星形成ゼミ
2019/03/01
SFN 313 #41-#48

長谷川哲夫（国立天文台）

1

10

National Radio Astronomy Observatory, 520 Edgemont Rd., Charlottesville, VA 22903, USA
11
Istituto di Radio Astronomia, INAF-IRA, Via Pietro Gobetti 101, I-40126, Bologna, Italy
12
Centre for Astronomy, Faculty of Physics, Astronomy and Informatics, Nicolaus Copernicus University, Grudziądzka 5, 87-100 Toruń, Poland
13
NASA Goddard Space Flight Center, 8800 Greenbelt Rd., Greenbelt, MD 20771, USA
313 Iden3fying14Young
Stellar Objects in the Outer Galaxy: l = 224◦ Region in Canis Major
Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD 21218, USA
Sewi lo1,2,15Barbara A. Whitney3,4, Bosco H. K. Yung5,6, Thomas P. Robitaille7, Davide Elia8, Remy Indebetouw9,10, Eugenio#41 16Marta
INAF, Osservatorio Astroﬁsico di Arcetri, Largo E. Fermi 5, I-50125 Firenze, Italy
Schisano8,11,
RyszardAgata
Karska12,
Jennifer
MarthaJ. Fischer14,
NASA Jet PropulsionSzczerba5,
Laboratory,
California
Instituteof Wiseman13,
Technology, Brian
4800 Babler4,
Oak GroveDr., Boyer14,
Pasadena,William
CA 91109,
USA
Marilyn
Meade4,Olmi15,
Deborah2018
Padgeb16
and NataszaSio ́dmiak5
Received
2018Luca
August28; revised
November27; accepted2018 November 28; published 2019 February 1
The Astrophysical Journal Supplement Series, 240:26 (30pp), 2019 February

Abstract
We study a very young star-forming region in the outer Galaxy that is the most concentrated source of outﬂows in
the Spitzer Space Telescope GLIMPSE360 survey. This region, dubbed CMa–l224, is located in the Canis Major
OB1 association. CMa–l224 is relatively faint in the mid-infrared, but it shines brightly at the far-infrared
wavelengths as revealed by the Herschel Space Observatory data from the Hi-GAL survey. Using the 3.6 and4.5 μm data from the Spitzer/GLIMPSE360 survey, combined with the JHKs Two Micron All Sky Survey(2MASS) and the 70–500 μm Herschel/Hi-GAL data, we develop young stellar object (YSO) selection criteria
based on color–color cuts and ﬁtting of the YSO candidates’ spectral energy distributions with YSO 2D radiativetransfer models. We identify 293 YSO candidates and estimate physical parameters for 210 sources well ﬁt with
YSO models. We select an additional 47 sources with GLIMPSE360-only photometry as “possible YSO
candidates.” The vast majority of these sources are associated with high H2 column density regions and are goodtargets for follow-up studies. The distribution of YSO candidates at different evolutionary stages with respect toHerschel ﬁlaments supports the idea that stars are formed in the ﬁlaments and become more dispersed with time.
Both the supernova-induced and spontaneous star formation scenarios are plausible in the environmental context ofCMa–l224. However, our results indicate that a spontaneous gravitational collapse of ﬁlaments is a more likely
scenario. The methods developed for CMa–l224 can be used for larger regions in the Galactic plane where thesame set of photometry is available.
Key words: infrared: stars – ISM: jets and outﬂows – stars: formation – stars: pre-main sequence
Supporting material: ﬁgure set, machine-readable tables
1. Introduction
Since the advent of the Spitzer Space Telescope (Spitzer;
Werner et al. 2004), young stellar objects (YSOs) in the Galaxy
have been studied extensively (e.g., Gutermuth et al. 2009;
Povich et al. 2011; Megeath et al. 2012; Dunham et al. 2015).

relatively uncharted territories, in environments different thanin the inner Galaxy (e.g., reduced metallicity at larger
Galactocentric distances, lower cosmic-ray ﬂux;
Bloemen2
et al. 1984; Balser et al. 2011). The GLIMPSE360 survey
can detect low-mass star formation throughout the Perseus

GLIMPSE360
GLIMPSE360: Completing the Spitzer
Galactic Plane Survey

2019/03/01 2(49Completing the Spitzer Galactic Plane Survey
A Spitzer Cycle 6 Exploration Science Program

GLIMPSE360

Proposal Abstract
GLIMPSE360 proposal
GLIMPSE360 gaps proposal

Completing the Spitzer Galactic Plane Survey

GLIMPSE360 Team
Early Results
GLIMPSE360 Papers

A Spitzer Cycle 6 Exploration Science Program
GLIMPSE360 Observations

99.99% of the data have been taken. Observations started September 2009 and will hopefully conclude at the end of 2012. Source
GLIMPSE360: Completing the Spitzer Galactic Plane Survey
lists were delivered to IRSA for l=109.0-169.0 deg and will be available at IRSA shortly.

Proposal AbstractGLIMPSE360 Areal Coverage
GLIMPSE360 proposal
GLIMPSE360 gaps proposal
GLIMPSE360 Team
Early Results
GLIMPSE360 Papers

GLIMPSE360 Observations
99.99% of the data have been taken. Observations started September 2009 and will hopefully conclude at the end of 2012. Source
lists were delivered to IRSA for l=109.0-169.0 deg and will be available at IRSA shortly.

GLIMPSE360 Areal Coverage

MSX 8 micron image of the Outer Galactic plane from Price et al. 2001. The survey limits are shown with solid red lines and are
based on the midplane of the CO/HI warp at a Galactocentric radius of 13 kpc (Binney & Merrifield 1998); the dashed line shows
the position of the warp at 14 kpc. The filled circles show the mean latitude of red IRAS sources, binned by 10 deg in longitude.
Rectangular boxes show regions previously observed with IRAC. In the small programs, the exact orientation of the observing
strip is not shown, and is placed vertically on the figure.

hbp://www.astro.wisc.edu/sirg/glimpse360/“The GLIMPSE360 survey can detect low-mass star forma3on
throughout the Perseus spiral arm (at a distance of ∼3.6 kpc
http://www.astro.wisc.edu/sirtf/glimpse360/
toward l = 225°) and higher-mass star forma3on beyond.”

The Galactic Plane showing the areas covered by previous surveys in white and the GLIMPSE360 survey in blue. Th
centered onThethe Sunshows theGLIMPSE360
limit distancea solar type star (solid line) is about twice t
circles
centered
ondetection
the Sunshowsforthe(dashed line). The approximate positions of Galactic spiral arms (Taylor & Cordes 1993) are indicated in red. The dadistancesolar
arms have GLIMPSE360
tangency points that detec3onshow no excess limit
star counts(Benjamin for2007)aandmay not be density wave arms. The
and bar represent the approximate extent of the central triaxial bulge/bar (Gerhard 2002; Cole & Weinberg 2002) and
type
star
(solid
line)
is
about
twice
that
of
WISE
(Hammersley et al. 2001; Benjamin et al. 2005). The radius marking the expected truncation or break in the exponen
stellar disk(dashedis also shownline).
with a dashed line.
1 / 3 ページ

GLIMPSE360 AORs (1468 AORs)
Scheduled are 264,804 pointings x 2 bands per pointing x 2 FTs (HDR mode) = 1,059,216 frames

3
GLIMPSE,GLIMPSEII,GLIMPSE3D areal coverage

MSX 8 micron image of the Outer Galactic plane from Price et al. 2001. The survey limits are shown with solid red lines and are

The Astrophysical Journal Supplement Series, 240:26 (30pp), 2019 February

Sewiło et al.

Figure 1. Digital Sky Survey (DSS) R (left), WISE 12 μm (middle), and combined SHASSA Hα (“The Southern Hα Sky Survey Atlas”; Gaustad et al. 2001) and H2column density (based on the Herschel data; Elia et al. 2013; right) images of the Canis Major (CMa) OB1 association (Ruprecht 1966; Gregorio-Hetem et al. 2009).
All three images show the same area on the sky. Overlaid on the DSS R image in the left panel are the members of the CMa OB1 (Clariá, 1974b) and CMa R1(Shevchenko et al. 1999) associations, shown as blue and red crosses, respectively. The blue box shows the boundaries of the CMa OB1 association (Ruprecht 1966).
Several well-known open clusters (black circles; e.g., Dias et al. 2002; Moitinho et al. 2006; Kharchenko et al. 2013) and stars (red crosses; Gregorio-Hetem 2008 andreferences therein) are indicated in the left and middle panels, respectively. In the middle and right panels, the H II regions from the Sharpless (1959) catalog are shownas blue circles. The brown circle indicates the approximate position of the Hα ring. The orange rectangle in the right panel shows the region we study in this paper,
which we dubbed “CMa–l224”; it encloses the region rich in outﬂows (as traced by the 4.5 μm emission) and was selected based on the Herschel emission as
described in Section 1.

The Astrophysical Journal Supplement Series, 240:26 (30pp), 2019 February

Sewiło et al.

Figure 2. Three-color composite image combining the Herschel/Hi-GAL SPIRE 250 μm (red), WISE 12 μm (green), and Spitzer/GLIMIPSE360 IRAC 4.5 μm(blue) mosaics. The image shows a region indicated with the orange rectangle in the right panel of Figure 1.

The Astrophysical Journal Supplement Series, 240:26 (30pp), 2019 February

Sewiło et al.

data from previous near- to far-infrared (IR) surveys (seeSection 2.2). Table 1 includes a list of photometric catalogs usedin our analysis, together with the angular resolutions of theobservations.

Table Herbst
1& Assousa 1977). The studies on kinematics of theoutﬂow activity characteristic for the earliest, deeply embedded
neutralYSOs (e.g., He Photometric
et al. 2012; Cyganowski
et al. 2013).
Thus,StudytheI; Herbst &Assousa 1977of
andthereferences therein)
Data Usedin Thisandthe(HAngular
Resolutions
region at l∼224° is the site of a very young population ofand ionized (Hα and [N II]; Reynolds & Ogden 1978) gas
Surveys
protostars, actively feeding back on their surroundings.
revealed that the shell is expanding. The origin of the largeThe outﬂow area is located within the boundaries of thescale expanding shell of ionized gas in the CMa OB1
Canis MajorWavelength(CMa) OB1 association(222°<l<226°,
associationis still being debated. Oneof the hypotheses is
Instrument
SurveyFWHM−3°. 4<b<+0°. 7; see Figure 1; Ambartsumian 1949;
that the shell was produced by a supernova explosion that(μm)(arcsec)
Ruprecht 1966), ∼50′ northeast from the largest H II region
triggered star formation in the region (Herbst & Assousa 1977;
Nakano et al. 1984; Comeron et al. 1998); however,modelsin the area Sh 2-296 (hereafter S296; Sharpless 1959), at a
1.2
2MASS
All-Sky
2MASS
2–3Ha II region are not
involvingstrongstellar winds or an evolving
distance of ∼1 kpc (Clariá 1974a; Shevchenko et al. 1999;
Kaltcheva & Hilditch 2000). S296 is the most prominent
ruled out (Reynolds & Ogden 1978; Nakano et al. 1984).
1.6
2MASS
All-Sky 2MASS
2–3
feature in the images of the CMa OB1 association, forming a
CMa OB1 was covered by the 13CO(J=1–0) survey of the
2.2in the central part2MASS
All-Skyand
2MASS
2–3 deg2; 208°long arc of emissionof the association.
Monoceros
Canis Major region (∼560
Another striking feature, at optical wavelengths in particular, isl230°, −20°b10°) with the NANTEN millimeter3.6
Spitzer IRAC
GLIMPSE360
1.7
the H II region S292 (IC 2177)—a regularly shaped nebulosity
wave telescope at 2 6 resolution (Kim et al. 2004). Kim et al.
with a dark lane acrossand nearby
H II regions(2004)
detected a giant molecular cloud1.7
complex associated
4.5 its face. S292Spitzer
IRAC
GLIMPSE360
S293, S295, and S297 are associated with reﬂection nebulaewith the CMa OB1 association and comparable in size to S296.
8.3
MSX
SPIRIT
III
MSX∼20(RNe; Magakian 2003) and are members of the CMa R1
Kim et al. (2004) identiﬁed 115 CO clouds in the survey area;
Figure 3. Three-color composite image of CMa–l224 combining the WISE 12 μm (red), GLIMPSE360 4.5 μm (green), and GLIMPSE360 3.6 μm (blue) images. The
association. CMa R19is deﬁned by a groupof starsIRC
embedded in
11 arewithin theboundaries of the CMa OB1
AKARI
All-Sky
Survey5.5 association. The
region rich in outﬂows (traced by 4.5 μm emission) is located in the central part of the image. The white box indicates the area shown in Figure 4; it is very bright at
RNe, located within the boundaries of the CMa OB1
higher-resolution (43″) 12CO(J=1–0) survey with the SEST
Herschel wavelengths, indicating the earliest stages of star formation. The white contour corresponds to the H2 column density of 4×1021 cm−2. The region within
12
WISE
All-Sky
WISE/NEOWISE
6.5
association in the galactic latitude range from −3°. 4 to −2°. 0
telescope was conducted toward IRAS sources with colors of
the white circle represents the background emission in our analysis. The scale bar at the lower right corresponds to 5 pc assuming a distance of 1 kpc.
Sewiło et al.
(van den Bergh 12.1
1966). The CMa
and CMastar-forming
regions in the outer Galaxy
MSXOB1SPIRIT
III R1MSX∼20 (b10° and The Astrophysical Journal Supplement Series, 240:26 (30pp), 2019 February
associations are physically related (Clariá 1974b).
85°<l<280°) by Wouterloot & Brand (1989), detecting
formedbythetwobrightest
Herschelﬁlaments
intheregiontemperatureandH
columndensitymaps
derivedusing
the2
14.7 open clusters
MSX
III theMSX∼20
Nineteen conﬁrmed
are SPIRIT
located within
18 CO clouds
in the CMa OB1 association.
that coincide with the outﬂows (Figures 2–4). The size of the
160–500 μm data reveal a variety of conditions, with starboundaries of CMa 18
OB1 (e.g., Dias et al.
2002; MoitinhoThe
CMa–l224
outﬂow region overlaps
with the ﬁrst outer
AKARI
IRC et al.
All-Sky
Survey5.7
selected region is 2°. 5×1°. 4, with the center at (l, b)= (224°. 5,
formation preferentially distributed along ﬁlamentary struc2006; Kharchenko et al. 2013); three of them have been studied
Galaxy region (216°. 5l225°. 5,−2°b0°) studied−0°. 65); see Figures 1–3. We refer to this region as “CMa–tures. The ﬁlaments were identiﬁed in follow-up work by
21.3
SPIRIT
III 1).
MSX Space Observatory (Herschel;∼20 Pilbratt et al.
in detail: NGC 2335,
NGC 2343, andMSXNGC 2353(see Figureusing the Herschell224” throughout the paper.
Schisano et al. (2014) using the H2 column density map. The
These three open clusters
2010) datafrom the Open Time Key Project
The goal of this paper is to identify YSO candidates in CMa–
outﬂow region includes areas characterized by the lowest
22 with distances similar
WISEto that of the All-Sky
WISE/NEOWISE
12“Herschel infrared
CMa OB1 association were suspected to be physically related
Galactic Plane Survey”(Hi-GAL;
Molinarietal.
2010).
Thiss
l224 using the high-sensitivity GLIMPSE360 3.6 and 4.5 μm
columndensities
intheentiretemperaturesandthehighestH2bc
70 in early studies
Herschel
PACS1966);∼10
to the OB1/R1 complex(e.g., Ruprecht
Hi-GALHi-GAL
study by Elia et al. (2013) basedon the PACS 70 anddata combined with ancillary catalogs and images. We estimate
region studied by Elia et al. (2013) and Schisano et al. (2014).
however, the clusterin laterPACS
studies are
160 μm Hi-GAL(Poglitsch et al. 2010) and SPIRE
250, 350, andThe area studied in this work was selected based on thethe physical parameters of the YSO candidates and study their
160 ages estimated
Herschel
12c with
H2 column density map from Elia et al. (2013). It encloses
correlations with the dust and gas tracers. The methods
inconsistent with this hypothesis (∼15 to over 100 Myr versus
500 μm data (Grifﬁn et al. 2010), combined
the WISEca high column density ring-like structure (as traced by thedeveloped for this relatively small region can be applied to250
Herschel
SPIRE
Hi-GAL constructed a catalog 18∼3 Myr; e.g., Clariá
1974b; Fitzgerald
et al. 1990;
Lim et al.
22 μm photometry,of Herschel compact
21−2
4×10 cm contour); the western side of this structure is
the larger areas in the outer Galactic plane where Spitzer
c as protostellar
2011; Kharchenko et al. 2013).
sources (cores and clumps) that they classify
350
Herschel
SPIRE
Hi-GAL25
The optical and radio images of the CMa OB1 association(with a detection at 70 μm and/or at 22 and 160 μm; 255
c
3
SPIRE
36that most of thereveal a large-scale500
ring of emission Herschelnebulosity witha diametersources)Hi-GAL
and starless (688). They foundof ∼3° centered at (l, b)∼(225°. 5,−1°. 5), deﬁned by S296
protostellar sources are in the early accretion phase, while forand fainter nebulae to the north and east (see Figure 1;
prestellar sources the accretion has not started yet. The

2.1. Primary Data Set: Spitzer GLIMPSE360

Notes.
2aDepending on the atmospheric seeing.
b
The Hi-GAL source list from Elia et al. (2013).
c
Molinari et al. (2016).

photometry at longer wavelengths is not available, making themethods developed for the nearby molecular clouds unsuitable

The GLIMPSE360 data products include a highly reliable
IRAC Point Source Catalog and a more complete IRAC Point
Source Archive (see Meade et al. 2014 for details). In our
analysis, we use the IRAC Archive for completeness in boththe number of sources and the ﬂux measurements at eachwavelength. The sources included in the Archive fulﬁll lessstringent criteria than those included in the Catalog; however,these criteria were developed to ensure that each source is alegitimate astronomical source and that the ﬂuxes reported forthe IRAC bands are of high quality. The Spitzer GLIMPSE360
data are matched to the near-IR (JHK ) data from the Two
Micron All Sky Survey (2MASS; Skrutskie et al. 2006);
2MASS photometry is an integral part of the GLIMPSE360
Catalog and Archive.
About 212,000 GLIMPSE360 sources are located in CMa–l224 (see Table 2). The lower and upper sensitivity limits in the
3.6 μm band are 0.021 mJy (17.8 mag) and 1100 mJy (6.0mag), respectively. In the 4.5 μm band, the lower and upper
sensitivity limits are 0.022 mJy (17.3 mag) and 1100 mJy (5.5mag), respectively (Meade et al. 2014). The minimum signal4
Figure 4. ratio
Left: image(S/N)
showing theisregionwith a whitebox in Figure
3 in the samecombination ofcolors. This area harbors sources with outﬂows (traced by
to-noise
5 indicatedin bothbands.
Figure5 shows
exampleenhanced 4.5 μm emission). The source at (l, b)∼(224°. 6, −1°) is the reﬂection nebula Gy 3–7 associated with the embedded cluster and the high-luminosity IRAS
source IRAS 07069–1045 (Gy 3–7 #6; Tapia et al. 1997). Right: zoom-in on the area enclosed in the white box in the left panel.
color–colorand color–magnitude diagrams (CCDs and CMDs)
using
theW3Spitzer
and
2MASS
data.
18 μm bands are ∼5 5 and ∼5 7, respectively. For an S/N offor theand W5 giant
molecular
clouds andpart of the W4
region. They found that in WISE Band 3 (12 μm) and Band 4(22 μm), the AllWISE source extraction pipeline retrieves only

∼5, the detection limits in the MIR-S and MIR-L bands areabout 0.05 and 0.09 Jy, respectively (Kataza et al. 2010). Since

The Astrophysical Journal Supplement Series, 240:26 (30pp), 2019 February

処理の流れ

Table 2
Source Statistics for CMa–l224
Source TypeNo. of Sources

Spitzer/GLIMPSE360 sources
AllWISE sources
AKARI/IRS sources
MSX Point Source Catalog sources
Herschel/Hi-GAL sources (Elia et al. 2013)b
Herschel/Hi-GAL sources (Full)b

211,900
80,171 (74,726)a
102
78
273
787

GLIMPSE360 sources with valid 3.6 and 4.5 μm
photometry
AllWISE sources with valid 12 μm photometryc
AllWISE sources with valid 22 μm photometryc
AllWISE sources with valid 12 and 22 μm photometryc
Herschel/Hi-GAL sources (Elia et al. 2013) with valid
70 μm ﬂux
Herschel/Hi-GAL sources (Full) with valid 70 μm ﬂux

202,575

Matching Results

21,255 (18,919)a
8,200 (7064)a
6310 (5352)a
64
113

d

GLIMPSE360 (3.6 and 4.5 μm) with AllWISE 12 μm
matches
GLIMPSE360 (3.6 and 4.5 μm) with AllWISE 22 μm
matches
GLIMPSE360 (3.6 and 4.5 μm) with AllWISE 12 and22 μm matches
GLIMPSE360 (3.6 and 4.5 μm) with AKARI matches
GLIMPSE360 (3.6 and 4.5 μm) with MSX matches
GLIMPSE360 (3.6 and 4.5 μm) with Herschel matches(Elia et al. 2013)b
GLIMPSE360 (3.6 and 4.5 μm) with Herschel matches (Full)

Sewiło et al.

10,557/27,646e
2335/37,862e
1878/26,371e58
26
72 (40)f
206 (75)g

Notes.
a
The number in parentheses corresponds to point sources (ext_ﬂg=0).
b
Based on the Hi-GAL catalog from Elia et al. (2013). In order to be includedin the catalog, a source had to be detected in at least three consecutive bands.
c
AllWISE catalog sources with the ph_qual of “A,” “B,” or “C” (i.e., no upper
limits) in a corresponding band.
d
The matching radii between the GLIMPSE360 catalog and AllWISE, AKARI,
MSX, Herschel/E10, and Herschel/Full are 1″, 1″, 1 3, 5″, and 5″,
respectively. All the constraints described in Section 3 are applied.
e
The numbers represent the number of sources with valid data points (left) andupper limits (right) at a corresponding wavelength(s).
fThe number in parentheses corresponds to sources with valid 70 μm ﬂux. Outof 72 sources, 26 are SPIRE-only and are not used in the analysis.
g
The number in parentheses corresponds to sources with valid 70 μm ﬂux. Outof 206 sources, 111 are SPIRE-only and are not used in the analysis.

and bolometric luminosities), and their classiﬁcation. Twohundred and seventy-three Herschel sources from the Elia et al.
(2013) catalog are located in CMa–l224: 29 unbound, 171 pre-criteria (all ﬁve-band photometry is included regardless of the
1. Catalog Matchingcompleteness/regularityof the SED). We are able to identify
Spitzer/GLIMPSE360のカタログ（2MASSと対応済み）を、AllWISE,
unreliable photometry
in the subsequent analysis. There are 787
sources fromAKARI,
the FullMSX,
Herschelcatalog in CMa–l224.
Hi-GALと比較し対応天体を見つけ、長波長の情報を補う
In our analysis, we use the H2 column density and2. Distance
Determina3on
temperature
imagesderived by Elia et al. (2013) through a• graybody
付随するCOclumpのkinema3c
distance
pixel-to-pixel
ﬁt using
the Hi-GAL 160–500μm
images; they• represent
the properties of thecold component
99%の天体は距離0.5– 1.3kpc(median 0.92kpc)の範囲にある
of dust.
3. Ini3al Selec3on of YSO Candidates
The Herschel PACS and SPIRE images from the Hi-GAL
まず2MASSとSpitzer
dataからIR-excessのある天体を選び、それにフィsurvey havebeen downloaded fromthe Herschel Science
Archive (Proposal
ID:
OT1_smolinar_5).
ルターをかけてコンタミ（unresolved extragalac3c sources, PNe, evolvedstars）を除く
2.3. The
CO Data• Interstellar
ex3nc3onの見積り（HerschelのN(H2)より）
In our analysis,
we determine
kinematic –distancesthe – AllWISE)• Color-Colorplots (2MASS
Spitzer,toSpitzer
12
GLIMPSE360H-Ks
sourcesusing
the
NANTENCO(J=1–0)
vs. Ks-[4.5]等のプロットでIR-excessのある天体を選別
Galactic Plane Survey data (Mizuno & Fukui 2004) presented[3.6]-[4.5]vs.survey[4.5]-[12]のプロットがextragalac3c
sourcesの分別に有by Elia et al. (2013)
for their
area (see Section 1). The
効、など苦心がしのばれる
spatial resolutionof the survey is 2 6, signiﬁcantly lower thanthe4.
GLIMPSE360
SED Fiongresolution (∼2″). The spectra for ∣b∣ - 0 . 5
were obtained with a grid spacing of 4′. The CMa–l224 is
選ばれた候補に対し、YSO
modelsを用いたSED ﬁongを行うとともに、
covered in fullby this survey.
多波長のイメージを見て確認する（SED3.models）
The Catalog Matching•
Robitaille(2017)によるa new set of radia3ve transfer model SEDs
We positionally cross-matched the Spitzer/GLIMPSE360
catalog (alreadyRobitaille+(2006)から改良されている
matched to 2MASS) with the AllWISE,AKARI, MSX, 構成要素：aand Hi-GAL star,catalogs
provide envelope,
longerdisk,toinfalling
bipolar cavi3es, anwavelength data that, in combination with near- and mid-IR
ambient medium
data, would improve the YSO identiﬁcation process, including(Fitする対象)
the color–color/color–magnitudeselection and the SED ﬁtting.
• the12ミクロンが検出されていて、5波長以上で上限値でない測定がAs a result ofmuch higher sensitivity of the GLIMPSE360
survey in comparison
to other currently available surveys, only
ある天体→232天体（color-selected
sampleの73%）
a small fraction of GLIMPSE360 sources have matches in othercatalogs (see(Fit結果の選択）
Table 2).
ageを推定AllWISE: • To HR図上でPARSEC進化トラックと比較、mass,
match the GLIMPSE360 catalog with the
AllWISE catalog,
weadoptedamatching
radiusof
1″and→物理的におかしいﬁtを除く
selected the closest matches. We found 80,171 matches; we•
多次元→Bayesian approach
used the quality ﬂags provided in the AllWISE catalog to< χ2best+Fとなる解が多い• いきなりχ
exclude unreliable
sources.2best
In の解に飛びつくのでなく、χ2
our color–color/color–magni2
tude YSO selection
process,
we
use
the
12
and
22
μm
(w3
and
model setをまず選び、その中でχ が最小のmodelを選択
w4) photometry with a ﬂag ph_qual of “A” or “B,” excluding→ピーキーなﬁtやoverﬁongの可能性を減らす
low-S/N photometry and upper limits. The 12 and 22 μm ﬂux
5. limits
SIMBAD
serch
upper
will beused for the SED ﬁtting if both the w1 andw26.ﬂuxes
have an of
S/Nthan orinequal
to 10 Space (Class I/II)Classiﬁca3on
YSOlarger
Candidates
Color-Color5(ph_qual=“A”).
We
removeﬂuxes
ﬂaggedaslikelyspurious
7. [3.6][4.5]-only Sources (47 “possible YSO candidates”)
detections: Diffraction spikes (cc_ﬂag=“D”), Persistences(“P”), Halos (“H”), and Optical ghosts (“O”), and those ﬂagged
as extended sources with ext_ﬂg>1.

The Astrophysical Journal Supplement Series, 240:26 (30pp), 2019 February

The Astrophysical Journal Supplement Series, 240:26 (30pp), 2019 February
Sewiło et al.

Sewiło et al.

Figure 14. Example SEDs for sources at different evolutionary stages well ﬁt with Robitaille (2017) YSO models. The embedded objects with envelopes but no disks(Group e) are shown in the top left and top middle panels. The best-ﬁt model for the source in the top right panel contains both an envelope and a disk (Group d + e).
The bottom row shows SEDs of sources best ﬁt with models with disks, but not envelopes (Group d). The solid black line in each panel shows the best-ﬁt YSO model.
Filled circles and triangles are valid ﬂux values and ﬂux upper limits, respectively. The values of a reduced χ2, interstellar visual extinction, and a distance scale for the
best-ﬁt model are indicated in the plots. The ﬂux error bars are plotted if larger than the data points. See Section 6 for details.
The Astrophysical Journal Supplement Series, 240:26 (30pp), 2019 February
Sewiło et al.

between 1.4 and 4.9 GHz of 0.32 (Sν ∝ ν α; Vollmer et al.
YSOs, “envelope and disk” sources are Class I or ﬂat-spectrum
Table 4
YSOs, and “disk-only” sources are Class
II
and
Class
III
YSOs.
2010).
The extragalactic
Physical Parameters for a Subset of YSO
Candidatesin CMa–l224 nature of this source is not ruled out
The physical parameters estimated from the SED ﬁtting and(Petrov et al. 2011), but we keep it on the list of YSO
e
Model
Model
(nﬁts)setd with a RwarningTå the classiﬁcation
Lå
M
Agee
IRAC
Designation
npta for 210
the H-Rdiagram analysis
YSO candidates
are listed in(χ2minc/ candidatesåthatiså uncertain.
“SSTGLMC”
Setb
Namept)set(Re) on the(TSIMBAD(Lsearch,
(Myr)
e)
e)
e)
Table 4.
In summary, basedwe(Mremoved
one
from the list
candidates;0.443
293 sources
G223.4100–00.9594
5
spu−smi
O9yTeubb_01
0.363 source132
1.63of YSO3690
L remainLon
G223.4356–00.9369
6
sp− −s−ibRpK1icY_01
0.0262 the list.
461
0.939distribution
4180 is shown
0.243in Figure
L 15. All L
Their spatialthe7. 6SIMBAD
Search lcceyFSO_03
G223.4467–00.9365
spu−hmi3.69 photometry
30
6880 is provided
129
4
used 7.99
in the analysis
in Table
5. 0.766
G223.4519–00.9364
18
spubhmi for matches
UYPUxRp7_07
13
12.3
6270
210
5.2
0.275
We searched the SIMBAD
database
within 5″2.93
G223.4894–00.8837
7
sp− −hmi
68
1.4
5740
1.92
L
L
from the Spitzer/GLIMPSE360
positionsszXsXcnT_09
of 294 YSO0.743
G223.5416–00.7255
7 to checksp− −h−i
ddJqkFr2_03
42
2.03
5520
3.43
1.7
3.93
candidates from our list
whether thesources have1.41 8. Classiﬁcationof
YSO
Candidatesin
Color–ColorSpace
G223.5622–00.9594
7
sp− −s−ir9oRQITx_03
328
2.61
9610
52.2
2.4
3.85
been classiﬁed in the literature.
We have found
four matches0.135
G223.5881–00.9575
6
sp− −smi
5FMrwNtB_08
0.0817664 5, we 1.785900
3.46Class II L
L
In Table
provide theClass I andclassiﬁcation
with
known
sources
that
were
not
from
the
Elia
et
al.
(2013)
orG223.7198–00.8560
7
sp− −s−i
J29C1Wcg_04
0.581 for YSO
353 candidates
1.24 from 6300
2.16 on the L
L
our
list
based
classiﬁcation
Fischer et al. (2016) catalogs.
G223.7210+00.0052
7
sp− −hmi
74RteWgE_07
0.841 criteria 32from the0.975
5020For sources
0.544 with L
literature.
valid 12 L
μm
Source SSTGLMAG 224.4696–00.3780 is classiﬁed in thephotometry,
we
use
the
Fischeretal.
(2016)
criteria;
otherwise,
literature
as
a
Be
star
(HD
55135;
e.g.,
Frématetal.
2006);
we
Notes.
the classiﬁcation is based on the Gutermuth et al. (2009)
aremoved it from the list of YSO candidates.
npt is the number of valid data points used for the ﬁtting.
criteria. We include a ﬂag indicating which criteria were used.
b Source SSTGLMAG 225.3263–00.5314 has been identiﬁed
See Table 3 footnotes for the meaning of characters in the model set names.
c2Four sources do not fulﬁll either Fischer et al. (2016) orin(χ2minthe/pt)literaturegalaxy(2MASX
J07122445–1115336;
dataa point
for thebest-ﬁt modelfrom the most likely model set.
set is χ perasd2
et al.
(2009)e.g.,
et al. of2003),a χmassive
YSOpt )setanandextended(nﬁts)Paturel
ﬁts with/pt between(c 2minwith(c 2min pt )set + 2 Gutermuth
in the most likely
modelset. criteria; they are classiﬁed as YSO
set is the numberecandidates
disks
based(seeon Section
the SEDﬁtting.
We provide
masses(MSX6C(Må) and ages
(Age) only for sources Navaretewith the most reliable
estimates withof theseparameters
6.1 fordetails).
bipolarH2 stellar
emission
G225.3266–00.5318;
We use the Gutermuth et al. (2009) criteria for Class I and
et al.
andin itsa entirety
Herschel/Hi-GAL
protostellar
core
(This
table2015),
is available
in machine-readableform.)Class II based on the (Ks−[3.6])0 versus ([3.6]–[4.5])0 CCD(G225.32629–00.53244; Elia et al. 2013). Since the extra(see Figure 16),
where thesubscript
dereddenedgalactic
natureof this source
is notwe keep
thedistributionof Catalog
sourcesfrom “0”the indicates“off-position”
in theThe YSO
candidates
from ourlistconﬁrmed,
are overplottedon thecolors.CMD.
The extinction
is estimatedrepresents
based on thecolorssource on
the list
YSOclassiﬁcation:
candidates. (Ks−[3.6])0 versus
same
The “off-position”
the2MASS
background/
CCDs
used
for ofthe(theyrequire
at
least
H
and
K
photometry
andmagnitude
Source SSTGLMAG
224.3005–00.8417
wasidentiﬁed
asas color–magnitude cuts listed
foregroundpopulations.
We
apply([3.6]–[4.5])
inFigure16(Gutermuth
et
al.
2009)
and
[3.6]–
0
uncertainties<0.1
minimize thecontaminationfrommemberof theEX Lupi-likeclass17of(Fischer
YSO outbursts
by Millerin
Equation (2)to mag).
removeTobackgroundgalaxies(see the right[4.5]
versus[4.5]–[12]in Figureet al. 2016).
extragalactic
sources,
apply thetheremaining
brightness47limit
using
et al.
(2015):
iPTFshow
15afq.CCDs
Thesewith
outbursts,
with a duration
from
panel
of Figure18). Wethey
considersourcesas
The top panels
YSO candidates
identiﬁed
dereddened
3.6candidates”
μm photometry: [3.6]0<15.0 and [3.6]0<a few
weeksettoal.
severalare interpreted
as middle),
the accretion“possible
YSO
inDunham(2015;months,
left), Fanget al. (2013;
and
6
14.5
for
Class
I
and
Class
II
YSO
candidates,
respectively.
events leftdrivendisk right).
instabilities(e.g., Audard
et al. 2014).
The
Megeathetpanel)al. by(2012;
Figure 13. [3.6]–[4.5] vs. [4.5]–[12] CCD showing the YSO (top panel) and non-YSO (bottom
sourcesfrom the literature(as indicatedin the legends),
andspatial distribution of the [3.6][4.5]-only possible YSO
The Fischer
et al. (2016)
modiﬁed version
EX right
Lupipanel;
objects
have5.3).
circumstellar
disks9 caption
and nofor references.
candidates
is shown
overlaidcriteria,
on theaHslightly
the sources from the “off-position” with the Spitzer–AllWISE selection criteria overlaidThe(bottom
SectionSee the Figure
2 column density imageof Figurethe Koenig& vastLeisawitz(2014)
the WISE
envelopes; this is consistent with the results of the SED ﬁtting.
in
19. Themajorityof thecriteria,
sources use
are associated
9. [3.6][4.5]-only
Sources coincides with a
w1 -thew 2 highest
2 -2)wregions,
3 CCD (seeversus wN(HFigure with
17). Wereplace
Source SSTGLMAG
225.1412–01.3289
withcoincident
clustersofthe WeWISE
bands w1
w2 withGLIMIPSE360
3.6 μm andradiosource,(Wright
et al.
1996;
Vollmer
Weingartner&(2001)
extinction
lawwithfor
RV=4.0.
are included, and thus the UV and optical ﬂuxes for
YSO
candidates;
thisandlocation
supports
their classiﬁcation
asAbout
56% PMN
of Draine
theJ0709–1128
GLIMPSE360
ArchivesourcesbothThe Astrophysical Journal Supplement Series, 240:26 (30pp), 2019 February

Figure 21. Left: same image as in the bottom panel in Figure 20, zoomed in on the three brightest ﬁlam
et al.
solidSewiłolines) andYSO candidate positions (plus signs) indicated.

prestellar cosupported by(e.g.,4.5
AndréFigure 25. Three-color composite images combining the WISE 12 μm (red), Spitzer/IRAC
μm (gTh
C18O (right). The white dashed lines show the coverage of the Olmi et al. (2016) CO observations.
One of the90, 95)% of thwhite contours in the left panel correspond to the 12CO emission at the level of (70, 80,shock-compr
H2O maser is indicated with a yellow diamond. The image in the right panel is slightly zoomed in
ﬂows are velo
correspond to the C18O emission; the contour levels are (40, 50, 60, 70, 80, 90)% of the C18O peak of 2
main
axiswithocontaining sources associated with the extended 4.5 μm emission (likely outﬂows) areindicatedﬁlaments (e.gSuchkinemaobjects like AGB stars is very low because those would be
was support
observations
1367334). T
more evenly distributed.
etal. 2013)
appointment
We compare our list of YSO candidates to the list of YSO
observed
Goddard aloSpcandidates from Fischer et al. (2016) identiﬁed in Canis Majoretal. (2016)
Space
Rese
based on the WISE data. Eleven out of 93 Fischer et al. (2016)
thetransvers
NASA
and
WISE YSO candidates located within the boundaries of CMa–higher-resolu
Thiswork wl224 are not on our ﬁnal list of YSO candidates: two were notconclusions.
grant 2014/1
detected by Spitzer, one was resolved into two Spitzer sources,
In an alternFigure22.
Histogramsofdistancesfromtheﬁlament
branchesfor
YSO
observations
three were removed as non-YSOs, and ﬁve did not fulﬁll the
candidates well ﬁt with YSO models containing envelopes only (red),
have
been
operatedby
2MASS–Spitzer(four)
or[3.6][4.5]-only(one)criteria(theenvelopes+disks (green), and disks only (blue). The vertical dotted line
ﬁlaments
or,of Technolo
AllWISE
used ofsincecorresponds photometry
to 31″ or 0.15 was
pc at nota distance
1 kpc.the distance between
asalsosomemakesstuu
the AllWISE and Spitzer sources is larger than 1″). We foundcollapseof th
Sky Survey,
that
Group
00 of
WISE
YSO
candidates
fromFischer
et al.
ﬁlament
might
have
startedat thelocation ofCluster
B, where
ages
of
the
Y
Massachuset(2016)
corresponds
to three
YSO
candidateclusters
mass
would
be accreted
fromSpitzer
both themainand secondary
Figure 23. Two-color composite image combining the WISE 12 μm image (red) and the H2 column density image (blue). The YSO candidates wellﬁt with
YSO
the
YSO
can
Center/Calif
distributed
along the main ﬁlament (A, B, and C). The selection
ﬁlaments.
models with the envelopes only (“envelopes”; orange circles), envelopes and disks (“envelopes + disks”; green circles), and disks only (“disks”; whitecircles) areof
the supern
NationalAe
criteria
based
on
the
GLIMPSE360
data
recovered
more
of
the
The
AstrophysicalJournalSeries,240:26 (30pp), 2019 February
Sewiło et al.
The
positional
coincidence
betweenthemainﬁlament
in).
concentrated
in regionswith Supplement
the highest N(H21998).
National
Sc
evolvedYSO
candidates(“disk-only”).
CMa–l224 and the Hα shell indicates that there may be a
Survey Explo
We discussstar formationin CMa–l224
contextof theconnection
betweenthe supernova
event in inthea CMa
OB1/R1
California, Llargerregion—theCMa
OB1
association.
The
mainﬁlament
inassociation and star formation in CMa–l224. According to one
California In
CMa–l224
lies ofat therim offormation,
the Hα shell,hinting atarea possibleof the theories
ﬁlament
the ﬁlaments
formedAeronautics
connection
betweenthesupernova
event
in
the
CMa
OB1
Using Spitby the large-scale compression of the interstellar material inspace2MASS
obser
associationand star formation inﬂows,
the region.
Webediscuss
withsupersonic magnetohydrodynamic
which canturbuEuropean-led
possiblestar formation(1)andthereferences
formationtherein).
of theHi-GAL 70–
lent or nonturbulent(e.g.,scenarios:
André 2017
tant2participa
ﬁlaments
in from
CMa–l224
as a resultexplosionof the supernova
explosion
indegregion
The wavesthe supernova
are possible
drivingtheparticipatheCMaforOB1
the gravitational
dubbedCMasources
the association,
large-scale followed
turbulenceby(e.g.,
Rybakin &Experiment;
fragmentation
andthus,the theformation
prestellar could
cores;have(2)
tionof sourc
Goryachev 2018);
ﬁlaments ofin CMa–l224
Missile Defesupernova-induced
star formationin the preexistingoutﬂowtrac
been formed as a resultof the supernova
event in ﬁlaments;
the CMa
NASA Ofﬁc
and
(3)
the
spontaneousgravitational
collapseoftheﬁlaments.
coincides
wi
OB1/R1 association. The formation of the ﬁlament is followed
theOur
SIMBADAll
formationfragmentation
scenarios seemand
to be
YSO
by the stargravitational
theplausible;
formationhowever,of preour
results
indicatethat the spontaneousgravitational
collapseand Spitzer–
stellarcores,
the precursors
of stars. Sucha two-step
starofﬁlaments
is the(i.e.,
mostscenario.
Moreevidenceisvisual inspec
formation
scenario
thelikely
ﬁlament
formation
followedby theneeded to draw stronger conclusions.
Catalogs fr
25
We thank the anonymous referee for reading the paper
carefully and providing useful comments. We thank Leo
We usedYSO candid
Girardi for providing the TRILEGAL simulation. This work
Figure 24. Left and right: three-color composite image combining the WISE 12 μm (red), Spitzer/GLIMPSE360 4.5 μm (green), and Herschel/Hi-GAL PACS 70 μm

12 from Fischer et al. (2016)(left),
andandYSO
paper,
those(blue) images.
The positionof theIRASsources,thePlanckcores,an H2O maser
Figure25. Three-color
composite
imagescombining
WISE cold
12 μm(red),and
Spitzer/IRAC
4.5 μm(green),COcandidates(blue; Olmifrom
et al. this2016)
images:CO (left) and18
only,
and those
et al.
(2013)
are overlaid
indicatedin theCOlegends.
The greensolidimages
circlesareinscaled
the leftpanel indicate
IRAS sources
associated with theO (right).
The from
white Elia
dashedlines
show only
the coverageof theasOlmi
et al. (2016)
observations.
The COto highlightthe brightest
regions. The
C
12
−1
−1
CS emission.
See
for details.
white
contours in
theSection
left panel12correspondto the CO emission at the level of (70, 80, 90, 95)% of the peak emission of 29.8 Jy beam km s . The position of the
H2O maser is indicated with a yellow diamond. The image in the right panel is slightly zoomed in on the main ﬁlament. The white contours in the right panel
correspond to the C18O emission; the contour levels are (40, 50, 60, 70, 80, 90)% of the C18O peak of 2.9 Jy beam−1 km s−1. The approximate positions of the regions
containing sources associated with the extended 4.5 μm emission (likely outﬂows) are indicated with the green circles.

non-YSO contaminants. The ﬁnal list of YSO candidates inages). We divide sources into three groups depending on the

3.6 and 4.5 μm photometry only (with no AllWISE matches or

The Herschel/SPIRE emission in CMa–l224 is highly

presence byof thetheSpitzer
envelope
and/or disk:“envelope-only”(16),
CMa–l224
sources.
objects
like contains
AGB stars293
is very
low because those would be
was supported
GLIMPSE360
project(RSA
We
identify
anadditional
47
sourceswiththe
GLIMPSE360“envelope
and
disk”
(21),
and
“disk-only”(173).
1367334). The work of M.S. was also supported by anmore evenly distributed.

27

7

5

INAF/Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125, Firenze, Italy
Dipartimento di Fisica “Enrico Fermi”, Università di Pisa, Largo Pontecorvo 3, 56127 Pisa , Italy
7
Department of Astronomy & Astrophysics, 525 Davey Laboratory, Pennsylvania State University,
University Park PA 16802, USA
IRAMLaboratory,
and GaiaM/S
views
of mul3-episodic
starforma3onin IC1396A:
The origin and
8313
Jet Propulsion
180-703,4800 Oak Grove Drive,
Pasadena,
CA 91109,USA
6#42 dynamics of the Class 0 protostar at the edge of an HII region

Received 11
AprilSicilia-Aguilar1,
2018 / Accepted
17 December
AuroraNimeshPatel2, Min2018
Fang3, Veronica Roccatagliata4, Konstan3n Getman5 and Paul Goldsmith6
A&A 622, A118 (2019)

ABSTRACT
Context. IC 1396A is a cometary globule that contains the Class 0 source IC 1396A-PACS-1, which was discovered with Herschel.
Aims. We use IRAM 30m telescope and Gaia DR2 data to explore the star formation history of IC 1396A and investigate the possi-

bilities of triggered star formation.
Methods. IRAM and Herschel continuum data were used to obtain dust temperature and column density maps. Heterodyne data reveal
the velocity structure of the gas. Gaia DR2 proper motions for the stars complete the kinematics of the region.
Results. IC 1396A-PACS-1 presents molecular emission similar to a hot corino with warm carbon chain chemistry due to the UV
irradiation. The source is embedded in a dense clump surrounded by gas at velocities that are significantly different from the velocitiesof the Tr 37 cluster. CN emission reveals photoevaporation, while continuum data and high-density tracers (C18 O, HCO+ , DCO+ , andN2 D+ ) reveal distinct gaseous structures with a range of densities and masses.
Conclusions. By combining the velocity, column density, and temperature information and Gaia DR2 kinematics, we confirm thatthe globule has experienced various episodes of star formation. IC 1396A-PACS-1 is probably the last intermediate-mass protostar
that will form within IC 1396A; it shows evidence of being triggered by radiation-driven implosion. Chemical signatures such as CCS
place IC 1396A-PACS-1 among the youngest known protostars. Gaia DR2 data reveal velocities in the plane of the sky ⇠4 km s 1
for IC 1396A with respect to Tr 37. The total velocity difference (8 km s 1 ) between the Tr 37 cluster and IC 1396A is too small for
IC 1396A to have undergone substantial rocket acceleration, which imposes constraints on the distance to the ionizing source in time
and the possibilities of triggered star formation. The three stellar populations in the globule reveal that objects located within relatively
close distances (<0.5 pc) can be formed in various star-forming episodes within ⇠1–2 Myr. Once the remaining cloud disperses, we
expect substantial differences in evolutionary stage and initial conditions for the resulting objects and their protoplanetary disks, which
may affect their evolution. Finally, evidence for short-range feedback from the embedded protostars, and in particular, the A-type star
V390 Cep, is also observed.
Key words. stars: protostars – photon-dominated region – stars: individual: IC 1396A-PACS-1 – stars: individual: HD 206267 –open clusters and associations: individual: Tr37 – molecular data

1. Introduction
The IC 1396A dark globule is one of the classical examplesof bright-rimmed clouds (BRC) at the edge of an H II region(Sharpless 1959; Osterbrock 1989; Patel et al. 1995). The globule is illuminated by the O6.5 trapezium-like system HD 206267

beautifully imaged in H↵ (Barentsen et al. 2011) and by IR space
missions, such as AKARI (Huang & Li 2013).
Early molecular-line observations suggested the 8presence ofhighly embedded sources and the potential of the region toundergo a substantial episode of star formation (Loren et al.
1975). The total gas content was estimated to be about 200 M

1Stephen L. Skinner1, Kimberly R. Sokal2, and Manuel Güdel3

Center for Astrophysics and Space Astronomy (CASA), Univ. of Colorado, Boulder, CO 80309-0389, USA; stephen.skinner@colorado.edu
2
Department
ofthe
Astronomy,
Univ.Star-Formingof Texas, Austin, TX
78712-1205,
USA2
313 Chandra 3Observa3onsofMassive
RegionOnsalaDepartment of Astrophysics, Univ. of Vienna, Türkenschanzstr. 17, A-1180 Vienna, Austria
L. Skinner1, Kimberly R. Sokal2 and Manuel Guedel3#43 Stephen
Received 2018 September 2; revised 2018 December 4; accepted 2018 December 5; published 2019 January 25
The Astrophysical Journal, 871:116 (20pp), 2019 January 20

Abstract
Previous radio and infrared observations have revealed an obscured region of high-mass star formation in Cygnus
known as Onsala 2 (ON2). Within this region lies the optically revealed young stellar cluster Berkeley 87, whichcontains several OB stars and the rare oxygen-type Wolf–Rayet star WR 142. Previous radio studies of ON2 havealso discovered masers and several H II regions excited by embedded OB stars. Radio and Gaia parallaxes havenow shown that the H II regions are more distant than Berkeley 87. We summarize two Chandra X-ray
observations of ON2, which detected more than 300 X-ray sources. Several optically identiﬁed stars in Berkeley 87
were detected, including massive OB stars and WR 142, the latter being a faint hard source whose X-ray emission
likely arises in hot thermal plasma. Intense X-ray emission was detected near the compact H II regions G75.77+0.34 and G75.84+0.40, consisting of numerous point sources and diffuse emission. Heavily absorbed X-ray
sources and their near-IR counterparts that may be associated with the exciting OB stars of the H II regions are
identiﬁed. Shocked winds from embedded massive stars offer a plausible explanation of the diffuse emission.
Young stellar object candidates in the ON2 region are identiﬁed using near-IR colors, but surprisingly few
counterparts of X-ray sources have near-IR excesses typical of classical T Tauri stars.
Key words: HII regions – open clusters and associations: individual (Berkeley 87) – stars: individual (WR 142) –stars: massive – X-rays: diffuse background – X-rays: stars
1. Introduction
Massive young stars spend their short pre-main-sequence
lifetimes embedded in the dense cloud material from which
they formed and are invisible to optical telescopes. Thus,
previous studies have traditionally focused on radio andinfrared observations to probe the gas and dust environmentsof massive newborn stars. X-ray observations have emerged asanother technique for studying young heavily obscured stars
since X-rays at energies of a few keV can penetrate high visualextinction of up to AV∼50 mag. X-rays provide information

previous faint X-ray detection of WR 142 (Oskinova et al.
2009; Sokal et al. 2010). This study extends our earlier
Chandra X-ray study of optically revealed massive stars in thesouthern part of ON2 (Sokal et al. 2010). The Chandra data
provide lower background emission, higher spatial resolution,and more reliable source identiﬁcation compared to theprevious XMM-Newton X-ray observation. However, XMM9
Newton provides better sensitivity to soft X-rays
at energies
below ≈1–2 keV.
An optical study of the region by Turner & Forbes (1982,

The Astrophysical Journal, 871:116 (20pp), 2019 January 20

Figure 1. POSS2 red image of ON2 with Chandra ACIS-I detector footprints(17′×17′) overlaid. Radio positions of H II regions are shown. Positions of
three optically identiﬁed members of the Berkeley 87 cluster are marked (BD+36 4032, HD 229059, and WR 142). The M3 supergiant BC Cyg is IR-brightbut is not detected in X-rays. The outer bounding box has dimensions
30′×30′. N is up, E is left.
The Astrophysical Journal, 871:116 (20pp), 2019 January 20

The OH maser lies in close proximity to the ultracompact
H II (ucH II) region G75.78+0.34, which has been the subjectof many radio continuum studies (e.g., Matthews et al. 1973,
hereafter M73; Harris 1976; Matthews et al. 1977; Matthews &Spoelstra 1983; Wood & Churchwell 1989). Because of its
ultracompact morphology, G75.78+0.34 is thought to be a
massive star at a very early evolutionary stage. Radio
trigonometric parallax measurements give a distance of
3.56 kpc to maser G75.78+0.34 (Xu et al. 2013), so thisucH II region lies twice as far away as Berkeley 87 and is not
physically associated with the optically revealed cluster.
Several other H II regions in ON2 have been identiﬁed andtheir positions are marked in Figure 1. Two of these, G75.77+0.34 and G75.84+0.40, have dense groupings of X-ray
sources in their immediate vicinity suggestive of recent starformation, as shown below. The compact H II (cH II) region
G75.77+0.34 has a double-peaked radio morphology with aseparation of ≈19″ between the radio peaks (M73, Harris 1976;
Matthews et al. 1977). This region emits brightly in the IR at
10.7 and 19 μm (Hefele et al. 1977). Lying in the northern partof ON2 is the cH II region G75.84+0.40, which also has
double-peaked radio morphology with a separation between
peaks of ≈20″ (M73). Previous radio studies of G75.77+0.34
and G75.84+0.40 have usually assumed a distance ∼5–6 kpc
based on the kinematic distance d=5.5 kpc to ON2 estimatedby Reifenstein et al. (1970), but the study of Garay et al. (1993)

The Astrophysical Journal, 871:116 (20pp), 2019 January
20 Sokal, & Güdel
Skinner,

Skinner, Sokal, & Güdel

toward G75.77+0.34 by Matthews et al. (1986) and interpreted
as possible interaction between the cH II region and its+
The AstrophysicalJournal, 871:116 (20pp), 2019 January 20
surroundings. Dense gas was mapped in NH3 and HCO
byDent et al. (1988) who found evidence for local gas heatingnear the embedded stars that excite the H II regions. Shepherd
et al. (1997) obtained high spatial resolution molecular line andmillimeter continuum maps of the ON2 core region. They
detected at least four massive molecular outﬂows and a 3 mm
continuum peak at G75.78+0.34. Two fainter continuum
sources probably associated with dust-enshrouded young stellar
objects (YSOs) were found. They also noted an enhanced SiO
abundance, which they hypothesized could be due to grain
destruction caused by shocks from winds or outﬂows impacting
surrounding material. Submillimeter mapping at 450 and
800 μm by Hunter (1998) detected several compact sources(candidate protostars) near the H II regions.
The Chandra observations discussed here reveal densegroupings of X-ray point sources near G75.77+0.34 andG75.84+0.40, some of which are likely to be young stars. We
ﬁnd evidence for diffuse X-ray emission near these H II regions(see also Oskinova et al. 2010) with spectral properties
suggestive of a thermal origin. Diffuse X-ray emissioninterpreted as thermal emission from shocked stellar winds
Figure 2. Left:Mosaicked Chandra ACIS-I images (0.2–8 keV) of ON2 from observation ID (ObsID) 9914 (ON2S) and ObsID 18083 (ON2N). Ellipses (3σ) markhas also been detected by Chandra in other massive starChandra X-ray source positions. Crosses (+) show H II region mean radio continuum positions. WR 142 lies in the overlap region and was detected in bothforming Nregions,
the Arches(Yusef-Zadeh
observations.
is up and Eincluding
is left. Right:Same
as at left cluster
but heavily
binned and smoothed to bring out the regions with brightest X-ray emission. X-ray properties of
theet
B1al.
Ia starHD 229059
were given
in Sokal
et al. (2010).
No(Townsley
optical or TwoetMicron2002),and M17and theRosetteNebulaal. All Sky Survey (2MASS) counterpart was found for the X-ray source marked“anon”(Chandra
J202152.99+372139.33)
but
it
is
coincident
withafaintNebula
Spitzer mid-IR source (Section 3.7). The image is binned by a factor of 8 (binned pixel2003). Diffuse thermal X-ray plasma in the Orionsize=3 94). Contours are at (4, 5, 6, 8, and 10) counts per binned pixel.
region excited by massive Trapezium stars was also reported on
the basis of XMM-Newton observations by Güdel et al. (2008).
In contrast, the diffuse X-ray emission
detected
Chandra in
Skinner,
Sokal, &by
Güdel
Table as1 nonthermal (Wolk et al. 2002), separately for each of the four ACIS-I CCDs using the input
RCW 38 has been interpreted
Chandra X-Ray Observations of ON2
and diffuse emission in Westerlund 1 is consistent with either a value falsesrc=1.0 to assure no more than one false source
Parameter
nonthermal or hot thermal plasma origin (Muno et al. 2006). per scale size. Each run was executed using four scale sizes
We conﬁrm here the previous
142 as a hard (wavelet radii): 1, 2, 4, and 8 pixels with a pixel size of 0 492.
RegionON2S detection of WR
ON2N
We thus anticipate no more than four false sources perab
X-ray source and characterize
the emission 18083
of thebrightest
ObsID
9914
ACIS-I CCD in each observation. Checks for X-ray variability
Startdate/time(TT) in 2009
Feb 2/02:41:24
13/09:29:27
X-raysourceON2,
which has 2016
no Augcataloged
opticalwere made using the CIAO glvary tool (Gregory & Loredo
Stopdate/time
(TT)
2009
Feb
2/22:52:55
2016
Aug
14/05:31:06
counterpart.
1992, 1996). For brighter sources, spectra and associatedccInstrument
Livetime (s)Frame time (s)ACIS-I
ACIS-I
70,14869,066
2. Observations
3.1 and Data Reduction
3.2

instrument response ﬁles were extracted using CIAO specextract. Spectra were ﬁtted using HEASARC’s XSPEC ﬁtting
package (v. 12.8.2) to obtain estimates of absorption column
A Chandra observation of the southern portion of Onsala 2 density (NH), mean plasma energy (kT) for thermal models,Notes.
a (ON2S) was obtained in 2009 with pointing
centerednear WR photon power-law index (Γph) for power-law models, normalNominal pointing position was (J2000) R.A.=20h21m
46 42, decl.=142, whichlies about 3′ south of G75.77+0.34 (Figure 2; ization (norm), and X-ray ﬂux Fx (absorbed and unabsorbed).
9.
+37°22′44bhSokal etpointing
al. 2010).
Thewasnorthern(ON2N)
wasdecl.=observed We used the CIAO Image of Diffuse Emission thread to search
Nominalposition(J2000) portion
R.A.=20
21m38 39,
1. with pointing centered ≈1′ south of G75.84+0.40. for evidence of diffuse emission near the cH II regions+37°30′04
in2016c
ACIS-I
consists of four used1024×1024
pixel CCDswith a pixelofBoth observations
the ACIS-I(Advanced
CCDsizeImaging(Section 3.8).
The combined
FOV array,is 16 9×16
The energy
range
is
0 492.
Spectrometer)
imaging
which 9.capturesa ≈16
9×169E≈0.5–10 keV. For an on-axis point source the 90% encircled energy
ﬁeld of view (FOV). There was some overlap in the FOV
3. Results
radius is R90≈0 9. The energy resolution at 1.49 keV is ≈0.13 keV.
between the two exposures as shown in Figure 2. Table 1
3.1.
X-Ray
SourcesON2keV) of a 1′×1′
Figure 3. Top:Mosaicked Chandra ACIS-I
image in
(0.2–8summarizes the observations and ACIS-I detector properties.
regionnear
G75.77+0.34,
combining
data
from
both
observations.
Chandraimageswith
Using the Portable Interactive Multi-Mission Simulator, the Figure 2 (left) is a mosaic of the two Chandra10
at ∼3.5 kpc, the detection limits are higher and correspond to source
more position
errorsource
ellipsespositions(3σ) are shownin black(ObsID9914)
and blue
X-rayoverlaid.
In
total,
201
sources
wereintrinsic (unabsorbed) X-ray luminosity detection limits for an on(ObsID 18083).
The×symbols
mark209
UGPS
near-IRcounterpart
positionsmassive cTTS with M*>1 Me (Sections 4.1 and 4.3).
detectedin ON2N andsourcesin ON2S.
A comparison
axis point source and a 5-count detection threshold (0.3–8keV)Rectangles
are position error boxes for Westerbork radio continuum
Data analysis was basedﬁles 2).
−1 on standard processing (Table−1 of the X-ray
positionsofthesources
in
the
two
ﬁelds
reveals
were log L =29.79 erg s (ON2S) and log Lx= 29.88
erg s G75.77+0.34A and B from Matthews et al. (1973). The diamondsources
provided by xthe Chandra X-ray Center (CXC) andfollowed
34
sourcesintheoverlap
region
whose
R.A.
and
decl.
differ(ON2N) at a characteristic distance of 1.75 kpc. These
valuespeak A marks the position of 2MASS J202142.13+372554.3. The
standard CXC processing threads using version 4.9 ofoverlappingtheby 1″.source
Thesesources are probableduplicatedetections.
assume a generic T Tauri star thermal plasma spectrum
withpositionof Chandra
J202140.54+372548.3
overlapsthat of Westerbork

Figure 4. Top:S

WR 142
WR 142 is a Wolf-Rayet star in the constellation Cygnus, an extremely rare star on the WO oxygen sequence.

WR 142
Observation data
Epoch J2000.0

Contents

Constellation

Discovery
Features
Evolutionary status
See also

Equinox J2000.0

CygnusRight ascension

20h 21m 44.3s[1]Declination+37° 22′ 30.56″[1]Apparent magnitude (V)12.94[2]Characteristics

References

Spectral type

WO2[3]

Apparent magnitude (J)9.538[1]Apparent magnitude (H)8.889[1]Apparent magnitude (K)8.596[1]highly-ionised OVI.[8] Because of the unusual oxygen lines, seen in only a handful of other stars, it was given theU−B color index−0.29[4]spectral type WC5pec in the Sixth Catalogue of Galactic Wolf-Rayet Stars.[5]B−V color index+1.43[5]Discovery
In 1966, a search for Wolf-Rayet stars in the northern celestial hemisphere discovered seven new examples.
One, designated as Stephenson 3, was classified as WC.[7] It was later found to show unusual emission lines ofAstrometry

In 1981, described as a WC-OVI star, it was identified as being associated with the active star-forming region
ON2,[9] and then a heavily-obscured open cluster designated Berkeley 87, 9.5′ south of the red supergiant BC

Proper motion (μ)

RA: −6.270[6] mas/yr
Dec.: −3.422[6] mas/yr

Parallax (π)

0.5755 ± 0.0284[6] mas

Distance

5,700 ± 300 ly

Cygni.[4]WR 142 - Wikipedia

2019/02/26 14)17

In 1982, the WC-OVI stars were grouped as members of
https://en.wikipedia.org/wiki/WR_142

the new WO class. The class at that time consisted of fivestars, two of which were in the Magellanic Clouds and(1,740 ± 90 pc)

one of which was later found to be the central star of aplanetary nebula.[10]20[2] M☉

Radius0.40[3] R☉

Luminosity (bolometric)

245,000[3] L☉

Luminosity (visual, LV)

847[2] L☉very well known but thought to be around 1.23

Temperature

200,000[3] K

kiloparsecs (4,000 light-years). As with its home cluster

Metallicity [Fe/H]

0.0[3] dex

its light is very reddened and extinguished by interstellar

Rotational velocity

1,000[2] km/s

WR 142 is usually assumed to be a member of the open
Location of WR 142, circled (the bright star at
the centre is γ Cygni and north is to the right)

Details
Mass

Featurescluster Berkeley 87, whose distance from the Sun is notdust.[11](v sin i)

This star, of spectral classification WO2, is one of the very few known oxygen-sequence Wolf-Rayet stars, justfour in the Milky Way galaxy and five in external galaxies. It is also one of the hottest known with a surface
temperature of 200,000 K.[3] Modelling the atmosphere gives a luminosity around 245,000 L, while
calculations from brightness and distance give luminosities of 500,000 Lor more. It is a very small dense star,with a radius 40% of the Sun's but a mass 20 times greater. Very strong stellar winds, with a terminal velocity of
5,000 kilometers per second are causing WR 142 to lose 10−5 M/year.[2] For comparison, the Sun loses (2-3) x

10−14 solar masses per year due to its solar wind, several hundred million times less than WR 142.

Other designations
WR 142, 2MASS J20214434+3722306,
GSC 02684-00001, Sand 5, St 3,
UCAC2 44891902

WO Wolf-Rayet stars are the lastevolu3onary stage of the most
massive stars before exploding as
supernovae, possibly with agamma-ray burst (GRB).[12] It isvery likely that WR 142 is on its
last stages of nuclear fusion, nearor beyond the end of helium
burning.[13] It is es3mated to
explode as a supernova inapproximately 2,000 years. The
mass and rapid rota3on make a
GRB likely.[3]Database references
SIMBAD

Hard X-Ray emission has been detected from this star with the help of the Chandra space telescope, that has
been suggested to be caused by the presence of a companion, a B-type main sequence star located at a distanceof 1 AU from WR 142. There is no other indication of a companion and other reasons for the x-ray luminosity are considered more likely.[11]Evolutionary status

1 / 4 ページ

Absolute magnitude (MV) -2.49[2]Evolu3onary Status

data (http://simbad.u-st
rasbg.fr/simbad/sim-id?Ident=WR+142)

11

4

Max Planck Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany
Argelander Institute for Astronomy, University of Bonn, Auf dem Hügel 71, D-53121 Bonn, Germany
6
Department of Astronomy, University of Florida, Gainesville, FL 32611, USA
7
Department of Physics, University of Florida, Gainesville, FL 32611, USA8
Mul3component
kinema3cs
in a massive
INAF-Osservatorio
Astroﬁsico
di Arcetri, ﬁlamentary
Largo E. Fermi 5,IRDC
I-50125 Firenze, Italy
9
Vlas Sokolov1,
Ke
Wang2,3,
Jaime
E.
Pineda1,
Paola
Caselli1,
JonathanD.
Henshaw4,
AshleyTorrejón
T. Barnes5,
C. Tan6,7,
Centro de Astrobiología (CSIC/INTA), Ctra de Torrejón a Ajalvir, km 4, E-28850
de Jonathan
Ardoz, Spain
Francesco
Fontani8
andIzaskunJiménez-Serra9
Received 2018 September 30; revised 2018 December 16; accepted 2018 December 18; published 2019 February 7
The Astrophysical Journal, 872:30 (17pp), 2019 February 10
5

313
#44

Abstract
To probe the initial conditions for high-mass star and cluster formation, we investigate the properties of dense
ﬁlaments within the infrared dark cloud (IRDC) G035.39–00.33 (G035.39) in a combined Very Large Array andGreen Bank Telescope mosaic tracing the NH3 (1, 1) and (2, 2) emission down to 0.08 pc scales. Using
agglomerative hierarchical clustering on multiple line-of-sight velocity component ﬁtting results, we identify sevenextended velocity-coherent components in our data, likely representing spatially coherent physical structures, someexhibiting complex gas motions. The velocity gradient magnitude distribution peaks at its mode of
0.35 km s-1 pc-1 and has a long tail extending into higher values of 1.5–2 km s-1 pc-1, and it is generally
consistent with those found toward the same cloud in other molecular tracers and with the values found toward
nearby low-mass dense cloud cores at the same scales. Contrary to observational and theoretical expectations, we
ﬁnd the nonthermal ammonia line widths to be systematically narrower (by about 20%) than those of N2H+ (1–0)
line transition observed with similar resolution. If the observed ordered velocity gradients represent the coreenvelope solid-body rotation, we estimate the speciﬁc angular momentum to be about 2×1021 cm2 s−1, similar tothe low-mass star-forming cores. Together with the previous ﬁnding of subsonic motions in G035.39, our results
demonstrate high levels of similarity between kinematics of a high-mass star-forming IRDC and the low-mass starformation regime.
Key words: ISM: clouds – ISM: individual objects (G035.39–00.33) – ISM: kinematics and dynamics – stars:
formation
Supporting material: data behind ﬁgures, interactive ﬁgure
1. Introduction
The formation of stars from the diffuse ISM is a multistage
process of gas channeling and mass assembly. In the wellestablished low-mass star formation picture, individual stars areformed from dense cores, which become gravitationallyunstable once enough mass has been accumulated from theirpoorly constrained by observations. As these conditions
directly regulate the subsequent fragmentation and infall of
the star-forming material(Motte et al. 2018; Tan 2018), highangular resolution observations of the dense infrareddark gas
12with little evidence of disruptive feedback can provide
improved constraints on the theoretical models and numerical

The Astrophysical Journal, 872:30 (17pp), 2019 February 10

Sokolov et al.

Figure 1. Channel maps of the ammonia (1, 1) line for the IRDC G035.39. Each panel shows the brightness temperature of the NH3 (1, 1) spectral cube channel of
0.2 km s-1 width, with text boxes at the bottom indicating the velocity centroid of the channel. The velocity range of 42–47 km s-1 was chosen to capture the main
hyperﬁne group of the inversion transition. The light-gray contours show the infrared extinction contours starting from AV=30 mag and progressing inward in stepsof 30 mag (Kainulainen & Tan 2013).

studies of the cloud(labeled F3 in Henshaw et al. 2013, 2014;
Jiménez-Serra et al. 2014). A ﬁlamentary-like feature connects

kinematical structure exhibited in the high-resolution ammonia
emission, standard algorithms commonly used to ﬁt the line

13

The Astrophysical Journal, 872:30 (17pp), 2019 February 10

Sokolov et al.

The Astrophysical Journal, 872:30 (17pp), 2019 February 10
The Astrophysical Journal, 872:30 (17pp), 2019 February 10

Sokolov et al.

Sokolov et al.

Figure 2. PPV diagram of all ﬁtted velocity components within IRDC G035.39. Every point in the diagram represents a best-ﬁt velocity centroid for a single ammonialine component, and the darker shade of black indicates greater density of points along the on-screen projection. When constructing this diagram, we required the S/Nof both (1, 1) and (2, 2) lines to be greater than 3. The coordinates axes are speciﬁed relative to the offset at α(J2000)=18h57m08s, δ(J2000)=+2°10′30″. An
interactive version of the diagram is available with this ﬁgure, while a separate copy will be maintained under https://vlas-sokolov.github.io/post/cloudh-ppv/.

The Astrophysical Journal, 872:30 (17pp), 2019 February 10

Sokolov et al.

excessive amount of human interaction, resulting in a low level
FWHMs are not allowed to overlap), the model isrejected. This restriction is imposed to avoid lineof reproducibility, while methods for global nonlinear regresblending.
sion are prohibitively computationally expensive for such a
3. A rejected model is replaced with a simpler one: thehigh volume of data. Therefore, to ﬁt multiple line-of-sight line
three-component model is replaced with the two-compcomponents, we employ the same ﬁtting strategy as in Sokolov
onent one, and a rejected two-component model is
et al. (2017), where an open-source Python package10 is used toreplaced with a one-component model.
assist the Levenberg–Marquardt ﬁtter to ﬁnd an optimal
4. The replaced models are required to meet the samestarting point for exploring the parameter space of the spectral
heuristical criteria as in step 2.
model. The free parameters for the ammonia model(same
5. If none of the ﬁtted models are valid, we consider theformalism as in, e.g., Friesen et al. 2017), implemented within
spectrum to be a nondetection.
the PYSPECKIT package(Ginsburg & Mirocha 2011), are split
6. Steps 1–5 are repeated for all the spectra in the combined
into a multidimensional parameter grid, and a synthetic
spectral cube, until a collection of ammonia ﬁtting
ammonia model is then computed for each. Subsequently,
parameters is assembled in a position–position–velocity
every spectrum in our VLA data is then cross-checked with
(PPV) space.
every model computed to calculate the corresponding residual.
The model parameters for a given spectrum that yield the best
In the above procedure, the second step requires imposing
residual (i.e., the least sum of squared residuals) are taken to beconstraints for both (1, 1) and (2, 2) lines. This is done to assure
the starting value for the Levenberg–Marquardt ﬁttingthat the physical properties that depend on the population ratio
algorithm.
of the two transitions (kinetic temperature and ammonia
Three independently derived sets of best-ﬁtting parameterscolumndensity)Theare
reliably
constrained.
Figure3. PPV diagramof theaﬁtted
velocityfor
components
withinIRDC
G035.39
along theR.A. projection.
coordinates
are givenin arcsecondThroughout
offset relative tothis
theareobtainedinsuchfashion,
one-,
two-,
and
threethis istothe
methodwe refervelocity
to as components
the strictare
censoringδ(J2000)=+2°10′30″ coordinate. All the data are plotted in paper,
black, similarly
Figure2, and individual
marked inα(J2000)=18h57m08s, 11
component
models.
The
threesets aremergedareinto
different colors.
The data not
found
to ﬁttingbe associatedwiththenany clusters
plotted in criterion
gray. The ﬁgure
shows
the
clusteringobtainedwiththestrictmasking
criteria
for arriving at the PPV structure in the IRDC.
one
following
a set
criteriathat determine(introducedin Section2.2);oftheheuristicalequivalent ﬁgure
for relaxedcensoring canthe
be found However,in Appendixwe
C. In
addition
R.A. projection,a projection
decl. isnotethat tointhe
principle,
shouldonly thealongkinematic
shown in theaxis for the components
F1 ﬁlament.
numberof insetline-of-sightin a given pixel of theinformation be needed from our ammonia data, either one ofspectral cube. The heuristical procedure is described below:
the inversion line detections would constrain it in our
1. First, the three-component
model is considered to be
3. Results
valid.
3.1.
Velocity
Components
2. If any of the three components are below S/N of 3—for
andspace(2, 2)with
transitions—themodel isboth2 NH
Figureshowsthe1)PPVthe best-ﬁt line-of-sight3 (1,marked which
as rejected.
Similarly,theshouldhalf ofthe sum ofcomponents,
demonstrates
complexkinematics
any neighboring
components’
FWHM
thanmaythethe G035.39
cloud. As
different partsof bethelesscloud
separation
their velocity
centroids(i.e.,it the
represent
unique,between
noninteracting
physicalstructures,isnecessary to identify the components that are linked in both10
https://github.com/vlas-sokolov/multicube
positionand velocity. The inability to separate such unique
11
Upon closewould
inspection,no spectra
containinganalyses
more than on
threethevelocity
structureshinderquantitativegas
components were identiﬁed.

-1). ﬁtting.
simultaneous
Therefore,
we
consideran additional
Figure4. MapoflinederivedVlsr values
biggestvelocity
components(lefttopanel:
F2;F1,
right
top to bottom:
F1,velocity
F5, F4),gradient
overlaid with the velocity gradientdata (0.2Figure
Additionally,
links
with
widthsof for the (left
km sline
4. Mapof derived Vlsrno
valuesforthe biggest
velocity
components
panel:
F2; right
panels, topbottom:F5,panels,
F4), overlaid
with themethod,which
weresolution
call relaxedcensoring,which
naturallyThe
directions
ofthearrowsthe steepest
velocity ﬁeld
change whilein thetheredshifted
direction,
while thelengths
twice thearrows.
velocity) wereallowedtovelocity
be point
km
s-1pointThe
directionsof(0.4the arrows.
arrowstothesteepest
ﬁeld tochangein the redshifted
direction,
arrow lengths
represent
the arrowrelative
vectorrepresent the relative vector
magnitudesof thegradient.
The
open
ﬁlled redstars
the positionsof 70Nguyen
μm Herschelfrom Nguyen
Luong
et 20
al. (2011)
results
in magnitudes
a PPV
structurewith a The
largeraveragenumberlineof thespatially
gradient.
open
anddistinct
ﬁlled
red
starsofdenotethe andpositionsof 70
μmdenote
Herschelsources from
Luongsources
et al. (2011)
below andabove
Me, below and above 20 Me,
establishedto prevent
close
yet
components
respectively,
and
the
red
diamonds
indicatetheposition
of
cores
from
Butler&Tan(2009,
2012).
The
overlaid
contours
indicatethehighest extinction contours fromrespectively,
theour
red diamonds
indicate
position
ofpoints
coresof-sight
components.
the differences
inPPV
these
twofrom Butler & Tan (2009, 2012). The overlaid contours indicate the highest extinction contours fromfrom being
merged.andInDespite
ACORNS
run,&thethe
=40
to
120
mag,
progressing
in
steps
of
20
mag.
The
relative
position
on
the
plane
of
skyisof the components shown is
KainulainenTan(2013),
rangingfromA
V in steps of 20 mag. The relative position on the plane of sky of the components shown
120derive
mag, progressing
Kainulainen& Tan (2013),
rangingfrom AV=40
approaches,
the average
properties
we to
will
below S/N
signiﬁcanceofkinematic
3 wereset
tobeintoexcluded.
Appendix
E.ﬁgureThe data
used to create this ﬁgure are available.
displayedinAppendix E. changed
Thedisplayed
data usedcreate thisare available.
14
later
will
not
be
signiﬁcantly
ifarelaxedcensoring
We obtain two sets of results from the ACORNS clustering
methodis used(Appendix
C provides
quantitative
comparison
for the twomasking
strategies
introducedin Section
2.2. Thebetweenthe two
approaches);
therefore,clusters,
we use7 thestrict-1 -1strict censoring
results
in 13 hierarchical
of km
which
s pc
the IRDC.
Figure6 shows ofan theoverall
distribution of the kinetic
magnitudesof
0.3–2.5, which are thebroadly
consistent1
1
s pc us, VLA
IRDC.
Figure 6 shows
an overall
distribution
kinetic
ofbecause
0.3–2.5
which
are broadly
consistentcensoringapproachallowsto discussthecover an magnitudes
area
in excessof ﬁveitkm
synthesized
beams(i.e.,
with
ourﬁndingsinG035.39(Figure 5).
temperature
values,whileofthethe
spatialdistribution of the ammoniawith
ourﬁndingsinG035.39(Figure5).
temperature
structureofthevelocity
components
inaddition
temperature
values,while
thespatialdistribution
ammoniahave physical properties constrained in more than 116 pixels).
derived values
gas kinetic
temperature
values7isforpresented in Figure 7 fortokinematics.
derived gas kinetic temperature
is presented
in Figure
Thetheiroverallview of the structures identiﬁed is shown in1College of Physics, Jilin University, Changchun, Jilin 130012, People’s Republic of China
2
Max-Planck-Institut für Astronomie, Königstuhl 17, 69117 Heidelberg, Germany
3
NAOCUKZN Computational Astrophysics Centre (NUCAC), University of KwaZulu-Natal, Durban 4000, South Africa
4Center for High Energy Physics, Peking University, Beijing 100871, People’s Republic of China
5313
On theGravita3onal
Instabili3es
Protoplanetary
DisksUSA; min.li@unlv.edu, minphyli@163.com
Departmentof Physics
and Astronomy,
University of of
Nevada,Las Vegas, NV 89154,
Sui1,2, Ping He1,3,4,
MinAugustLi5 16; accepted 2018 December 3; published 2019 January 29
Receivedand2018#45 Ning
Publica3ons of the Astronomical Society of the Paciﬁc, 131:034301 (11pp), 2019 March

Abstract
Gravitational instabilities are important to the evolution of the disks and the planet formation in the disks. We
calculate the evolution of the disks that form from the collapse of the molecular cloud cores. By changing
the properties of the cloud cores and the hydrodynamical viscosity parameters, we explore their effects on theproperties of the gravitational instabilities. We ﬁnd that the disk is unstable when the angular velocity of themolecular cloud core is larger than a critical value. The time duration of the instability increases as the angularvelocity of the core increases. The increase of the hydrodynamical viscosity parameter hardly affects the stability ofthe disk, but decreases the time duration of the critical state of the gravitational instability in the disk. The
instability of the disks can happen at very early time of evolution of the disk, which is consistent with theobservations.
Key words: accretion – accretion disks – protoplanetary disks – stars: formation – stars: pre-main sequence

1. Introduction
The gravitational instabilities of the protoplanetary disks are
signiﬁcant in the evolution of the disks. The instability
contributes to the transport of angular momentum in the disk
(Tomley et al. 1994; Bodenheimer 1995; Ruden 1995; Durisen
et al. 2007; Forgan et al. 2011; Armitage 2011; Rice et al.
2011). FU Orionis events that occur at protostar stage (Bellet al. 1995) can be explained by the gravitational instability(Armitage et al. 2001). Gravitational instability also plays animportant role in the formation of planetesimals, gas giant
planets, brown dwarf companions, and binary stars (Laughlin& Rozyczka 1996; Mayer et al. 2004; Matzner & Levin 2005;
Rice et al. 2006; Boss 2011, 2012). Observations show thatmomentum leads to a disk which is more likely to be unstable.
Vorobyov & Basu (2010) gave the same conclusion and further
showed that high temperature of the cloud core moderates thefragmentation of the resulting disk.
The instability of the disks also relates to the viscosities inthe disks which drive the evolution of the protoplanetary disks(Balbus & Hawley 1991; Sui & Li 2016). Generally, theviscosities are parameterized by the α-prescription (Shakura &Sunyaev 1973). Usually, the α-parameters are considered to beconstant with time and radius. Using these kinds of disks,
Vorobyov & Basu (2009) showed that if15the viscosity
parameter α increases, the disk becomes more gravitationallystable.

Sümeyye T. Suri , Álvaro Sánchez-Monge , Peter Schilke , Seamus D. Clarke , Rowan J. Smith ,Volker Ossenkopf-Okada1 , Ralf Klessen4 , Paolo Padoan5,6 , Paul Goldsmith7 , Héctor G. Arce8 , John Bally9 ,John M. Carpenter10 , Adam Ginsburg11 , Doug Johnstone12,13 , Jens Kau↵mann14 , Shuo Kong8 , Dariusz C. Lis15,16 ,
17 , Thushara Pillai18 , Jaime E. Pineda19 , Ana Duarte-Cabral20
Steve Mairs
313 The CARMA-NRO
Orion
Survey: The ﬁlamentary structure as seen in C18O emission

iv:1901.00176v1 [astro-ph.GA] 1 Jan 2019#46

S. T. Suri1, A. Sanchez-Monge1, P. Schilke1,
S. D. Clarke1,
R. J. Smith2,
V. references)Ossenkopf-Okada1, R. Klessen3, P. Padoan4,5, P.
(Affiliations
can be foundafter theGoldsmith6, H. G. Arce7, J. Bally8, J. M. Carpenter9, A. Ginsburg10, D. Johnstone11,12, J. Kauﬀmann13, S. Kong7, D. C. Lis14,15, S.
Received 08, 2018; accepted 12, 2018
Mairs16, T. Pillai17, J. E. Pineda18 and A. Duarte-Cabral19
arXiv 1901.00176
ABSTRACT

Context. We present an initial overview of the filamentary structure in the Orion A molecular cloud utilizing a high angular and velocity resolution

C18 O(1–0) emission map that was recently produced as part of the CARMA-NRO Orion Survey.
Aims. The main goal of this study is to build a credible method to study varying widths of filaments which has previously been linked to starformation in molecular clouds. Due to the diverse star forming activities taking place throughout its ⇠20 pc length, together with its proximityof 388 pc, the Orion A molecular cloud provides an excellent laboratory for such an experiment to be carried out with high resolution and highsensitivity.
Methods. Using the widely-known structure identification algorithm, DisPerSE, on a 3-dimensional (PPV) C18 O cube, we identified 625 relatively
short (the longest being 1.74 pc) filaments over the entire cloud. We study the distribution of filament widths using FilChaP, a python packagethat we have developed and made publicly available.
Results. We find that the filaments identified in a 2 square degree PPV cube do not overlap spatially, except for the complex OMC-4 region
that shows distinct velocity components along the line of sight. The filament widths vary between 0.02 and 0.3 pc depending on the amount of
substructure that a filament possesses. The more substructure a filament has, the larger is its width. We also find that despite this variation, thefilament width shows no anticorrelation with the central column density which is in agreement with previous Herschel observations.
Key words. stars: formation – ISM: clouds – ISM: structure – ISM: individual objects: Orion A – methods: statistical

1. Introduction

Over the past decade, it has become clear that stars preferentially form along and at the junctions of filaments inside cold,dense molecular clouds (e.g. Schneider et al. 2012). Filaments
pervade molecular clouds whether these clouds form high-mass
or low-mass stars, or lack star formation entirely (André et al.
2010, 2016). Earlier studies of dust continuum observed with
the Herschel Space Observatory suggest a “universal width” ofthe observed filaments (Arzoumanian et al. 2011). The distribution of widths of 90 filaments having column densities of 1020 –1023 cm 2 peaks at 0.1 pc with very little dispersion (±0.03 pc,
Arzoumanian et al. 2011). However, work by Smith et al. (2014)shows that the calculated filament width is dependent on the fitting range considered when analyzing the radial profiles. They

vations allow studying intertwined filamentary networks which
otherwise appear as single objects when observed in dust emission due to projection of 3-dimensional information (positionposition-velocity, hereafter; PPV) on to 2-dimensional column
density space, therefore providing a more accurate determination of the filament properties. Understanding the role of filaments in the process of star formation requires high spatial andspectral resolution studies of, ideally, multiple molecular emission lines that trace di↵erent densities within the cloud, as wellas continuum emission. This way, filament properties can be calculated for di↵erent tracers that probe a variety of physical and16 propchemical conditions. Investigating physical and kinematic
erties of filaments weighted by the environmental and chemical
conditions allows us to gain an understanding of how filaments

Suri et al.: The CARMA-NRO Orion Survey: The filamentary structure as seen in C18 O emission

OMC-3
OMC-2
OMC-1

OMC-4

OMC-5

L1641-N
V380

HH34

HH1/2

Fig. 1. Moment maps of C18 O J = 1–0. From left to right the panels display the integrated intensity emission, the centroid velocity, and the velocity
dispersion. Relevant star-forming regions are marked on the integrated intensity map. Beam-size shown in the bottom left corner is five times thereal beam-size (5x0.015 pc) of the observations for demonstration purposes.

ferent properties in the northern part of the cloud compared tothe more di↵use filaments identified in the southern part (see thefollowing sections and Figures 5 and 10). A large number offilaments identified by DisPerSE have short lengths (< 0.15 pc),
and seem to be continuations of one another, resulting in a longer
but fragmented (as seen in C18 O) filament. We therefore do not
exclude them from further analysis, in particular for the width
determination. However, properties of only those filaments withan aspect ratio larger than 2:1 are shown in the Appendix B.

The pixels that lie on or adjacent to the calculated perpendicularlines are used to produce the radial intensity profiles in a similar
fashion to Duarte-Cabral & Dobbs (2016).
This process is shown in Figure 3. On the top panel of the
figure the peak intensity map towards the OMC-4 region that isoverlaid with an identified filament spine is shown. Two perpendicular slices from which we extract intensity profiles are also
overlaid on the filament spine in red. An intensity profile is constructed using intensities of the pixels that are perpendicular to17

A&A proofs: manuscript no. main

Dec (J2000)

Dec (J2000)

Suri et al.: The
CARMA-NRO
Orion Survey: The filamentaryies). Instead, the intensities
fromeach perpendicularslicestructure
are as seen in C18 O emission
integrated over the two neighbouring velocity channels aroundthe velocity at which the filament is identified. For example, if a
-5°32'of0.1
point along the filament is at a velocity
5.5pckm s 1 the intensity
profile that is perpendicular to this 33'
point is calculated from theintegrated intensity between 5.28 km s 1 and 5.72 km s 1 with
34'km s 1 . This helps reducethe velocity channel width being 0.22
the confusion in PPV space by discarding
emission other than
35'
that which belongs to the identified filament and it makes sure36'
that the emission included in the constructionof the intensity
profiles is not separated from the velocity
at which the filament is37'
identified by more than 2 times the sound speed (approximately
24.00s12.00s5h35m00.00s
resolution),provid0.2 km s 1 at 10 K , close to our velocity
RA (J2000)
ing coherence. The velocity coordinate along the filaments
oftenchange, therefore, the integrated intensities are computed indi0.1 pc-5°34'12.0"
vidually for each PPV point along
a filament.

30.0"

3.1.1. Filament width determination methods
48.0"⌃ p (r) = A p "1+

Fig. 2. Filaments identified in the Orion A molecular cloud. Greyscaleshows the C18 O peak intensity emission, while the colored segments
indicate the filaments identified by DisPerSE.

distance from the spine. The negative distances from the spine
point to the north-east direction, and positive distances to southwest of the filament. The grey thin lines indicate the di↵erent
profiles extracted from each slice. Blue, yellow and red curves
fitted to the averaged intensity profiles in the right panel denote
Gaussian and Plummer–like fits with p=2 and 4, respectively.
Grey dashed lines indicate the fitting boundaries. Green dashedlines indicate the shoulders of the averaged profile. A shoulder is
identified where there is a local minimum in the second derivative of an intensity profile. Within FilChaP the shoulders are always identified on the smoothed intensity profiles. For this study,⇢c R f lat⇣⌘◆2 # p2 1✓r

R f lat

Integrated Intensity [K.km/s]

In order to calculate the widths, we employ four di↵erent meth35'06.0"
ods, a Gaussian fit, two fits with
Plummer–like functions withpower law indices p=2 and 4, and24.0"
the second moment of distribution. Plummer-like functions have been used to describe thecolumn density of a filament with a dense andflat inner por12.00s
5h35m07.20s
RA (J2000)
tion and a power-law decline at larger radii (Arzoumanian
et al.
2011):
5
4

(1)

3
2

where ⇢c is the central density, R f lat 1is the radius within which
the density is uniform, p is the power law index, and A p is a finite
0constant that is related to the filament’s inclination compared to
0.4
0.2
0.0
0.2
0.4
the plane of the sky (here we assume that thisangle
is zero).
The
Distance from the ridge [pc]thn moment of a distribution I is given by:
P3. Top: C18 O peak intensityFig.
emission overlaid with a filament spine (in orange) and 2 perpendicular slices (in red and yellow). The yellown(xslicesindicate the extent of emission fitted with FilChaP. The green box marks the zoom-in region shown in the middle panel.
x̄)
1portioniN ofIitheiemission levels of 2, 3, 5, 7, 9, 11 K. Middle: Expanded view of the(2)perpendicular slices. The perpendicular pixels used to
mn = Contours indicaten
at each radial bin are marked with crosses. Bottom: The intensity profiles (grey lines) extracted from the slices overlaid on
Nextract the intensities
the map in the middle panel. Blue, yellow and red curves fitted to the averaged intensity profile (black) indicate Gaussian and Plummer-like fitsandintensity
4, respectively.weighted
The vertical grey
dashed lines
show the ofpositionsthe identified minima considered during the fitting procedure
where x̄withisp=2
themean
positionthe ofprofile,(see Sect. 3.1). Vertical green dashed lines correspond to shoulders as explained in Sect. 4.1.
Ii is the intensity at position xi , and is the intensity weighted
standard3.1.2.
deviationof the profile. The second momententire
is thevariTreating the bias caused by the averaging length
length of a filament. Although averaging helps to obtain asmoothradial of
profile, it has the disadvantage that any
ance of the distribution
andtherefore
wecalculate
the
width
p
p (e.g. Arzoumanian et al. 2011) calculated fila- information thatintensity
could be used to understand how and why thePrevious
the profileas studies
m2 ⇥ 2 2ln2. For the fits, we use thecurve_fit
widthchangesalongasinglefilament is lost. Therefore, we firstment widths using intensity profiles that are averaged over thefunction contained within the scipy.optimize package with the
Article number, page 5 of 19
default Levenberg–Marquardt algorithm.
We refrain from fitting the power law index of the Plummerlike function due to the degeneracy between p and R f lat . Instead,18
we use the literature values of the index; p=4 for the densityof a homogeneous isothermal cylinder (Ostriker 1964) and p=2,
the value derived in Herschel studies when p is fitted as a freeparameter (Arzoumanian et al. 2011), which is attributed to non-Suri et al.: The CARMA-NRO Orion Survey: The filamentary structure as seen in C18 O emission

Fig. 9. Filament width (left panel) and number of shoulders (right panel) for the filaments idnetified in Orion A overlayed on the C18 O peakintensity map. The color of each point along filaments represents the width (left) and the number of shoulders (right) at that point, corresponding
to the colorbar shown in the right side of the panels. The arrows indicate the positions of hubs where filaments converge.

acteristics of filaments; FilChaP, a python based algorithm. In
this method, while calculating the filament widths: (i) the radialintensity profiles are not averaged over the length of an entire
filament, but over a number of consecutive slices that represent a
region 3 beam-sizes in length, (ii) there is no fixed fitting range
for these profiles that would force a prefixed width. Instead, eachprofile has its own fitting range. This range is set by the min-critical filament widths are coupled to the column density.
The Heitsch (2013) model with magnetized and externallypressured filaments best reproduces the observed decorrelation.The model agrees with our observations except for highcolumn densities (1023 cm 2 ) where narrower widths are predicted. We found that even at these high column densities
traced by our filaments the lack of correlation persists, how-19313 The forma3on of rings and gaps in wind-launching non-ideal MHD disks: three#47 2018
dimensional
December 13. Received
November 27; insimula3ons
original form 2018 October 4
Scob S. Suriano1,2, Zhi-Yun Li2, Ruben Krasnopolsky3, Takeru K. Suzuki1 and Hsien Shang3
MNRAS 484, 107–124 (2019)

ABSTRACTPrevious axisymmetric investigations in two dimensions (2D) have shown that rings and gaps
develop naturally in non-ideal magnetohydrodynamic disc–wind systems, especially in thepresence of ambipolar diffusion. Here, we extend these 2D simulations to three dimensions(3D) and find that rings and gaps are still formed in the presence of a moderately strong
ambipolar diffusion. The rings and gaps form from the same basic mechanism that was
identified in the 2D simulations, namely, the redistribution of the poloidal magnetic flux
relative to the disc material as a result of the reconnection of a sharply pinched poloidal
magnetic field lines. Thus, the less dense gaps are more strongly magnetized with a large
poloidal magnetic field compared to the less magnetized (dense) rings. The rings and gaps
start out rather smoothly in 3D simulations that have axisymmetric initial conditions. Nonaxisymmetric variations arise spontaneously at later times, but they do not grow to such an
extent as to disrupt the rings and gaps. These disc substructures persist to the end of thesimulations, lasting up to 3000 orbital periods at the inner edge of the simulated disc. The
longevity of the perturbed yet still coherent rings make them attractive sites for trapping large
grains that would otherwise be lost to rapid radial migration due to gas drag. As the ambipolar
diffusivity decreases, both the disc and the wind become increasingly turbulent, driven by themagnetorotational instability, with tightly wound spiral arms becoming more prominent in thedisc.
Key words: accretion, accretion discs – (magnetohydrodynamics) MHD – protoplanetary
discs – ISM: jets and outflows.

om https://academic.oup.com/mnras/article-abstract/484/1/107/5266837 by National AstronomicalObservatory, J

Astronomy, University of Virginia, Charlottesville, VA 22904, USA
ca, Institute of Astronomy and Astrophysics, Taipei 10617, Taiwan

20

et al. 2017), rapid pebble growth at the condensation fronts of313 2018
TheDecember
young10;star
popula3on
L118820
December 11. Received
in originalform 2018 of
September
Szegedi-Elek1, M. Kun1, A. Mo ́or1, G. Marton1 and B. Reipurth2#48 E.MNRAS
484, 1800–1814 (2019)
ABSTRACT

We present new results on the young star population of the Lynds 1188 molecular cloud,which is associated with the Cepheus Bubble, a giant interstellar shell around the association
Cep OB 2. In order to reveal the star-forming scenario of the molecular cloud located on the
supershell, and to understand the history of star formation in the region, we identified andcharacterized young star candidates based on an Hα emission survey and various publishedphotometric data sets. Using Gaia Data Release 2 astrometry, we studied the spatial distributionof the young star candidates and we isolated three groups based on their distances. We
constructed spectral energy distributions of our target stars, based on photometric data from thePanoramic Survey Telescope and Rapid Response System (Pan-STARRS), the Two-Micron
All-Sky Survey (2MASS), the Spitzer Space Telescope and the Wide-field Infrared Survey
Explorer (WISE), estimating their spectral types, extinctions and luminosities. We estimatedmasses by means of pre-main-sequence evolutionary models, and we derived accretion rates
from the equivalent width of the Hα line. We studied the structure of the cloud by constructing
a new extinction map, based on Pan-STARRS data. Our results show that the distribution of
low-mass young stars in L1188 is well correlated with that of the dust and molecular gas.
We identified two small, compact clusters and a loose aggregate of young stars. We foundthat star formation in L1188 started about 5 Myr ago. The apparent age gradient of youngstars across the cloud and the ammonia cores located to the east of the optically visible youngstellar groups support the scenario of star formation propagating away from the centre of the
Cepheus Bubble.
Key words: stars: formation – stars: pre-main-sequence – stars: variables: T Tauri, Herbig
Ae/Be – ISM: clouds – ISM: individual objects: Lynds 1188.

ded from https://academic.oup.com/mnras/article-abstract/484/2/1800/5300113 by National Astronomical Obser

stronomy, University of Hawaii at Manoa, 640 N. Aohoku Place, Hilo, HI 96720, USA

21