THE STAR FORMATIONNEWSLETTER
No. 261
塚越崇 (茨城大)
09 Oct. 2014

A critical analysis of shock models for chondrule formation
S. M. Stammler and C. P. Dullemond• コンドリュール生成には急激な加熱と冷却が必要• ʻsolar nebulaʼでの重力不安定・惑星などによるショックによる生成がポピュラー• nebularショックモデルによるコンドリュール生成の計算• downstream boundary conditionとsideways radiative coolingに着目• モデル• 1D (Desch&Connolly 2002)• 平面円盤仮定で無限遠に広がる• Particle dynamics• 温度、速度、ダストサイズを仮定• Gas dynamics• 様々な分子種を含む。温度と速度は共通。

• Desch&Connolly(2002)と同じインプットを使用して計算• 1350Kを越えるとダストが融けはじめる• ショック面ではガスのみパラ

Temperature [K]

ダストが融ける

Velocity [• ショック前面でも温度上昇

Temperature [K]sink in this case. Instead of raising the temperature, energy is con4
sumed in breaking molecular bonds. Simulations withoutchemical
3000
3000
reactions show an overall increase in temperature downstream by
2500
2500
!100 K. We want to point out that only a small fraction of the H2 is22000
dissociated. This demonstrates that it is very important2000
to include
chemistry in such simulations: the temperature deviation
can
1500
1500
0
-10
-5
0
5
10
make the difference between losing particles via evaporation
or
1000
1000
not. Whether or not other chemical reactions are equally important
Time [h]500500
has to be investigated.
Fig. 5.0Velocity of the particles.
Fig. 4 shows the cooling rates of the particles in the 0standard
-10
-5
0
5
10
-10
-5
0
5
10
shock at temperatures T > 1500 K which is the important temperTime [h]Time [h]ature regime for crystallization. As seen here the peak of the cool1900
ing rates is at the upper limit of Fig.
whatis allowedby experimental
3. Left:
Particle temperature
in a shock with shock speed V shock ¼ 6:5 km/s. Right: Same as left but with gas temperature. See text fo
constraints. The standard shock itself is too weak to completely
1800
Tpeak
melt the particles. In addition to that the cooling rates drop to zero
Tpost
8
1700
at already 1600 K. This is discussed
to reachlater
1350inK, thisthensection.
there’s also a lack of the ‘‘opacity knee’’
noticeable.
Another interesting feature
is shown in Fig. 5: the velocity of
1600
6
The
chemical
reactionof hydrogendissociation
works as energy
the particles in the simulation. As soon
as thepre-heatingsetsinsinkbecause
in this case.
of raising
the temperature, energy is conthe particles’ velocity decreases
the Insteadgas velocity
decreases.
1500
4
sumed in breaking molecular bonds. Simulations without chemical
From initially 6.5 km/s to roughly 5.5 km/s shortly before the shock
reactions show an overall increase in temperature downstream 1400
byfront. After the gas velocity is changed according the Rankine–!100 K. We want to point out that only a small fraction of the H2 is2
Hugoniot conditions at the dissociated.
shock frontThis
thedemonstrates
particles are
rapidly
1300
that it is very important to include
5.4 5.6 5.8 6 6.2 6.4 6.6 6.8
decelerated by the drag forcechemistry
within minutes!1.0 km/s.the temperature deviation canin such tosimulations:
8
S.M. Stammler, C.P.
Shock 0-10
speed [km/s]
The Mach
number of the gas
is initiallyof the between
order of losing
5. But particles
due
-5
0
make
the difference
via evaporation or
重力不安定性など
to the increase in temperature
decreasevelocity
it isonly are equally important
not.and
Whetheror notinotherchemical
reactionsTime [h]Fig. 6. The peak andpost-shock temperatures of the standard shock for di
2000
has
to
be
investigated.
roughly 2 at the position of the shock, where the Mach number
shock speeds. The upper solid line denotes a temperature of 1820 K, whi
Fig. 5. Velocity of the particle
Fig. 4conditions.
shows the cooling rates of the particles
in theis applied in the Rankine–Hugoniot
assumethestandard
particles have to reach to be completely molten. The lower solid
shock
at temperatures> 1500K whichisis the atimportanttemper1400 K. Thisis thetemperature the particles have to reach at least after the
The problem of chondrule
formation
withinT this
scenario
1500
ature regime for crystallization. As seen here the
peak
of
the
coolto retain their volatiles. This is not possible in the standard case.
summarized in Fig. 6. If we assume that the particles have to reach1900
ing rates is at the upper limit of what is allowed by experimental
a temperature of at least 1820 K to be completely molten and at

メータ変化。ダストは摩擦ででコンドリュールを生成するのは難しい=>光学的に厚いグローバルなショックでは難しい

Velocity [km/s]• 急速な加熱と冷却が共存せず、この条件

Temperature [K]加熱されるが、冷却も効く•• 光学的に薄いとcoolingが効くので再現[K/h]

Temperature [K]constraints. The standard shock itself is too weak to completely
1000
1800
Tpeak
melt the particles. In addition to that the coolingthe
rates
drop
to
zero
same time to cool rapidly back down
below at least 14
Tpost
1700
at already 1600 K. This is discussed later in thisto
section.
retain volatile materials which are observed in them, th
500
Another interesting feature is shown in Fig.not
5: thevelocity
possiblein ofa one-dimensional 1600
model of an optically thick3
微惑星のbow
shockみたいなケース
theparticlesinthesimulation.
As
soon
as
the
pre-heatingsets
in
6.0 x 10If we increase the gas mass density (see Fig. 7) the pro
the particles’ velocity decreases because the gas velocity decreases. 0
1500
improves
slightly
and
the
two
lines
0
5 of peak and10post-shock
From initially 6.5 km/s to roughly 5.5 km/s shortly before the shock
1400 at unrealistic densiti
perature
are further apart. But Time
only
front. After the gas velocity is changed according
the Rankine–
[h]

出来る
•

Temperature [K]!

di↵erent evolutionary stages and to compare these results withother starforming regions observed by HIFI. This paper presentso-H2 O 110 ! 101 and C17 O 3 ! 2 observationsof four
R.F. Shipman
et al.:starformThe physical conditions in IRDC clumps
ing clumps from two di↵erent clouds G28.34+0.06 and G11.110.12.
Thereveals
active clumps
are chosenclouds
based (Marseilleon their brighttion also
many foregroundet al.submm
2010).
continuum(Rathborne
et
al.
2006)
while
the
quiescentclumpsWater abundances are sensitive to temperature as the water
subarechosenbased
their prominentline emission
limatesfromgrainonsurfaces(Herpin etNH
al.3 2012)and show(Pillai
very
etal. 2006a)
The sources
have been
relativelyclumps
well studiedincomplexstructures
within massive
starforming.
the literature which allows our observations of the clumps to be
The goal of this paper is to present the results of Herschelplaced into an existing context.
HIFI and APEX observations of a sample of IRDCs clumps at
Figures
1 and 2 show
images ofthese
the results
two IRDCs
di↵erent
evolutionary
stages850
andµm
to compare
with
2
.
The
two
clumpsin
G28.34+0.06
arelabelled
G28-NH3 forother starforming regions observed by HIFI. This paper presents
17 on NH emission (Pillai et al. 2006a)
thechosen basedo-Hposition
2 O 110 ! 101 and C O 3 ! 23 observations of four starformand
G28-MM
fortwothedi↵erent
positionclouds
chosenbased on theing clumps
fromG28.34+0.06
andstrongest
G11.11submm
emission(Rathborne
et al. based
2006).onLikewise,two
posi0.12. The
active clumps
are chosen
their bright
submm
tions
in G11.11-0.12
Pillai3 (NH3 peak,continuum(Rathborneare
et labelled
al. 2006)G11-NH
while thequiescentclumpset2006a)basedand G11-MM(continuum
Rathborne et al.
areal.chosenon their prominent
NHpeak,3 line emission (Pillai
2006).
et al. 2006a) The sources have been relatively well studied in Fig. 1: SCUBA 850 µm image of G28.34+0.06. The positions ofgeneralclump
properties.
systemictoveliteraturewhich
our observationsof the clumps
be the observed clumps are labelled as ”SCUBA” for the submm
01 the Table
218 1 lists
10allows
01
2 +Clump
locity,are based on NH3 observations while position and ”NH3” for the NH position. The image is colour
placed gas
into temperature
an existing context.
3
clump
mass andtotal
density
are basedon dustFigures
1 and
2 column
show 850µm imagesof thetwo continIRDCs scaled from black (331 mJy/beam) to red (5951 mJy/beam).
2
uum
emission(CareyinetG28.34+0.06
al. 1998; Pillai
al. 2006b).
Kinematic. The
two clumps
areetlabelled
G28-NH
3 for
Every change in colour is roughly an increase of 331 mJy/beam.
distancesto theclumpsareontaken
Carey
et al.et(1998)
and At a distance of 4.8 kpc, one arcminute is about 2 pc
the position
chosenbased
NH3 fromemission(Pillaial. 2006a)arethe radial
velocityof molecularline emissions
and theandbased
G28-MMfor thepositionchosen basedon the strongest
galactic
rotationcurve.
The
positionsandHerschelobservations
submm emission(Rathborne et al. 2006). Likewise, two posi17 are
34
shown
in Table
2. are labelled G11-NH
3 3 (NH3 peak, Pillai
tions
in G11.11-0.12
et al.
2006a)
andG11-MMpeak,stages.
Rathborneal.
The
clumpsspan
a range(continuumof evolutionary
Both et
G282006).
MM
and G11-MM stand out as prominent point sources in PACS
Fig. 1: SCUBA 850 µm image of G28.34+0.06. The positions ofphotometry(Raganet al.clump2012).
Both sources
also assoTable 1 lists
generalproperties.
Clumparesystemicve- the observed clumps are labelled as ”SCUBA” for the submm
ciatedCH3 OH and/or
H2 O maser(Pillai while
et al.
locity, withgas temperatureare basedon NHemission
3 observations
position and ”NH3” for the NH3 position. The image is colour
2006b;
Wangandet al.
2008,
2014)
and both
are seen
to fragment
clump mass
total
columndensity
are basedon dustcontin- scaled from black (331 mJy/beam)
to red (5951 mJy/beam).
into
at high
al. 2009;
uumsubstructures
emission (Careyet al.spatial
1998; resolution
Pillai et al.(Zhang
2006b).etKinematicEverychangeincolourisroughlyanincreaseof 331 mJy/beam.
outflow
observed
for(1998)
G11-MMWang
et al.to2014).
A CHaredistancesthe clumps
taken
fromisCareyet al.
and At a distance of 4.8 kpc, one arcminute is about 2 pc
3 OH
(Leurinial.radial
2007;velocity
Gómez of
et molecularal. 2011; Wanget al. 2014)
are basedettheline emissions
and andthe
G28-MMwasfoundtoharboura”hot-core”
Zhang
et
al.
(2009).
galactic rotation curve. The positions and Herschel observations
These
clumpsare active.
are shownin Table
2.
The
together
withBothG28-MM
The clump
clumpsG28-NH
span a range
evolutionary
stages.
G283 wasofstudied
andto be of astandyounger
evolutionary
based onindeuterMMfoundand G11-MM
out asprominent stage
point sources
PACS
ation
level (Chenet al.
2010).
Neither
identified
photometry(Raganet al.
2012).
BothNHsourcesareisalso
asso3 clump
in
Ragan
et
al.
(2012).
Based
on
their
maps,
G28-NHG11ciated with CH3 OH and/or H2 O maser emission (Pillaial.
3 and et
NH
are
seen
at
SPIRE
wavelengths
but
do
not
stand
out
as
point
2006b;
Wang et al. 2008, 2014) and both are seen to fragment
3sources
at PACS wavelengths
implyinglow (Zhang
temperatures16
into substructuresat high spatial
resolution
et al. (<
2009;
K).
These
not 3active.
OH outflow is observed for G11-MM
Wanget al.clumps2014).are
A CH(Leurini
et al.
2007;et theal. 2011;
Wang etandal.data2014)
andIn Sect.
2 wegiveGómez
details ofobservations
reducG28-MM
was
found
to
harboura”hot-core”
Zhang
et
al.
(2009).
tion. The line fitting procedure and results are presented in Sect.
These
active.of H2 O , N2 H+ , C17 O , and CH3 OH as
3.
The clumps
column are
densities+
wellTheas gasdensities
basedon Nstudied
CH3 OHwith
transitions
areclump G28-NH
G28-MM
2 H andtogether
3 wasdiscussed
In Sect. 5evolutionary
the various stage
dynamic
structuresinand foundin
toSect.
be of4.
a younger
based
on deuterthe
clumpsare discussed
and anevolutionary
sequence
is pro- Fig. 2: SCUBA 850 µm image of G11.11-0.12. The positions of
ation
level (Chenet al. 2010).
Neither
NH3 clumpis identified
posed
andettheareon
comparedwithother starformingin Raganal.H(2012).
Basedtheir maps,
G28-NH2 O results
3 and G11the observed clumps are labelled as ”SCUBA” for the submm

The physical conditions in IRDC clumps from Herschel HIFI observations of H2O
R. F. Shipman et al.

• 若い大質量星候補であるIRDCを観測的に調べる• Herschel HIFI

• o-H2O(110-1 ), H

• APEX

O(1 -1 ), N H (6-5)• N2H+(3-2), C O(3-2),C S(7-6),CH OH• 二つのIRDC: G11.11, G28.34• それぞれMMピークとNH3ピークを観測• 多くの放射＆吸収を検出• 複数ガウシアンでフィットして切り分け• 物理量を導出• 線幅や柱密度• 温度: CH3OH rotation diagram• 前方ガスによる吸収

R.F. Shipman et al.: The physical conditions in IRDC clumps• クランプ共通で見える線
幅が狭くて弱い成分• infall• H2Oのredshiftした吸収• 降着率は~1E-3 M⦿/yr• outﬂow• H2Oの線幅が広い成分• 3天体で見られる• フィラメント• C17O, N2H+, H218Oで見える線幅の狭い成分

G28-NH3

G28-MM

G11-NH3

G11-MM

Fig. 3: Data for HIFI and APEX (C17 O and N2 H+ ) observations for each clumpand transition are shown in blue. Gaussian fits areindicated in red except for the water transition where only the continuum level is shown in red. The systemic velocity of each clump
is indicated with a green line. The complex profiles of the main H2 O line are clearly visible. The full model decompositions of bothemission and absorption are shown in Figs. 4 and 5 respectively.

• IRCD内のクランプは星orクラスターを形成• 進化の流れAn estimation of the gas density is possible based on analysisof the two di↵erent transitions of N2 H+ . The RADEX code isused to determine transition line flux ratios (where flux=width⇥ peak intensity) of N2 H+ (3-2) to N2 H+ (6-5) for isothermal,
uniform density and optically thin clouds. Figure 6 shows thetemperature vs density trends for the N2 H+ line flux ratios from10 to 130. At a given gas temperature, the figure can be used tonarrow the appropriate range of densities.

We have explored the e↵ect of increasing optical depth byincreasing the cloud column density in the RADEX calculations
by an order of magnitude. An increase in the optical depth will
decrease the density by roughly an order of magnitude for agiven temperature and line ratio putting the range of densities
for the clumps between 3 ⇥ 105 cm 3 to 3 ⇥ 106 cm 3 .
These densities are consistent with earlier published resultson N2 H+ in G28.34+0.06 (Chen et al. 2010) and C34 S in G11-

G11-NH3 -> G28-NH3 -> G11-MM -> G28-MM

[FeII] 1.64µm Imaging Observations of the Outﬂow Features around Ultracompact
HII Regions in the 1st Galactic Quadrant
J. H. Shinn et al.

• outﬂowを知る事は大質量星形成を理解する上で重要• [FeII]によるイメージングサーベイと5GHzカタログの精査•

UCHII周りのoutﬂowの有無を調査した•

範囲は0<longitude<90º• 5/237の検出率•

COより遥かに低い, outﬂowはCOの観測のほうが良いかも• [FeII]の起源• PDRでは無くoutﬂowショック起源• 検出率が低い理由•

UCHIIの減光量見積もりエラー•

ショックを起こす物質が既に少ない•[FeII]を出している柱密度が低く、感度的に受からない•

outﬂowが遅くショックを起こしにくい•

outﬂow materialがclumpy

Evidence for Large Grains in the Star-forming Filament OMC-2/3
10

S. Schnee et al.

S. Schnee et al.

• 複数波長の連続波観測からβのマップを出す事が可能• κ=κ0*(ν/ν0)ß -> ダスト成長を調べる• 連続波マップの取得• MUSTANG on GBT - 3.3mm• MAMBO on IRAM30m - 1.2mm• VLA&GBTでNH3マップを取得• ßは領域平均で0.9• 要因は色々考えられるが、ダスト成長を示しているか、オパシティがpower-lawから外れているのだろう• 温度とß反相関もリアルかも

Evidence for Large Grains in the
Figure 2. Temperature (left) and emissivity spectral index (right) mapsof OMC-2/3. The temperature map comes from Li et al. (2013) and issmoothed and regridded to the 11′′ resolution of the 1 mm and 3 mm continuum maps. Temperature is only plotted where the emissivity spectral index
is calculated. The emissivity spectral index in each pixel is derived from theratio of the 1 mm and 3 mm fluxes, using the respective temperature map
values as the isothermal temperature along the line-of-sight. See Sections
3.1 and 4.6 for details. The symbols (plus signs, asterisks, and diamonds)
are as in Figure 1.

spectral index• 微細構造線解析から温度分布へ

c 0000 RAS, MNRAS 000, 000–000⃝

Temperature

Figure 5. (left) Plot of the emissivity spectral index derived from 1 mm andTrigonometric Parallaxes of Star Forming Regions in the Scutum Spiral Arm
M. Sato et al.

oﬀset-oﬀset

epoch-oﬀset

epoch-oﬀset

固有運動+視差

solid:東西方向

上部:視差成分のみ

線はモデルフィット dashed:南北方向• 天の川銀河のスパイラル構造は良く分かっていない• 距離の不定性が主要因• VLBIのメーザー観測(三角視差測定)がキー• NRAOのBeSSeLサーベイから、たて座アームの結果を紹介• 銀河中心近傍のアーム• VLBAでメタノールメーザー2
点、水メーザー4点を観測• マルチエポック、phase reference• 各エポックでイメージングし視差
を導出

下部:残差

six HMSFRs of this study, we assigned 10 other
BeSSeL/VLBA maser sources, whose trigonometric parallaxes have been reported by Xu et al.
(2011), Immer et al. (2013), and Zhang et al.
(2014), to the Scutum arm (Table 1). We thus
have a total of 16 HMSFRs with trigonometricparallaxes for the Scutum arm, which enables us
to examine the geometry of the inner part of the
spiral arm with a substantial sample.
Table 1 lists the measured parallaxes of eachstar-forming region and the estimated proper motion of the central star. Figure 3 shows the locations of all these sources in the Milky Way.
Note that the 4 most distant sources in Fig. 3(G029.95−00.01, G029.86−00.04, G027.36−00.16,
and G025.70+00.04) are the same ones that would
be expected to be most distant based on theirl − v locations in the upper (far-side) sectionof the Scutum loop in Fig. 2. Even their relative distances are correct, since the distance tothe far side should increase with decreasing longitude. This agreement supports the l − v locus
of Scutum proposed by Shane (1972) and others subsequently, and suggests that the arm continues for some distance past its tangency near
the end of the central bar, as recently shown by
Ellsworth-Bowers et al. (2013).

• 対数周期のアームを仮定してピッチアングルを導出• ln(R/Rref)=-(ß-ßref)*tanφ• φ=198º 31º• 過去の見積もりや他アームのアングルより大きい

log(R/kpc)=>銀河中心に近いところなので曲率が大きい？

4.1.2.

Fig. 3.— Locations of 16 HMSFRs in the Scutum armin the Milky Way for which trigonometric parallaxes havebeen measured. 1σ distance error bars are indicated. Six
HMSFRs from this study are shown by red dots and labeled with the galactic longitudes. Ten other HMSFRs
from BeSSeL/VLBA papers are marked by green dots (Table 1). R0 = 8.34 kpc was adopted (Reid et al. 2014).
The positions of the Galactic center (GC) and the Sun
are marked by an asterisk (∗) and ⊙, respectively. The
background is the spiral arm model of the Milky Way by
Taylor & Cordes (1993) with a slightly modified (more axisymmetric) curvature in the near part of the Sagittarius
arm. Spiral arm names are labeled in blue. The grey ellipses show the locations of the central bar (solid curve)
from Gerhard (2002) and Churchwell et al. (2009) andthe long bar (dashed curve) from Benjamin (2005) in the
Galactic center region.

Pitch Angle

For an ideal log-periodic spiral arm, the arm
geometry can be defined by ln(R/Rref ) = −(β −βref ) tan ψ in a Galactocentric reference frame,
where R is the radius from the Galactic center, β is
the Galactocentric
towardthe Sun and
Fig. 4.— Pitch angle of the Scutum arm. The logarithm of Galactocentric
radius (R) in azimuth
kpc units is(0plottedagainstGalactocentric azimuth (β), for R0 of 8.34 kpc. A total of 16 HMSFRs in the Scutum arm with parallax measurements are
increaseswithgalactic
longitude),andR is theplotted with 1σ parallax uncertainties (black curves): red dots for 6 sources from this study and green dots for 10 sources from refother BeSSeL/VLBA papers (Table 1). Red solid line shows the straightline at(spiralarm) fitted bya Bayesian proceduretoradiusa reference
longitudeof βref (Reidet al.
minimize distance of each source perpendicular to the line (see text).
2009b). The pitch angle ψ of the arm is constantfor a logarithmic arm, defined as the angle between
Table 2the spiral arm and a tangent to a Galactocentric
Motions of High-Mass Star-Forming Regions in the Scutum Arm
circle at the same radius. In a plot of the logaSourceR
βUV
W radius, R (in kpc units),rithmof Galactocentric
(kpc)
(deg)
(km s )(km s )(km s )
against Galactic azimuth, β, a log-periodic spiralG005.88−00.39
5.37
3.3−2.5 ± 3.8−11.5 ± 8.5−9.2 ± 5.1
arm
sectiona straight
line, with the slope of
7.8
2.7
± 8.5
8.8 is
± 9.3−11.0 ± 5.7
G011.91−00.61
5.09
5.0
41.1
−3.7 ± 13.0
3.2
G012.68−00.18
6.02
the± 10.3
line proportional
to± 10.9
the tangent of the pitch
6.7
5.6 ± 7.9
6.4 ± 11.9
9.6 ± 9.9
G012.80−00.20
5.54
angle,In± Fig.
4,plot log (R/1kpc) ver5.4
10.6± 7.6 ψ. −1.4
8.7
−5.6we± 3.9
G012.88+00.48
5.93
5.3
16.3
±
10.4
−1.4
±
11.6
−7.1
± 9.5in the same manner as
G012.90−00.24
5.98
sus β for the Scutum arm,ß (deg)• たて座アームの運動は特異• non-circular• 銀河中心に向かう速度が大きい
s

+0.23
−0.25
+0.33
−0.42
+0.21
−0.23
+0.30
−0.39
+0.26
−0.32
+0.21
−0.22
5.91+0.24−0.27
4.62+0.34−0.42
5.01+0.28−0.33
4.51+3.00−0.68
3.88+1.99
−0.08
+0.25

+0.4
−0.3
+1.9
−1.1
+0.6
−0.4
+1.5
−0.9
+1.0
−0.6
+0.5
−0.4
5.5+0.8−0.5
11.8+2.9−1.6
11.8+2.1−1.4
100.9+20.1−28.5
71.4+33.3
−25.4
+13.3

s
−1

s
−1

s
−1

• 中心のbarに近い為 or 太陽近傍のmotion(LSR
G012.90−00.26
G013.87+00.28
G016.58−00.05
G025.70+00.04
G027.36−00.16
G029.86−00.04

frame)が正しくない
4.28

46.3

18.5 ± 10.4
7.9 ± 20.8
21.3 ± 8.0−33.7 ± 38.5
−82.5 ± 50.2
7.3 ± 22.8

−1.1 ± 12.1
−12.3 ± 35.1
−8.6 ± 9.5−19.5 ± 21.5
−48.4 ± 27.4
−16.0 ± 7.8−0.1 ± 9.8−9.0 ± 37.6
5.1 ± 6.0
−4.5 ± 8.2−1.6 ± 6.4−1.5 ± 4.0

8

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