fiducial
model
7
which isareasonably
consistentwiththeobserved
10
and
CO
lines
to
the
observations,
we
took
the
outputtomidparameters.
Parameters inderednear
are the non-fiducial
valuesinfrared.
forthe case that opticallysion.K, Finally,
backgroundcosmicray
ionization
rate
We
note
PAH
formed
with
the
model
disk
at
an
inclination
of
35.2
30
1 keplerian velocity structure of the diskaandraytracof no inner disk.
X-ray luminosity of LX = 1 3 ⇥ 10 erg s (Dionatos
Within•観測結果とfiducialモデルの⽐較

Figure 5. Comparison of the DM Tau observations (black lines and red circles for SED photometry) and fiducial model (bluelines). Top Row: deprojected and azimuthally averaged mm continuum intensity, Overall SED, and SED zoomed in on the nearto mid infrared region. Middle Row: deprojected and azimuthally averaged intensity for the moment 0 maps of the 12 CO, 13 CO
and C18 O lines. Bottom row: Flux integrated over the disk vs velocity for the 12 CO, 13 CO and C18 O line cubes. The spectral
resolution of the observed line cubes is 1 km s 1 , while the models are sampled in 0.5 km s 1 steps

of the dust continuum emission at 21 au
disk to shieldthe
thedisk
midplanefromradiation,tgenerale↵ectsofvariations
intheismodel
parametershig
near-infraredexcessoftheSEDslightly•観測結果とfiducialモデルの⽐較
on
the RESULTS
physical conditions of the disk.
5.
ghter while the emission in the outer disk structure,
freezeoutzone
isentirelyabsentinthecavity,
a
gas
and
dust
temperature,andCOabundanc
the the 2.5 au case but underproduced in the
5.1. Fiducialmodel
au is unchanged; the SED is entirely miss- underscores
onset
beginsafew
aufurther
out
in
the
disk
fro
the
valuableroleofthermochemical
mod
case,
suggesting
that
r
is
well
constrained
ov
inner
In
this
Section,we
describethefiducial
DALI
mThe
parametersofthefiducial
DALI
model
fit
are
proSurface
densityatstructureandselected
2D physicalquantities
fromedgethe fiducialDM
Tau
modelmodel.
plottedInfrom
fromsilicates
10
µm
and
now
peaks
in
21
au
cavitythan
inthefiducialcoellingforinterpretingline
emissionfromdisks.
vided
in Table
3. To anshowincreased
the accuracy with
which
12
13
18
largervalue
ofrcauses
gas
temp
inner
=
1
surfaces
for
CO,
CO
and
C
O
lines
are
overlaid
on
each
2D
DALI
panel
as
the
solid,
dashed,
an
for
DM
Tau,
explore
how
well
the
inner
disk
and
the
fiducialmodel
reproducestheobservations,theoutar-infrared rather than at ⇠ 20 µm. The
to the results of DALI modelling by Bruderer (201
(Figureresultinginatheslight< rinnerput isofthe
DALI 23),
raytracer
is comparedde- onectively while the approximate freezeout temperatureatofradii
CO (25.5K)
overlaid
as theorange with
contournner
disk
causesthesee↵ectsbyallowing
find
CO
is
not
photodissociatedwithinthedust
structureareconstrainedbyobservations,andconprojected
azimuthally
averaged
continuum,COisotopo5.2.
Parameterspace
e panels.
in the model CO
withinthe 21cavit
logues,emission
SED and CO line
fluxes integrated
overauthe diskradiation to reach the cavity edge atthe
21 au
without
an
inner
disk,
which
we
discuss
furtherigenerale↵ectsofvariations
the
model
param
in
Figure5. Ae↵ects
directincomparisonofthe output
images ke
Wefurther
exploretheofvaryingvarious
ever,
this
is
poorlyconstrained
by
the
CO
obser
andtheCOcavity,
abundance
E↵ect
Ring
and5.2.1.
moment maps
from ofthe Inner
raytracer Dustis provided
in Figg, gas
theandhotdustdusttemperature,component inas parameterstion 6. for thefiducialmodel
to demonstrate hoonthephysicalconditions
the
disk.
ure of19
ofresolution.
theof
appendix.
due
to
the
lack
res
of thermochemical
modure the21. valuable
The CO roleemissionis brighter in the
The radial profile of the dust emission is reasonably
Dustsurfacedensity
5.2.2.
wellconstrainedbytheobservations
thefiducial
mod
12
well
described
by themodel, however,thedescribedmodel under- ininterpretingline
emissionfromdisks.
9,
but
for
the
e↵ectsofchanging
thewithin3.5
auin the
DALI
model
Many
transition
disks
are
observedtohaveacularly
the more optically thick dust
CO depletion
estimates
the emission
inthe outerdisk beyond⇠ 60 ga
is and The
investigate
the
e↵ect
on
the
radiation
field,
structure
withinthecavitya tranGas
surfacedensity
5.2.3.
au.within
The
ALMA
observationsof cavity
Hashimotoet ofal. (2021)
disk dustor5.1.
ringthecentral(Francis&
Fiducial
model
otopologues, which reach 2-3 times higher and dust
temperature,andCOabundance
within
th
have
shown
severallow
contrast
dust
rings
are
present
at
5.2.
Parameter
space
disk
can 2020),
play animportantrole toin e↵ectregulating
t
Marel
which
is
expected
the
su
at
their
⌧
=
1
surfaces;
in
the
outer
disk
these
pointsin theouterscale
disk,
which
we density
do not attempt
culate the vertically averaged gas The parametersDiskheight5.2.4.
cavity.
The
structure
of
the
gas
surfacewitofthefiducial
DALI
model
fit
ar
diationfield
throughout
the
disk
and
is
typicallto
reproduce
as
our
focus
is
on
the
dust
cavity
and
inner
ther
explorethee↵ects of Without
varying various
key
CO in the cavity by shielding the gas in the cav
ion
islargely
unchanged.
an
inner
disk.
The model
SEDcanalso agrees
well withimportan
the obserd by the gas density for the fiducial21
aucavityofDM
Tauprovidestrainedbythemmcontinuum
emission
and
SE
vided
in
Table
3.
To
show
the
accuracy
withrsFigure
for9. Comparison
the fiducialmodel
todemonstratehowstellarradiation(Bruderer2013).
Here,
we
exp
vations,although
the
exact
shape
of
the
near-infraredof the fiducial DM Tau DALI model and a model with no inner dust disk (
=
= 10
for
modelsofplanetdisk
andaccret
↵
is
by
with
the
DMpoorly
Tau observations,constrained
as described in Section 5.2.1. Top
Row:obserdeprojected and azimuthally averaged
mm continuum
The
scale
height
of
the
gas
disk
depends
silicatefeatureis notinteraction
reproduced,
specifically,thefluxontrainedbyzoomed
thein onobservations
thefiducial
model
of
the
inner
dust
ring
at
3.5
au
in
DM
intensity, and SED
the near to mid infrared region.
MiddleRow: deprojected
and azimuthally averagede↵ect
intensity
thefiducialmodel
reproduces
observations,at ⇠ 7 microns is theoverestimated,suggesting some de-th
for the moment 0 maps of the CO, CO and C O lines.
inthecavity.
Here,
we
vary
and
wit
ks,
but
is
typically
assumed
to
be
gas1
gas2
drostatic
equilibriumbalancebetweenthegasof the inner fromdisk structure
missing. Hashimoto
vestigate the e↵ect on the radiation field, gasremoving ittailsentirely
our arefiducial
model,
put
of
the
DALI
raytraceriscomparedwithth
et modelal. (2021) have
performed
radiative
transfermodeling
fiducial
DALI
toassess
how
well
the
rad
au
is one of few
where turbua↵ect
thestellvtemperature,
anddisks
CO abundance
within the and gravity,ting bothandandto amounte↵ectively
zero
dust2
ofdust1
thewill
same
SED
and
show
that
the dip at ⇠ 7 ofmicrons
50 profiles
projected
azimuthally
averagedCOis e↵ectsensitive
to constrain
thedustcontinuum,
size distribution
between1 andisoemission
can
the
gas
surfacebeen detected through high velocity
.
The
on
the
observations
is
shown
i
10
ation absorbed3 bythedisk
surfaceasafunctio
au, and can be reproduced by a model with a depledust1

12

13

dust2

50

18

and
comparetheresults
with
the
gas
surfacetion
inthesurfacedensityofmicronsized
dust
grainslogues, SED
and
CO
line
fluxesintegrated
overth
ns of the CO line (Flaherty et al.
dius. In our DALI
model
setup,
the
scale
heigh
of 10
(their Figure 5). We have also experimented
impliedby
theaccretion
rate
for
an
↵-diskmodbulent broadening of 0.25-0.33 c , orwith varying the amount of small dust grains from the
124

7. CONCLUSIONS

the disk continuum emission, 12 CO, 13 CO, and C18 O
In this paper, we have used thermochemical modelling
es,withSED,
and stellar
accretion rate
to constrain
the DALI
code in combination
with ALMA
images theoperties of the gas cavity and inner disk of DM Tau.
ur major conclusions are the following:

Hashimoto et al. (2021) and Calvet et al. (2005)• The
planetmassspeciesof . 1M
shallow
J implied
Most
volatiles
arelikely inby
thethe
gas
phase
gas
gap and↵=0.078±0.02
meawithin
3.5 aucavity,and mayfrom
thusturbulence
be incorporated
surements(Flaherty
al. 2020), is inconsistent
into planets
forming etwithin.
with the larger planet masses required for dust
trapping in a highly viscous disk. A deeper but
narrower gap unresolved by our ALMA observations could allow for a higher planet mass consistent with dust trapping by a giant planet.

• The gas surface density within the 21 au dust cavity18.
is appears
to=show
only a shallow
depletionbyigure
Midplane (z0) temperature(top panel)andbundanceof various
species
into ices(bottoma factorof ⇠ volatile
10, whichis locked
consistentwithestimates • The inner disk of DM Tau absorbs much of theanel)ofin thethe fiducial
DM Taudensity
model within
10 au.
X/H+H
2
stellar radiation, lowering the gas temperature ingas surface
in the
inner
disk
from
enotes the ratio of the number densities of ice X and thethe cavity. In contrast to the DALI transition disk
thestellaraccretion
rateandviscous↵-diskmodotal gas density. N.B.: the gas and dust temperature at themodelling of Bruderer (2013), CO is able to survive
midplaneare essentiallyin thisregionofof thethe disk.
elling,however,identical
the gasdensity
inner disk
in the shallow gas cavity of DM Tau even without
itself is poorly constrained and could be enhanced
an inner dust disk providing shielding from pho.5 au,
midplaneis cold100(⇠ relative
30 K) andices gapby thea factorof 10
to mosta shallow
todissociation, which is likely the result of H2 forhouldbe frozen
out.
As the innercavityat 3.5 auandstill beconsistentwithouredgeinterpretationofmation at higher gas densities by the H 1 + H !
approached, the midplane temperature rises to ⇠ 200
H2 + e reaction, which provides an alternative
the CO observations and the accretion rate.

K, sufficient to remove essentially all of the ices on therains. Any planets formed within 3.5 au would thus be• The dead zone + photoevaporation models of
xpected to incorporate solids and gas which are respecetdepletedal. (2021)
can produce
transition disks
velyGárate
relatively
or enrichedin volatiles.

with an inner dust disk, outer ring, and high ac7. but
CONCLUSIONS
cretion rate,
are still unlikely as an explanaIn thiswe gaphave inused
thermochemical
modellingdisk, as
tionpaper,for thethe
DM Tau transition
with they
the DALI
code reproduce
in combination
ALMA
imagesdo notthewithhighstellaraccretion
rate or bright mm emission in the dust disk. However, non-viscous angular momentum transport in
the inner disk and uncertainties in the model dust
opacities may contribute to the disagreement between their models and the observations.
• The planet mass of . 1MJ implied by the shallow• Our SED modelling indicates DM Tau is extremely
depleted in PAHs relative to the ISM value, andhas a PAH fraction of at most 10 5 that of theISM, assuming that all PAHs are in the form ofneutral C100 H25 and that the PAH distribution
is uniform throughout the disk. Unlike previous
studies of PAH survival in transition disks (Geerset al. 2006, 2007), our modelling includes UV flux
from accretion in assessing the PAH abundance,
resulting in a lower inferred value.

pathway to CO formation.

• Our modelling of the inner disk indicates a ⇠3.5 au cavity, consistent with previous studies by
Hashimoto et al. (2021) and Calvet et al. (2005)
Most volatiles species are likely in the gas phasewithin 3.5 au cavity, and may thus be incorporated
into planets forming within.

APP

3 Frontier Research Institute for Interdisciplinary Sciences,

University, Sendai 980-8578,
Cosmic-ray ionization rateTohoku
in protoplanetarydisksJapan
with sheared magnetic fields
4 Astronomical Institute, Graduate School of Science,

1, 2
TohokuUniversity,
980-8578,
Japan
Yuri I. Fujii
andSendai
ShigeoS. Kimura

3, 4

1 Graduate School of Human and Environmental Studies,

ABSTRACT

Kyoto University, Yoshida-Nihonmatsu, Sakyo, Kyoto 606-8501, Japan
We investigate the e↵ects of magnetic field configurations on the ionization rate by cosmic rays in
2 Department of Physics, Nagoya University, Furo-cho, Chikusa-ku, Nagoya, Aichi 464-8602
protoplanetary disks. First, we consider cosmic-ray propagation from the interstellar medium (ISM)
3 Frontier Research Institute for Interdisciplinary Sciences,to the protoplanetary disks and showed that the cosmic-ray density around the disk should be 4
TohokuUniversity,
SendaiJapan
times lower than the ISM value.
Then,
we computethe 980-8578,
attenuationof cosmic rays in protoplanetary
4 Astronomical Institute, Graduate School of Science,disks. The magnetic fields in the disk are stretched to the azimuthal directions, and cosmic rays need
University,WeSendaito detour while propagating Tohoku
to the midplane.
show980-8578,
that the Japandetouring e↵ectively enhances thecolumn density by about two orders of magnitudes. In the case of the disk around IM lup, this increases
the ionization rate over an order of magnitude
for r & 100 au. On the other hand, for r . 100 au, theABSTRACT
cosmic rays are shielded at the disk midplane while the ionization rate is also enhanced at z ⇠ 2H.
We investigatethee↵ectsofmagnetic
field
configurations
ontheionization
ratebycosmicrays
Our results are consistent with the recent ALMA observation that indicates the radial gradient in therotoplanetaryFirst,
weTheconsidercosmic-ray
propagationfrominterstellar
mediumcosmic-raydisks.
ionization
rate.
elevated
ionization ratein the outerradiitheof disks
may activatethe (ISMo the magnetorotational
protoplanetary disks
and thatshowedthat thecosmic-ray
around thedisk should
instability
was thought
to besuppresseddensity
due to ambipolardi↵usion.
These beresultsthan
will havea strong
influenceon wethe computedynamical theand attenuation
chemical evolutionsof protoplanetarydisks.
mes lower
the ISMvalue.
Then,of cosmic
rays in protoplaneta

sks. The magnetic fields in the disk are stretched to the azimuthal directions, and cosmic rays neeo detour while propagating to the midplane. We show that the detouring e↵ectively enhances th
Keywords: Cosmic rays (329), Protoplanetary disks (1300), Magnetic fields (994), Ionization (2068)
olumn density by about two orders of magnitudes. In the case of the disk around IM lup, this increas

Chemical evolution in ices on drifting, planet-forming pebbles
1
1
Christian Eistrup and Thomas Henning

Max Planck Institute for Astronomy (MPIA), Königstuhl 17, Heidelberg, Germany
e-mail: eistrup@proton.me
August 17, 2022
ABSTRACT
Context. Planets and their atmospheres are built from gas and solid material in protoplanetary disks. Recent results suggest thatsolid material like pebbles may contribute significantly to building up planetary atmospheres. In order to link observed exoplanet
atmospheres and their compositions to their formation histories, it is important to understand how icy pebbles may change theircomposition when they drift radially inwards in disks.
Aims. Our goal is to model the compositional evolution of ices on pebbles as they drift in disks, and track how their chemical evolution
en-route changes the ice composition relative to the ice composition of the pebbles in the region where they grew from micron-sizedgrains.
Methods. A state-of-the-art chemical kinetics code is utilised for modelling chemical evolution. This code accounts for the timeevolving sizes of the solids that drift. Chemical evolution is modelled locally for 0.1 Myr at two starting radii, with the micron-sized
solids growing into pebbles simultaneously. The pebbles and local gas, isolated as a parcel, is then exposed to changing physicalconditions, intended to mimic the pebbles drifting inwards in the disk midplane, moving to 1 AU on three di↵erent timescales. A
modelling simplification is that the pebbles are not moved through, and exposed to new gas, but stay in the same chemical gas
surroundings in all models.
Results. For ice species with initial abundances relative to hydrogen of > 10 5 , such as H2 O, CO2 , CH3 OH and NH3 , the abundanceschange by less than 20% for both radii of origin, and for the two smaller drift timescales (10kyr and 100kyr). For less abundant ice
species, and the longest drift timescale (1Myr), the changes are larger. Pebble drift chemistry generally increases the ice abundancesof CO2 , HCN and SO, at the expense of decreasing abundances of other volatile molecules.
Key words. planet formation – astrochemistry – pebble accretion

CO(J = 1-0) Observations toward the Filamentary Cloud in the Galactic Region of 153. 60 6 l 6156. 50 and 1. 85 6 b 6 3. 50

6v1 [astro-ph.GA] 16 Aug 2022

Weihua Guo,1, 2 Xuepeng Chen,1, 2 Jiancheng Feng,1, 2 Li Sun,1, 2 Shiyu Zhang,1, 2 Chen Wang,1 Yang Su,1, 2
Yan Sun,1, 2 Qingzeng Yan,1 Shaobo Zhang,1 Xin Zhou,1 MiaoMiao Zhang,1 Min Fang,1 and Ji Yang1, 2
1 Purple Mountain Observatory and Key Laboratory of Radio Astronomy, Chinese Academy of Sciences, 10 Yuanhua Road, Nanjing

210034, China
2 School of Astronomy and Space Science, University of Science and Technology of China, Hefei, Anhui 230026, China

ABSTRACT
We present observations of J=1-0 transition lines of 12 CO, 13 CO, and C18 O towards the Galactic
region of 153. 606 l 6 156. 50 and 1. 856 b 6 3. 50, using the Purple Mountain Observatory (PMO)
13.7 m millimeter telescope. Based on the 13 CO data, one main filament and five sub-filaments arefound together as a network structure in the velocity interval of [ 42.5, 30.0] km s 1 . The kinematic
distance of this molecular cloud (MC) is estimated to be ⇠ 4.5 kpc. The median length, width,
excitation temperature, line mass of these filaments are ⇠ 49 pc, ⇠ 2.9 pc, ⇠ 8.9 K, and ⇠ 39 M pc 1 ,respectively. The velocity structures along these filaments exhibit oscillatory patterns, which are likely
caused by the fragmentation or accretion process along these filaments. The maximum accretion rateis estimated to be as high as ⇠ 700 M pc 1 . A total of ⇠ 162 13 CO clumps and ⇠ 103 young stellar
objects (YSOs) are identified in this region. Most of the clumps are in gravitationally bound states.
Three H ii regions (G154.359+2.606, SH2-211, SH2-212) are found to be located in the apexes of thefilaments. Intense star forming activities are found along the entire filamentary cloud. The observed
results may help us to better understand the link between filaments and massive star formation.

Keywords: Molecular clouds(1072) — Interstellar filaments(842) — Star forming regions(1565) —Young stellar objects(1834)