SF newsletter 343 #19-23
M. Momose (Ibaraki Univ.)#19-22のポイントを紹介＃23 は アブストラクトと一言メモ19.

should decrease again for M dwarfs near the sub-stellar boundary like Trappist-1.

Small Protoplanetary Disks in the Orion Nebula Cluster and OMC1 with ALMA

Justin Otter, Adam Ginsburg, Nicholas P. Ballering, John Bally, J. A. Eisner, Ciriaco Goddi, Richard Plambeck,
Melvyn Wright F The Orion Nebula Cluster (ONC) is the nearest dense star-forming region at ⇠400 pc away, making it an
ideal target to study the impact of high stellar density and proximity to massive stars (the Trapezium) on protoplanetary disk
evolution. The OMC1 molecular cloud is a region of high extinction situated behind the Trapezium in which actively formingstars are shielded from the Trapezium’s strong radiation. In this work, we survey disks at high resolution with ALMA at three
wavelengths with resolutions of 0.095 arcsec (3 mm; Band 3), 0.048 arcsec (1.3 mm; Band 6), and 0.030 arcsec (0.85 mm;
Band 7) centered on radio Source I. We detect 127 sources, including 15 new sources that have not previously been detected
at any wavelength. 72 sources are spatially resolved at 3 mm, with sizes from ⇠8 - 100 AU. We classify 76 infrared-detectedsources as foreground ONC disks and the remainder as embedded OMC1 disks. The two samples have similar disk sizes, but the
OMC1 sources have a dense and centrally concentrated spatial distribution, indicating they may constitute a spatially distinct
subcluster. We find smaller disk sizes and a lack of large (>75 AU) disks in both our samples compared to other nearby starforming regions, indicating that environmental disk truncation processes are significant. While photoevaporation from nearby
massive Trapezium stars may account for the smaller disks in the ONC, the embedded sources in OMC1 are hidden from thisradiation and thus must truncated by some other mechanism, possibly dynamical truncation or accretion-driven contraction.

7• Trapezium
cluster背後も含めた若い星々に付随するダスト円盤サーベイ
V899 Mon: a peculiar eruptive young star close to the end of its outburst

Sunkyung Park, Ágnes Kóspál, Fernando Cruz-Sáenz de Miera, Michał Siwak, Marek Dróżdż, Bernadett• ALMA
B3 Daniel(0.095”38au),
B6Könyves-Tóth,(0.048”), B7Levente(0.030”)，source
Iを視野中心に
Ignácz,T. =Jaﬀe,
Réka
Kriskovics, Jae-Joon
Lee, Jeong-Eun Lee, Gregory N.
Mace, Waldemar Ogłoza, András Pál, Stephen B. Potter, Zsófia Marianna Szabó, Ramotholo Sefako, Hannah• 127天体で検出(うち15天体が新検出)，λ=3mmで72天体が空間分解された
L. Worters F V899 Mon is an eruptive young star showing characteristics of both FUors and EXors. It reached a peakbrightness in 2010, then briefly faded in 2011, followed by a second outburst. We conducted multi-filter optical photometric• 検出天体のうち76天体は近赤外線で検出済み=
手前にある天体群
monitoring, as well as optical and near-infrared spectroscopic observations of V899 Mon. The light curves and color-magnitude
diagrams show that V899 Mon has been gradually fading after its second outburst peak in 2018, but smaller accretion bursts• 残りは分子雲成分に埋もれたもの，より空間的に集中しているが，円盤性質に顕著な差はない
are still happening. Our spectroscopic observations taken with Gemini/IGRINS and VLT/MUSE show a number of emission
lines, unlike during the outbursting stage. We used the emission line fluxes to estimate the accretion rate and found that it has• 他の領域に比べると，円盤サイズ小さめ。半径>75auのものが欠けている
significantly decreased compared to the outbursting stage. The mass loss rate is also weakening. Our 2D spectro-astrometric• Trapezium
clusteranalysis of emission
linesによるPhotoevaporation＝手前側のものは説明可能
recovered jet and disk emission of V899 Mon. We found the emission from permitted metallic lines andthe CO bandheads can be modeled well with a disk in Keplerian rotation, which also gives a tight constraint for the dynamical• 埋もれている円盤も半径が小さいのは，他の効果（例えばdynamical
truncation）も考える必要あり
stellar mass of 2 M . After a discussion of the physical changes that led to the changes in
the observed properties of V899
Mon, we suggest this object is finishing its second outburst.

観測の概要
Small disks in Orion
Band

3
6
7

CentralFrequency

Primary
Beam
FWHM

Beam
Major⇥Minor

GHz

00

00

98.0
223.5
339.8

59.4
26.1
17.1

Beam
PA

⇥ 00

0.097 ⇥ 0.071
0.049 ⇥ 0.037
0.030 ⇥ 0.024

43.4
80.4
-69.6

RMS
(central-edge)

Min. Mdisk a(central-edge)

mJy beam 1

M0.04 - 1.0
0.2 - 20.12 - 0.45

2.3 - 57
5 - 502.0 - 7.4

3
Flux
RobustThresholdb
mJybeam 10.5
0.5
0.5

Baseline
Lengths
m
(k )

0.05

21-14854(6.9-4855)

1

150-12147(112-9055)

0.05

210-1124(238-1273)

Table 1. Imaging parameters and information for each image. a Minimum detectable disk dust mass, calculated according toEquation 2, see Section 3.9. b the flux level in the residual image to end tclean.

disk sizes and masses to those in other regions. Lastly, in Section 6, we discuss di↵erent disk truncation mechanisms
in the context of our results in Orion.
2. OBSERVATIONS AND METHODSThis work presents single pointing ALMA continuum observations of the Orion Nebula Cluster in bands 3, 6, and 7,

4

Otter and Ginsburg et. al.

観測領域○が今回検出された天体観測波長でPrimary beam
(視野の広さ)が異なることに注意
緑 B3 only
ピンク B3/B6
赤 B3/B6/B7
Figure 1. An IR image mosaic from Gemini of the ONC, with our detected sources circled (Bally et al. 2015). The dottedshapes show the fields of view of our data in di↵erent bands, 3 mm, 1.3 mm, and 0.85 mm in descending size. Sources circled
in red are detected at 0.85 mm, pink at 1.3 mm, and green at 3 mm.

source has a peak signal-to-noise ratio S/N

5 in one or more bands to remain in our final catalog. We estimate theBand7で捉えられた円盤の全イメージ
Small disks in Orion

6

5

Otter and Ginsburg et. al.

Small disks in Orion

Figure2. (3/3)
Imagessources in the band 7 (0.85 mm) field of view, from left to right: band 3 (3 mm), band
Figure2. (2/3)
Imagessources in the band 7 (0.85 mm) field of view, from left to right: band
3 (3 mm),band
6 (1.3ofmm),
2. (1/3) Images of sources in the band 7 (0.85 mm) field of view, from left to right: band
3 (3 mm),band
6 (1.3ofmm),
00
00
and
7. The(200yellowband
7. Theyellow
is band
0.5 acrossAU).ellipses
The shows the synthesized beam size for each image. Each cutout is 0.500 across (20d 7. The yellow ellipses shows the synthesized beam size for each image. Each cutoutand
is 0.5across(200
AU).ellipses
The shows the synthesized beam size for each image. Each cutout
number in
the topNondetections
left is the source ID, and source numbers boxed in red correspond to newly detected sources. N
numbersourcesin the (seetop leftis thesource ID, and source numbers boxed in red correspond to newly detectedsources.
n the top left is the source ID, and source numbers boxed in red correspond to newly detected
Section3.5).
are shown as well.
are shown as well.
ctions are shown as well.

この他，B6/3で捉えられている天体，B3のみで捉えられている天体がある

We compute
thetoRMS
following the method in Section 3, in an annulus of radius 0.100 with a width of 15
approximately 12 mas yr 1 (e.g. Dzib et al. 2017; Goddi et al. 2011; Rodriguez et al. 2005),
amountinga distance3.3. Correspondence with IR sources: the ONC and OMC1 samples
00
include these upper limits in Table 3.3.
of
0.2
thisthosetime inperiod,parate the sources into two groups: OMC1, those in or behind the molecular cloud, andoverONC,
front indicating that proper motions are significant00for the fastest-moving sources. Our results
are J,
notto thecatalog
choice of matching radius: a radius as small as 0.3 yields 71 source matches, and a radius asciated with the nebula cluster. We match our source locations to the near-IR (NIR)
H,sensitiveandKs band00
3.4. Correspondence with other wavelength dataonly
as 1 0.7
00 adds 2 more source matches with 78 total matches.
of the ALMA
ch et al. (2002) (hereafter MLLA), where we consider a NIR source a match if itlarge
is within統計的性質：disk mass & radius
Small disks in Orion

21

Figure 15. Left: log disk dust mass as a function of projected distance to ✓1 Ori C, for our ONC (blue) and OMC1 (orange)
オレンジがOMC(クラウド中)，青はOrionsamples, with Spearman
correlations of -0.092 and 0.020 for the ONC and OMC1 samples Nebularespectively. Cluster
Right: The same but weAndrews et al. (2018)との比較
22

Otter and Ginsburg et. al.

ピンクは，SMAデータ＋ALMAのLupusデータ(Ansdell et al. 2016他)から求めた
Nearby protoplanetary disks（ただし遷移円盤は除く?）のデータ，及びL*=1L⦿に対する
サイズ・ルミノシティ関係−1/4(2−q) 1/(2−q)
Reﬀ ∝ L*
Lmm

Figure 16. A plot of disk HWHM versus scaled millimeter luminosity. Green points are our band 3 measurements, and pink
points are data from Andrews et al. (2018). Triangles represent upper limits. The orange dashed line shows the radius andluminosity scaling relation derived by A18 for disks with solar luminosity protostars. The blue solid line is a linear fit to ourband 3 measurements, with sources I and BN removed as outliers. This fit has a slope of 0.17 ± 0.05 AU/Jy and correlation
r2 = 0.15.

20.

radiation and thus must truncated by some other mechanism, possibly dynamical truncation or accretion-driven contraction.

V899 Mon: a peculiar eruptive young star close to the end of its outburst

Sunkyung Park, Ágnes Kóspál, Fernando Cruz-Sáenz de Miera, Michał Siwak, Marek Dróżdż, Bernadett
Ignácz, Daniel T. Jaﬀe, Réka Könyves-Tóth, Levente Kriskovics, Jae-Joon Lee, Jeong-Eun Lee, Gregory N.
Mace, Waldemar Ogłoza, András Pál, Stephen B. Potter, Zsófia Marianna Szabó, Ramotholo Sefako, Hannah
L. Worters F V899 Mon is an eruptive young star showing characteristics of both FUors and EXors. It reached a peakbrightness in 2010, then briefly faded in 2011, followed by a second outburst. We conducted multi-filter optical photometricmonitoring, as well as optical and near-infrared spectroscopic observations of V899 Mon. The light curves and color-magnitude
diagrams show that V899 Mon has been gradually fading after its second outburst peak in 2018, but smaller accretion bursts
are still happening. Our spectroscopic observations taken with Gemini/IGRINS and VLT/MUSE show a number of emission
lines, unlike during the outbursting stage. We used the emission line fluxes to estimate the accretion rate and found that it has
significantly decreased compared to the outbursting stage. The mass loss rate is also weakening. Our 2D spectro-astrometricanalysis of emission lines recovered jet and disk emission of V899 Mon. We found the emission from permitted metallic lines andthe CO bandheads can be modeled well with a disk in Keplerian rotation, which also gives a tight constraint for the dynamicalstellar mass of 2 M . After a discussion of the physical changes that led to the changes in the observed properties of V899
Mon, we suggest this object is finishing its second outburst.

High- and Low-Mass
Star Formation: Observation-Based Continuum分光観測
Modelling• FULinking
Ori /Ex Oriの特徴を示す
V899 Monのマルチフィルタモニタリング観測＋Optical/IR
and Physical Conditions
R. L. Pitts, L. E.&Kristensen,
J. K. Jørgensen, S. J. van der Walt F Astronomers have yet to establish whether high(Gemini/IGRINSVLT/MUSE)
mass
protostars
form
from
high-mass
prestellar
cores,
similarto
theirlower-masscounterparts,
orfromlower-massfragments
at• 2018年のsecond burst以降，ゆっくりと減光
the heart of a pre-protostellar cluster undergoing large-scale collapse. Part of the uncertainty is due to a shortage of envelope
structure data on protostars of a few tens of solar masses, where we expect to see a transition from intermediate-mass star
• 分光観測の結果・・・
formation to the high-mass process. We sought to derive the masses, luminosities, and envelope density profiles for eight sources• 質量降着率はアウトバースト時に比べて著しく減少（いくつかのline
in Cygnus-X, whose mass estimates in the literature placed them in the sampling gap. fluxesから）
Combining these sources with similarly
evolved sources in the literature enabled us to perform a meta-analysis of protostellar envelope parameters over six decades in≈ 2M⦿
• Metalの許容線やCOバンドヘッドは，ケプラー回転する降着円盤でよく説明，M*
source luminosity. We performed spectral energy distribution (SED) fitting on archival broadband photometric continuum datafrom 1.2 to 850 µm, to derive bolometric luminosities for our eight sources plus initial mass and radius estimates for modelling時間変化の様相(1)Park et al.

Figure 1. Top: light curve of V899 Mon. Bottom: light curve of V899 Mon for the 2020 September – 2021 April tim(pink box in the top panel). The bottom figure shows our own observations from Table 5 (Appendix A), complemendata from the ASAS-SN survey. The photometric uncertainty smaller than the symbol size is not presented. The num
brackets indicate the central wavelength of each filter in nm. The photometric data from the ASAS-SN (V and g; Shapp2014; Kochanek et al. 2017), ZTF (g and r; Masci et al. 2018), WISE (W1 and W1; Mainzer et al. 2011), 2MASS (JCutri et al. 2003), and U, B, V, R, I, J, H, and Ks from Ninan et al. (2015) were used. Circle and star symbols indicate
data and our data, respectively. Di↵erent colors indicate di↵erent bandpasses. The black dotted and dashed lines indi
spectroscopic observation date of IGRINS and MUSE. The dark gray dotted line in the top panel indicates the observatof Ninan et al. (2015). Dashed gray line in the top panel indicates the data point used in the SED in Figure 7.

ピンク部分の拡大

時間変化の様相(2)this paper accretion bursts and accompanying central dips, are indicatedsubplots show zoom-in accretion bursts and central dips in 2018 and 2021

周期解析 in 2020/2021

40
80
35
70
30

8

V & gバンド

combined light curvePark et al.

Period [d]

60

25

50

20

40
30

15

20

10

10

5

0

0
9100

9150

9200
HJD-2450000

9250

9300

Figurepresent
Figure 4. WWZ spectrum calculated for 2020/2021 season.
of thes
The colors represent the Z-statistic values. Edge e↵ects are
10
Parkdashedet alcontainedbeyond the white lines.
dots.

former line shows a relatively stronger blue-shifted absorption component to the latter. H↵ shows a shallow blue-shifted absorption component (> –332 km s 1
where the intensity is minimum) and a strong emission
profile (⇠ 60 km s 1 ) in our observation (black). The
blue-shifted absorption component is produced by an
outflowing wind (Hartmann & Kenyon 1996; Hartmann2009), and typically FUors show strongly blue-shifted

SED

Figure 3. V - and g-band combined light curve of V899 Mon for the 2014–2021 time period. The two, more widely studied inthis paper accretion bursts and accompanying central dips, are indicated by red and blue colors, respectively. Left and right
subplots show zoom-in accretion bursts and central dips in 2018 and 2021, respectively.

40

Th
absorpori
wasetdia
Novem
200
shiftedHo
the 201ou
also oT(2016)
line
cal
&P
lar
and
Ha
202
line
rela
Th
acc

分光観測解析の例: CO overtone(Δv＝2)

Figure 18. In the top panel with the black line, we show the CO o
The best-fitting model with Tex = 2482 K and NCO = 1022.18 cm 2 isTable 3. Disk parameters and dynamic

Parameteri( )Rin (au)Rout (au)M⇤ (M )↵

CO v = 3 1
52.54 ± 9.61
0.62 ± 0.16
3.55 ± 1.04
2.08 ± 0.57
2.00 (fixed)

v=2-0とv=3-1のPark et al.
スペクトル全体のフィット

18

CO v = 2 0
51.01 ± 10.82
0.55 ± 0.20
4.17 ± 1.34
2.00 ± 0.67
2.00 (fixed)

平均をとったもの
1.2

1.0

18

p
ra
of
st
th
et

CO v = 2 − 0 R−branch lines
CO v = 3 − 1 R−branch lines
Disk model for CO

Park et al.
Normalized flux

0.8

0.6

0.4

0.2

0.0
−100

−50

0
Velocity (km s−1)

50

100

Figure 19. Average line profiles of the v = 2 0 and v =Figure 18. In the top panel with the black line, we show the CO overtone bandhead emission in the spectrum of V899 Mon.
3 1 rovibrational lines of V899 Mon. The disk model thatfits the blue-shifted side of the line profiles is indicated with
The best-fitting model with Tex = 2482 K and NCO = 1022.18 cm 2 is plotted in red, while the lower panel shows the residuals.
the dashed curve. The corresponding model parameters are
Figure 18. In the top panel with the black line, we show the CO overtonein bandhead
Table 3. emission in the spectrum of V899 Mon.

Parameteri( )Rin (au)Rout (au)M⇤ (M )↵The best-fitting model with Tex = 2482 K and NCO = 1022.18 cm 2 is plotted in red, while the lower panel shows the residuals.
lerian peaks are asymmetric, but also the line is broader
Table 3. Disk parameters and dynamical stellar mass for V899 Mon.
in the red side and narrower in the blue side. Therefore,
Table 3. Disk parameters and dynamical stellar mass for V899 Mon.
we fitted the blue-shifted ( 60 km s 1 < v < 0 km s 1 )CO v = 2 0 CO v = 3 1 Average CO Atomic (blue-shifted) Atomic (red-shifted)
red-shifted(0 km s 1 < Atomic
v < 80 kms 1 ) sides of theParameter CO v = 2 0 CO v = 3 1 Average CO andAtomic(blue-shifted)(red-shifted)

51.01 ± 10.82
0.55 ± 0.20
4.17 ± 1.34
2.00 ± 0.67
2.00 (fixed)

52.54 ± 9.61
0.62 ± 0.16
3.55 ± 1.04
2.08 ± 0.57
2.00 (fixed)

51.77 ± 10.22
0.58 ± 0.18
3.86 ± 1.19
2.04 ± 0.42
2.00 (fixed)

求めたディスクパラメータ51.77 (fixed)i( )Rin (au)
0.41 ± 0.28
Rout (au)4.46 ± 1.94
M⇤ (M )
2.04 (fixed)↵
2.00 (fixed)

51.01 ±51.77
10.82 (fixed)
52.54 ± 9.61
0.55 ±0.20
0.20 ± 0.62
0.15± 0.16
4.17 ± 1.34
3.73 ± 3.55
1.86± 1.04
2.00 ± 0.67
2.08 ± 0.57
2.04 (fixed)
2.00 (fixed)
2.00 (fixed)

2.00 (fixed)

51.77 ± 10.22
0.58 ± 0.18
3.86 ± 1.19
2.04 ± 0.42
2.00 (fixed)

51.77 (fixed)
0.41 ± 0.28
4.46 ± 1.94
2.04 (fixed)
2.00 (fixed)

Avera
51.77 ±
0.58 ±
3.86 ±
2.04 ±
2.00

51.77 (fixed)
0.20 ± 0.15
3.73 ± 1.86
2.04 (fixed)
2.00 (fixed)

em
th
an
of
ca
li
th
on
d
w
ra
as
ga
st
n
if
th
fi
as

21.

stellar mass of 2 M . After a discussion of the physical changes that led to the changes in the observed properties of V899
Mon, we suggest this object is finishing its second outburst.

Linking High- and Low-Mass Star Formation: Observation-Based Continuum Modellingand Physical Conditions
R. L. Pitts, L. E. Kristensen, J. K. Jørgensen, S. J. van der Walt F Astronomers have yet to establish whether highmass protostars form from high-mass prestellar cores, similar to their lower-mass counterparts, or from lower-mass fragments at
the heart of a pre-protostellar cluster undergoing large-scale collapse. Part of the uncertainty is due to a shortage of envelope
structure data on protostars of a few tens of solar masses, where we expect to see a transition from intermediate-mass starformation to the high-mass process. We sought to derive the masses, luminosities, and envelope density profiles for eight sourcesin Cygnus-X, whose mass estimates in the literature placed them in the sampling gap. Combining these sources with similarly
evolved sources in the literature enabled us to perform a meta-analysis of protostellar envelope parameters over six decades insource luminosity. We performed spectral energy distribution (SED) fitting on archival broadband photometric continuum datafrom 1.2 to 850 µm, to derive bolometric luminosities for our eight sources plus initial mass and radius estimates for modelling
density and temperature profiles with the radiative transfer package Transphere. The envelope masses, densities at 1000 AU,
outer envelope radii, and density power law indices as functions of bolometric luminosity all follow established trends in theliterature spanning six decades in luminosity. Most of our sources occupy an intermediate to moderately high range of masses
and luminosities, which helps to more firmly establish the continuity between low- and high-mass star formation mechanisms.
Our density power law indices are consistent with observed values in literature, which show no discernible trends with luminosity.
Finally, we show that the trends in all of the envelope parameters for high-mass protostars are statistically indistinguishable
from trends in the same variables for low- and intermediate-mass protostars.

（注）star formation newsletterとarXivのアブストラクトの文章は違っています（内容的には大きな矛
Revisiting the Cygnus OB associations

盾はないように見えますが）。
Alexis L. Quintana, Nicholas J. Wright F OB associations play an important role in Galactic evolution, though their

origins and dynamics remain poorly studied, with only a small number of systems analysed in detail. In this paper we revisit• ブロードバンドSED及び，ダスト連続波のアーカイブデータから出した輝度分布(HerschelとかSCUBAthe existence and membership of the Cygnus OB associations. We find that of the historical OB associations only Cyg OB2
and OB3 stand out as real groups. We search for new OB stars using a combination of photometry, astrometry, evolutionary2とか)のフィッティングを球対称コアのradiative
packagemodels and an SED fitting process, identifying 4680 probabletransfer
OB stars witha reliability(Transphere)で行い，物理構造導出
of >90%. From this sample we search
for OB associations using a new and flexible clustering technique, identifying 6 new OB associations. Two of these are similar• 各種物理量（envelope
at 1000au，外径，power-law
index）は，bolometricluminosityに
to the associations Cygmass,
OB2 andρOB3,
though the others bear no relationship toany existing systems. We characterizetheproperties of the new associations, including their velocity dispersions and total stellar masses, all of which are consistent with
沿って滑らかに変化
つまりlow-massとhigh-mass
starformationは滑らかに繋がる（明確な違い
typical values for OB→
associations.
We search for evidenceof expansion
and findthat all are expanding, albeit anistropically,
with stronger and more significant expansion in the direction of Galactic longitude. We also identify two large-scale (160 pc andはない）
25 km s 1 ) kinematic expansion patterns across the Cygnus region, each including three of our new associations, and attribute
this to the eﬀects of feedback from a previous generation of stars. This work highlights the need to revisit the existence andサンプル：PILS-Cygnus sourcesのSMA画像
A&A proofs: manuscript no. PILSCygX-cont

N12

N30

N38

7"~10000 AU

7"~10000 AU

7"~10000 AU

N48

N51

N53

N54

7"~10000 AU

7"~10000 AU

N63

S8

S26

7"~10000 AU

7"~10000 AU

3"~10000 AU

7"~10000 AU

Fig. 1: SMA 868 µm (345 GHz) images of all 10 PILS-Cygnus sources (N53 and N54 are both featured in the middle right panel)
contoured in red at 10%, 40%, and 70% of the maximum flux. The cyan ellipses in the upper right corners show the FWHMa ofboth axes and position angle of the synthesised SMA beam, and the line segments alone the bottoms of each panel show how long
10000 AU and its corresponding angular length appears at the adopted distance to each object. We note the change of scale for S26.
All images are normalised to their peak flux densities, which are given in Table 1.
was preferable at short wavelengths, and at long wavelengths,

pected to contribute more flux to N53 than is accounted for by

SEDやradial profile fitの例

Fig. 3: Full SEDs for all sources, which were fitted, interpolated, and integrated by trapezoidal sums to calculate the bolometricluminosities. The curves running through data at
70 µm are the FIR/submm component of the SED fits discussed in §2.3, while
the rest of the data has been log-linearly interpolated in lieu of integrating over the less-physical warmer SED components. For
N30 and N51, no SED components passed through the 70 µm data point, so the SED fit was only used for
160 µm while the70 µm data were linearly interpolated like the shorter wavelength data. The light red filled curves show the margins of uncertainty
as constrained by the flux errors, and were likewise integrated to estimate the total uncertainty in luminosity.

SED ﬁtting

103
102

4

0.200

11

model, p=-1.0
N30 data

CygX-N30
p = 1.0±0.1
G/G0 =1

0.175

1

100
10 1
10 2
SED (Transphere)
SED (greybody)
N30 data
N30 data (greybody only)

10 3
10 4
100

101

102

103

[µm]

Fig. 4: SEDs of N30. The dashed red line is the modified greybody fit to the data, shown as black and grey markers with error
bars. The solid blue line is the best fit Transphere SED to only
the data shown in black. The grey points at < 70 µm were excluded from the fit because the local ISRF is poorly constrained,
and we were generally unable to set it large enough to fit the
short-wavelength data without crashing Transphere.

0.150

10

9

3

0.125

8

T (r) [K]F [Jy]

10

law index is related to the dust emissivity index via q = 2/(4 + )(Chandler et al. 1998; Kenyon et al. 1993).
10
10
3We had convolved the SCUBA legacy data to the radius of10
the secondary or “error” beam, which has a FWHM 4000 , andthen attempted the same Gaussian deconvolution routine as with
10
10
the Herschel data. Because these data were much more aggres10
sively background-subtracted than Herschel data, we allowed thefitted flat background to accommodate negative average back10
10 theground levels so we could add it to the peak amplitude of
10
Gaussian fit. We still integrated over a Gaussian with zero con10
stant background level. Nevertheless, we know that negative average background levels are more of an issue at 450 µm10 than
10
10
10
10
10
850 µm, so data at the shorter
wavelength
may
have
still
been
Distance from peak [AU]R [AU] at dist=1.4 kpc
boosted too much despite removing
the flat background from in(a) A.
(b) B.
tegration. That makes it less clear whether from the SCUBA
5: A. Radial
emissionreliable
profile comparison
betweenthe data(red dashedline) and the Transphere model (blue solid line), bothdata aloneFig.
would
be
more
than
the
valuesofwedenormalised by their respective maxima. B. Resulting model temperature (red solid line and red axis) and density (blue dash-dotted
rived, whicharebluedominatedline andaxis) profiles.by Herschel data. For most objects
that were isolated or separable, the results seemed reasonableture identified
in MBS07 as S7.of
Unfortunately,PSF fitting
the compared
After correcting for their older distance measurements, we
to within the
50% uncertainty
the 450 µm
dataofas
data was never an option since, at the native SCUBA found that 6 out of 8 sources had at least one mass measurement
to the 350SCUBA
and
500µm Herschel/SPIRE data. We also checked
and SCUBA-2 resolutions, all of the objects are resolved along consistent with the respective mass estimates in MBS07. For
the resultsatagainst
simple aperture photometry with SAO
DS9least one axis.
N30 and
N48, neither Mc nor Menv were as large as the envelope
computed by MBS07, though the discrepancy with N30
Given andthat SCUBA-2
the Transpheremodelsincorporate
more masses
for both SCUBA
data,
and
with
Herschel
imagesphysics, the masses from those models are probably more ac- was small. The rest of the sources seem to agree more closely
where the curate
clumpscould
befromisolated.
For S8alone,
we redidwith thethe
MBS07 masses when they have lower T dust , since thethan themassesmulti-component
greybody
fitting,
we found were typically higher than MBS07 assome components
not thenecessarily
physical.
The lu- temperatures
photometryof which
manually,
cuttingareout
brightspiral-arm-likefeasume. All of this makes sense given the methods they used to
0.100

0.075

7

2

6

0.050

5

0.025

1

0.000

0

5000

10000

15000

20000

25000

30000

minosities could only be fully accounted for by spline-fitting the

4

1

2

3

4

n(r) [cm 3 ]Menv /M =190±50
L /L =14000±3000
G/G0 =1

Normalized flux (integral = 1)

10

左が輝度分布のフィットと実際の観測の比較
Beltrán et al. 2002). Then in the optically thin Rayleigh-Jeans
右が温度分布と密度分布のプロット
limit, it can be shown analyticallythat the temperature powerR. L. Pitts, L. E. Kristensen, J. K. Jørgensen and S. J. van der Walt: Linking high- and low-mass star formation

4

R. L. Pitts, L. E. Kristensen, J. K. Jørgensen and S. J. van der Walt: Linking high- and low-mass star formation

主要な結果

10

5

single power law
broken power law

single power law
broken power law

10

105

3

Rout [AU]Menv [M ]

104

102
101

104

100
10 1

10 1

101

102

103

104

105

106

100

101

102

109

p = 1.4
single power law
p = 0.4|p| (power-law index)

2.0

103

104

105

106

104

105

106

Lbol [L ]

1.5

1.0

ちょうどhigh-massと10 1

Lbol [L ]2.5

それ以外のシンボルは
過去の研究

100

n(r=1000 AU) (cm 3 )黒の大きな丸が今回の結果

10 2

single power law
broken power law

108

107

106

0.5

Low-massを繋ぐところに
ハマった

105

0.0
10 1

100

101

102

103

Lbol [L ]

104

105

106

10 1

100

101

102

103

Lbol [L ]Fig. 6: Comparison of PILS-Cygnus Sources (large black circles with error bars) to objects in the literature from van der Tak et al.
2000 (purple x’s), Jørgensen et al. 2002 (magenta upward-pointing triangles), Mueller et al. 2002 (thin blue diamonds), Shirley
et al. 2002 (maroon stars), Hatchell & van der Tak 2003 (green squares), Williams et al. 2005 (small orange circles), Crimier et al.
2009, 2010a,b (wide navy diamonds), Kristensen et al. 2012 (cyan downward-pointing triangles), van der Tak et al. 2013 (red plus
signs), and Cao et al. 2019 (pale gold error bars). Solid and dashed black lines represent single and broken power law fits to thedata, respectively.

22.

Our density power law indices are consistent with observed values in literature, which show no discernible trends with luminosity.
Finally, we show that the trends in all of the envelope parameters for high-mass protostars are statistically indistinguishable
from trends in the same variables for low- and intermediate-mass protostars.

Revisiting the Cygnus OB associations

Alexis L. Quintana, Nicholas J. Wright F OB associations play an important role in Galactic evolution, though their
origins and dynamics remain poorly studied, with only a small number of systems analysed in detail. In this paper we revisit
the existence and membership of the Cygnus OB associations. We find that of the historical OB associations only Cyg OB2
and OB3 stand out as real groups. We search for new OB stars using a combination of photometry, astrometry, evolutionarymodels and an SED fitting process, identifying 4680 probable OB stars with a reliability of >90%. From this sample we search
for OB associations using a new and flexible clustering technique, identifying 6 new OB associations. Two of these are similar
to the associations Cyg OB2 and OB3, though the others bear no relationship to any existing systems. We characterize theproperties of the new associations, including their velocity dispersions and total stellar masses, all of which are consistent withtypical values for OB associations. We search for evidence of expansion and find that all are expanding, albeit anistropically,
with stronger and more significant expansion in the direction of Galactic longitude. We also identify two large-scale (160 pc and
25 km s 1 ) kinematic expansion patterns across the Cygnus region, each including three of our new associations, and attribute
this to the eﬀects of feedback from a previous generation of stars. This work highlights the need to revisit the existence andmembership of the historical OB associations, if they are to be used to study their properties and dynamics.

• Cyg OB associations の分類をモダンな方法（photometry + astrometry+evolutionary model+ SED fitting）を使って見直した• 歴史的に同定されていたもののうち，Cyg OB2とOB3しかリアルなグループでないと分かった• 新しく定義したassociationsの性質を調べ，全て非等方にexpandingしていることを確認
8• その他，2つのlarge scale (160pc & 25km/s)の膨張パターンも新たに確認：feedbackの影響？

in Cygnus remain poorly studied despite the available information(see e.g. Reipurth & Schneider 2008; Comerón & Pasquali 2012;
Mel’nik & Dambis 2017; Melnik & Dambis 2020) and deserve asa consequence a better understanding.
In this paper we investigate the 5 main OB associations in
Cygnus, consider their reality as genuine associations, and eventually propose the existence of 6 alternative structures in the region,
Revisiting
the Cygnus
associations
3
some of which
overlapwith theOB
historicalassociations.
In Section
2 we analyse the historical OB associations and show that, apartfrom Cyg OB2 and OB3, they do not stand out as real groups. In
Section 3 we describe how we identify OB stars across the region
using photometry, astrometry, evolutionary models and a new SED
fitting tool. In Section 4 we use a new flexible clustering techniqueto identify new associations. In Section 5 we characterize the identified groups by studying their physical and kinematic properties. In
Section 6 we discuss these findings and in Section 7, we wrap up
our results.

OB associationsの分布：Blaha &Humphreys (1989) + Behlanas et al. (2019)2

THE EXISTING OB ASSOCIATIONS IN CYGNUS

In this section we consider the stellar content and reality of theexisting OB associations in Cygnus using modern photometry andastrometry.
Our focus is on a 60 deg2 area in the Cygnus constellation
spanned by l = [71 ,81 ] and b = [-1 , 5 ]. This region contains
the OB associations Cyg OB1, OB2, OB3, OB8 and OB9, the five
most well-studied associations in Cygnus. A total of 178 OB stars
were identified by Blaha & Humphreys (1989) and divided between
these five associations based only on their positions on the sky. The
other OB associations in Cygnus, i.e. Cyg OB4, OB5, OB6 and OB7
are not considered here because they are much more extended andfound at greater longitudes. For this work we use the list of 164 OB
stars in Cyg OB1, OB3, OB8 and OB9 from Blaha & Humphreys(1989). For Cyg OB2 we use the updated census of OB stars from

Figure 1. Top panel: spatial distribution of Cygnus OB association members from Blaha & Humphreys (1989) and Berlanas et al. (2019), displayed with their

Cygnus, starting from the ground up.

3

IDENTIFYING AND CHARACTERISING OB STARS

In this section we outline our photometric SED fitting techniqueused to identify OB stars across our region of study and estimate
their physical parameters.
3.1

最新データを再解析

Data and selection process

We use photometry and astrometry from IGAPS (INT Galactic
Plane Survey, providing 6, A, 8 band photometry, Drew et al. 2005,
Monguió et al. 2020), 2MASS (2 Micron All Sky Survey, providing, , and B band photometry, Cutri et al. 2003), UKIDSS (United
Kingdom Infrared Deep Sky Survey, providing deeper ,andband photometry, Lucas et al. 2008), and Gaia EDR3 (Gaia
Collaboration et al. 2016, Gaia Collaboration et al. 2020, Rielloet al. 2020), providing ⌧, ⌧ ⌫% and ⌧ '% band photometry. In
addition we use parallaxes and proper motions from Gaia EDR3
(Fabricius et al. 2020), and the probabilistic distances derived fromthese parallaxes provided by Bailer-Jones et al. (2020). We correct
all parallaxes for the non-zero parallax zero-point, whose value
has been shown to depend on magnitude, ecliptic latitude, and thepseudocolour of the source, using the prescription in Lindegren
et al. (2020a).
There are a total of 8,474,415 sources in Gaia EDR3 in this region of study, 7,434,799 of which have a valid five or six-parameter
astrometric solution (suitable for our needs). We crossmatch thissample with 2MASS, UKIDSS and IGAPS, using a matching radius of 1", using the proper motions to account for the epoch difference. We found near-IR matches for 6,347,730 sources and IGAPS
matches for 6,158,261 sources. We discarded the 1,087,069 Gaia
sources that lacked a near-IR counterpart as a near-IR photometricdatapoint is necessary for our SED fitting. These sources generally

注意深い解析の末，4680個のOB型星(reliability > 90%)を選出

Revisiting the Cygnus OB associations

7

4680個の空間・固有運動分布
Revisiting the Cygnus OB associations

Figure 6. Comparison between the literature spectroscopic temperature withthat from our SED fits, for the 205 objects with reliable spectral types. Bluedashed lines correspond to the ’limit’ at ;>6 ()4 5 5 ) = 4 (dividing OB stars
from cooler stars). Green sources are from Blaha & Humphreys (1989) while
red ones are from Berlanas et al. (2019).

while their luminosities were under-estimated by ⇠0.01 dex. Bothof these effects will cause evolved stars with low surface gravities tobe fitted closed to the ZAMS than they should be, under-estimating
their ages.

Figure 6. Comparison between the literature spectroscopic temperature withthat from our SED fits, for the 205 objects with reliable spectral types. Bluedashed lines correspond to the ’limit’ at ;>6 ()4 5 5 ) = 4 (dividing OB stars
4from
IDENTIFYINGNEW
OBareASSOCIATIONS
cooler stars). Greensourcesfrom Blaha & Humphreys (1989) whilered
ones
arewithfromal. sample(2019). of OB stars to search for OB
pectroscopicIn
temperature
this
section
weBerlanas
exploit etour

h reliable spectral
types. Blueassociationsin the Cygnus region. For this purpose, we will use a
OB stars)4 5 5 ) = 4 (dividing
newmethod,the results of
whichwill be comparedwithdex.
existing
while
their
luminosities
wereunder-estimated
by
⇠0.01
Both
ha & Humphreys (1989) whileclustering
algorithms.
of these effects
will cause evolved stars with low surface gravities tobe fitted closed to the ZAMS than they should be, under-estimating
pectroscopic temperature
theirBroadages. with
4.1
kinematics distribution of OB stars
matedbyspectral⇠0.01 types.
dex. Bothh reliableBlue) = surface
4 (dividing
OB stars)ith
tobroad spatial and kinematic distribution of our SED4 5 5low
Fig. 7gravities
shows
the
ha
& Humphreys(1989)
while
hould
be, under-estimating
fittedOBstars(;>6()4 NEW
4). Thespatial distribution (top panel)
5 5 ) > OB
4 IDENTIFYINGASSOCIATIONS

7

新たな
同定結果
6つのグループ
Figure 10. Spatial distribution of all the identified new OB groups (coloured symbols), plotted against the field OB stars (black dots).

新たな
同定結果
6つのグループ

Revisiting the Cygnus OB associations

11

Table 1. Properties of our new OB associations. The first column indicates the parameter, where the subscript ’m’ indicates the median value and ’f’ thedispersion.
Parameters

UnitsGroup AGroup B

Group CGroup DGroup EGroup F' ( ⇠ '() <deg

301.45304.37305.47304.34308.08302.95⇡⇢ ( ⇠ '() <

deg

35.68

41.43

37.87

37.64

41.30

36.58

;<

deg

72.61

78.58

76.11

75.44

80.19

74.04

1<

deg

2.06

3.31

0.54

1.19

0.85

1.44

3<

pc

1894.51726.31713.1

2000.11674.01985.2`;<mas yr 1

-6.90

-5.47

-6.57

-5.55

-4.72

-6.11

f`;

mas yr 1

0.24

0.34

0.27

0.16

0.27

0.30

`1 <

mas yr 1

-1.35

-0.59

-1.19

-1.34

-0.96

-1.33

f `1

mas yr 1

0.17

0.27

0.12

0.14

0.25

0.1

km s 1

2.20+0.55
0.47

2.72+0.47
0.37

2.24+0.46
0.35

1.53+0.45
0.41

2.15+0.66
0.45

2.80+0.58
0.48

112+33

82+32

76+23

75+23

110+43

139+33

fE;
fE1

km s 1

Observed number of B stars
Observed number of O stars

3+11

Corrected number of O stars
Estimated total stellar mass

1.55+0.32
0.24

13+43

"

Velocity gradient (l)mas yr 1 deg 1

Velocity gradient (l)km s 1 pc 1

Velocity gradient (b)mas yr 1 deg 1

Velocity gradient (b)km s 1 pc 1

Expansion age (l)MyrExpansion age (b)Myr

2344+275
252

2.18+0.41
0.33
4+11

0.96+0.22
0.18
10+12

11+43

1978+245
224

12+43

2165+245
224

1.32+0.30
0.26
4+21

11+43

1569+215
192

1.96+0.52
0.29
35+23
23+65

4198+354
319

1.50+0.30
0.25
5+21

17+45

2955+349
269

0.24 ± 0.07

0.39 ± 0.04

0.38 ± 0.04

0.20 ± 0.05

0.43 ± 0.06

0.35 ± 0.03

0.34 ± 0.04

0.03 ± 0.04

0.10 ± 0.04

0.07 ± 0.03

0.16 ± 0.05

0.06 ± 0.03

0.07 ± 0.02
0.09 ± 0.01

13.98 ± 3.99

10.87 ± 1.21

0.11 ± 0.01

0.01 ± 0.01

8.89 ± 0.81-

0.10 ± 0.01
0.03 ± 0.01

7.93 ± 0.80

32.62 ± 10.87

0.05 ± 0.01
0.02 ± 0.01

19.57 ± 3.91-

0.12 ± 0.01
0.04 ± 0.01

8.37 ± 0.70

24.46 ± 6.11

0.09 ± 0.01
0.02 ± 0.01

10.87 ± 1.21-

we perform a Monte Carlo simulation using the mass functions
from Maschberger (2013) to sample stellar masses, counting thenumber of stars in our representative mass range as well as theHR図上でのプロットとisochrone比較
12

A. L. Quintana and N.J. Wright

Figure 13. HR diagrams for all the groups. Full coloured symbols correspond to those with spectroscopy while the empty symbols are those without. Isochrones
has been drawn from the rotational models in Ekström et al. (2012). From top to bottom these are the 3.16, 5, 10 and 20 Myr isochrones. Positions of some
spectral types have been indicated on the top horizontal axis for clarity.

23.

Establishing the evolutionary timescales of the massive star formation process throughchemistry

G. Sabatini, S. Bovino, A. Giannetti, T. Grassi, J. Brand, E. Schisano, F. Wyrowski, S. Leurini, K. M. Menten
F (Abridged) Understanding the details of the formation process of massive (i.e. M<8-10M ) stars is a long-standing problem
in astrophysics. [...] We present a method to derive accurate timescales of the diﬀerent evolutionary phases of the high-massstar formation process. We model a representative number of massive clumps of the ATLASGAL-TOP100 sample which cover
all the evolutionary stages. The models describe an isothermal collapse and the subsequent warm-up phase, for which we follow
their chemical evolution. The timescale of each phase is derived by comparing the results of the models with the properties ofthe sources of the ATLASGAL-TOP100 sample, taking into account the mass and luminosity of the clumps, and the column
densities of methyl acetylene (CH3 CCH), acetonitrile (CH3 CN), formaldehyde (H2 CO) and methanol (CH3 OH). We find thatthe chosen molecular tracers are aﬀected by the thermal evolution of the clumps, showing steep ice evaporation gradients from
103 to 105 AU during the warm-up phase. We succeed in reproducing the observed column densities of CH3 CCH and CH3 CN,
while H2 CO and CH3 OH show a poorer agreement with the observed values. The total (massive) star formation time is foundto be ⇠ 5.2 ⇥ 105 yr, which is defined by the timescales of the individual evolutionary phases of the ATLASGAL-TOP100
sample: ⇠ 5 ⇥ 104 yr for 70-µm weak, ⇠ 1.2 ⇥ 105 yr for mid-IR weak, ⇠ 2.4 ⇥ 105 yr for mid-IR bright and ⇠ 1.1 ⇥ 105 yr for
HII-regions phases. Our models, with an appropriate selection of molecular tracers that can act as chemical clocks, allow to getrobust estimates of the duration of the individual phases of the high-mass star formation process, with the advantage of being
capable to include additional tracers aimed at increasing the accuracy of the estimated timescales.

Stellar Rotation of T Tauri stars in the Orion Star-Forming Complex• 大質量星形成を念頭に，ATLASGAL-TOP100
sampleをモデルに，コアのコラプスと同時
JavierSerna, Jesús Hernández, Marina Kounkel, Ezequiel
Manzo-Martínez, Alexandre Roman-Lopes, Carlos

G. Román-Zuñiga, Maria Gracia Batista, Giovanni Pinzón, Nuria Calvet, Cesar Briceño, Mauricio Tapia,に化学反応を解いて，分子種の時間変化やそれをchemicalclockとして使えないかを議論
GenaroSuárez, Karla Peña Ramirez, Keivan G. Stassun, Kevin Covey, Jaime
Vargas-González, José Gregorio
Fernández-Trincado F We present a large-scale study of stellar rotation for T Tauri stars in the Orion Star-Forming Complex.
We use the projected rotational velocity (v sin(i)) estimations reported by the APOGEE-2 collaboration as well as individual
masses and ages derived from the position of the stars in the HR diagram, considering Gaia-EDR3 parallaxes and photometry
plus diverse evolutionary models. We find an empirical trend for v sin(i) decreasing with age for low-mass stars (0.4M < M⇤ <1.2M ). Our results support the existence of a mechanism linking v sin(i) to the presence of accreting protoplanetary disks,