THE STAR FORMATIONNEWSLETTER
No. 285 (#1-5)
塚越崇(茨城大)7 Oct. 2016

1

1. Sub-0.1′′ optical imaging of the Z CMa jets with SPHERE/ZIMPOL
2. Prompt planetesimal formation beyond the snow line
3. Turbulence in the Ionized Gas of the Orion Nebula
4. A Magnetic Ribbon Model for Star-Forming Filaments
5. Search for Compact Stellar Groups in the Vicinity of IRAS Sources2

Sub-0.1′′ optical imaging of the Z CMa jets with SPHERE/ZIMPOL
S. Antoniucci et al.
00
4

S. Antoniucci et al.: Sub-0.1 optical imaging of Z CM

• ZCMa• Herbig + FUoriの連星• それぞれにjetあり• VLTのSPHERE/ZIMPOL
でjetの根元を分解• Ha, [OII]• 0.03” resolution

Fig. 1. (a): Continuum-subtracted (top) and deconvolved (bottom) [O i] images of theshows only the di↵use emission as reconstructed by the MC-RL algorithm. The centroid
main knots of the FUor jet are indicated. Areas heavily corrupted by artefacts around the stfor the H↵ observations. (c): Fitted transverse positions of the peak of the FUor jet spatial
the exciting source for 3both the [O i] continuum-subtracted (blue squares) and deconvolv
red pluses indicate the measured profile width of the two images. The solid green line is

S. Antoniucci et al.: Sub-0.100 optical imaging of Z CMa• Herbigからはwide angle wind
• jetは感度的に見えていない• FUOriからはcollimated jet• wiggling: 1.2au, 4.2yrの連星(unresolved)に対応• mass ejection: M’jet=10-8~10-6 M⦿/yr,
M’jet/M’acc<0.02

4acted (top) and deconvolved (bottom) [O i] imagesof the Z CMa system. The deconvolved ima

f
y

,
n
s
l
-

)

g
5
e

l

outside the snow line, fall in the Epstein drag regime.
For particles of material density ρm and radius s, the2 dimensionless stopping time is,
Prompt planetesimal formation beyond the snow lineπ ρm2005) results
in
the
rapid
formationτ=s.of planetesimals. (3)sP.J. Armitage 2etΣ al.
Planetesimals formed from the streaming instability
The aerodynamic
drift (Johansen
rate is (Takeuchi& Lin2002),
et al. 2007,2012;
are characteristicallylarge
微惑星形成とradial
driftの関係
Simon et al. 2016), and toτ a−1good
approximation theys vr,gas − ηvK, in place. (4)vr =will be immune(タイムスケール)をモデル化
to aerodynamic
drift and stayτs + τs−1
The above sketch defines a model for “prompt” planwhere
vr,gas is thatthe radial
velocity
the gas,K is theetesimalformation
does notinvokeofhelpful
butvless
2
Keplerian velocity,
and(suchη = as−(h/r)− 3) beis a pawell-understoodprocesses“lucky(β/2growth”
the degreeof radial pressure
supportyondrametermaterialmeasuring
barriers, andlocal concentrationin zonal
forvortices
the gasor
disk.
We takethe limit
τs is small‘simpleなモデル’
flows,
persistentparticletraps,wherefor reviews
and |vr,gaset|≪
For a steady-stateparticles vK |.Armitageal.|ητ
2014;2015). Itradial
is obvisee Johansen
planetesimalwith accretion
rate Ṁ
surfaceindensity
Σp ,ouslyflow,
incomplete.
intentis ptoandhighlight
the simmetallicallyとstoping
timeの関係• Ourformation不可
plest possible modelΣthe critical
role
of
the
competition
3αṀp −1pτ .
(5)
between radial drift and=planetesimal formation.
Σ2(3 − β/2) Ṁ s
2. A MODEL FOR PLANETESIMAL FORMATIONThis expressionshows
how the solid-phase “metallicity”- tdrift• tgrow
FROM
DRIFTING
SOLIDSZ ≡ Σp /Σ depends upon the disk properties and ratio
Considera gas accretion
disk aroundstarof τmassM∗of solid torates,a forgivens.
that isWriting
parameterizedby physics
the accretionan-abovethe familiarof radialrate
driftṀin, thegularform
momentumtransportα
is useful because
the eﬃciency
conditionsparameterfor the streaming
連続の式から定常解を調べる(Shakura& Sunyaev
1973),and centralcollapse
temperature• leadinstabilitytoto gravitational
and plan−β
profileT (r) ∝ canr ,alsowithbeβ expressed
a constant.
etesimals
asIn
ansteady-state,allowed region in

•

•

Ṁ
,
(1)
3π
where the eﬀective viscosity is given in terms of thesound speed cs and angular velocity Ω by ν = αc2s /Ω,
and c2s = kB T /µmH . Here kB is the Boltzmann constant, µ ≃ 2.2 is the mean molecular weight, and mHνΣ =3Z-τs space. A polynomial fit to the simulation results ofandmassive
ofmetallicity as1 ,Carrera etal. (2015)
gives thebelt
critical• Broadlog Zcrit = 2.84 ×with
10−2 (logτs )4 + 0.11(logτs )3 edge
planetesimals
a sharpouter+0.38(log τs )2 + 0.6(log τs ) − 1.52. (6)particle diffusionとvapor• turbulent
Equation (5) then defines a line in that space that either does or does not admit planetesimal formation.

diffusionを入れた場合
Whether planetesimal formation is possible depends on•

α and the particle flux but not on the actual physical size to lineの外側でcm-sizeのダストがwhich solids grow. Plotting these tracks=>snowin Figure 1 for α = 10−2 (at the high end of observationally estimated values, Hartmann et al. 1998), we
堆積する傾向があるが微量なので微惑星
see that if Ṁp /Ṁ = 10−2 the metallicity is too low atany τs to allow prompt planetesimal formation. Plan形成に影響なし
etesimal formation is only possible in the outer disk if
Ṁp /Ṁ ! 3 × 10−2 , i.e. for relative accretion rates that=>入れない場合のモデルはrobust
exceed the fiducial dust-to-gas ratio of 10−2 . This implies that planetesimal formation would occur while theglobal average of the dust-to-gas ratio was decreasing
with time, presumably early on.
A steady-statemodel will only be valid, even approxsimple
modelの場合、prompt
imately, if particles can grow to a size set by material barriers before they drift significantly (otherwise
planetesimalformation(>snow line)はか
they will be in the regime of “drift-limited growth”,
Birnstiel et al. 2012). Birnstiel et al. (2012) estimate
なり限られた物理条件が必要
that a particle grows from radius s0 to s on a time scale,
! "
1 Σ
s
t
≈
ln
.
(7)
grow
=>このplanetesimalΩ Σp formationは円盤初
s0
Adopting s = µm, s = mm, and β = 0.5, we find that0期に起きる必要がある(?)and the drift time
the ratio between the growth time

tdrift = r/|vr | is,# α $−1
tgrow≈ 87
tdrift
10−2

!

h/r
0.05

"2 %

Ṁ /Ṁp
10

&

τs2 .

(8)

Turbulence in the Ionized Gas of the Orion Nebula
S.J. Arthur et al.
10• Orion領域のionized gasの乱流構造を調査• 理論モデル論文(Paper I)の検証• 可視光スペクトル(Garcia-Diaz+2008)から
速度構造を抽出 (Paper I提唱)
• Velocity Channel Analysis(空間構造)• Second-order Structure Functions(速度構造含む)

S. J. Arthur et al.

Turbulence in the Ionized Gas of the Orion Nebula13• (特に)Region II(8~22asec)に着目し、
VCA/Soのslopeを比較

4000=>Kolmogorov slopeとconsistent3000• 0.02-0.05pc相当• thick/thinどちらでもVCAのslope
breakは変化なし

A
2000

1000

0
−40

=>速度ではなくemissivity fluctuation
14

−20

0

20

40

60

80

100

Heliocentric velocity, km/s

S. J. Arthur et al.

• イオン化度↑で速度分散とslope↑pressing the e↵ects of scattering on the line widths is to use
the Full Width at Half Maximum (FWHM) instead of the
RMS width. Even a weak component with a velocity that isvery di↵erent from the line centroid can have a large e↵ect
on the los due to the (u ū)2 dependence of the second
velocity moment. But the same component will have almostno e↵ect on the FWHM if its amplitude is less than half that
Figure 10. Joint distributionof We
moment-derived
ve-e↵ecof the line core.
therefore
define line-of-sighta FWHM-basedpln 2 ⇡ FWHM/2.355,
with thetive mean
as ˆlos= FWHM/
locity dispersion andvelocity.
The 8thermal,
instrumental,
constantofproportionality
being
chosen
so
that
ˆ
= los
los
and fine-structure contributions to the velocity dispersion havefor a Gaussian profile. The results are shown in Figure 12,
been subtracted in quadrature. Mean velocities are shown with
where it can be seen that nearly all the lines show substan1
respect to the systemic
velocityinofthe
theline-of-sightstellar cluster(+7 kmtial reductions
line widths
withs respectin the Local Standard
Rest frame, or+25ofkms 110.inThetheexception
Heto theofmoment-derived
valuesFigureis H↵,etwherelarger thermal broadening means that ˆlos isliocentric frame, Tobin
al. 2009).
still contaminated by the back-scattered component due toblending. Note that ˆlos is insensitive to not only the scattered component, but also to any other weak component
3.4 Line-of-sightthat isvelocity
not blendeddispersions
with the line core. For the low-ionization• 22asec(0.05pc)はVCA/So共に特徴的なスケールとなっている=>密度揺らぎスケールかつ乱流のdriving scale

Table 5. Comparison of structure function slopes
2

Reference

Ion

Method (km2 s

O’Dell & Wen 1992
This paper
This paper
Wen & O’Dell 1993
This paper
This paper
Castañeda 1988
This paper[O i][S ii][N ii][S iii]
H↵[O iii][O iii][O iii]

Mean
Mean
Mean
Comp A
Mean
Mean
Comp A
Comp A

3:
5.4
5.6
13.8
9.4
10.2
13.7
15.3

2)

Range
(00 )
6–85
7–32
8–22
5–20
8–22
8–22
3–15
3–15

Slope
0.68:
0.80 ± 0.12
0.82 ± 0.11
0.92:
1.17 ± 0.08
1.18 ± 0.09
0.86 ± 0.05
0.73 ± 0.05

the structure function slope increasing with ionization. The
most directly comparable methodology to our own is thatof O’Dell & Wen 1992 who used the flux-weighted mean velocity of the [O i] 6300 line. The fitted range of 600 to 8500
extends to larger scales than in studies of other lines, andthis biases the slope determination towards lower values due

A Magnetic Ribbon Model for Star-Forming Filaments
S. Auddy et al.

Magnetic Ribbon M• ~0.1pcスケールの分子雲フィラメ

To observer

ントの構造解釈• 磁場があるとシリンダーではなくシート状に収縮するはず。また観測での0.1pcに対する大きい分散も説明できない• シート状フィラメント(ribbon)とviewing angleによる解釈• semi-analytic model2H sec θ = L csc θZL cos θθ

2H sin θθ = θcritx• local pressure balance in a subcritical• thermal pressure is neglected

To observer

2H

L(Kudoh&Basu2014)
cloud is assumed

y

⊗

Z

L

To observer

2H sin θ

θ > θcrit
2H

L cos θL csc θ

x

⊗

y

H

8⇡

= H0 ⇢0⇢v0 t0 +

8⇡

L ,

(1)

0

Observed Ribbon Width [pc](2011) to determine the mean filament w
By using equation (7) in equation (5), we can express
be used to estimate a mean spacing betw
re vt0theis theflowfilament
speed. asAssuming that the
finalnonlinear
width of thewhich we identify with L0 in our model.
リボンの構造•" ✓ into# 1
has adequate time to settleequi◆2hydrostatic
Arzoumanian et al. (2011) identifies 27 fi
vt0はnonlinear flow speed
vt0
dust emission map of the cloud IC 5146.
um along the z -direction,the (8)
L = L0 2 the half +thickness
1. v ofは初期アルフベン速度A0vA0same map to make an approximate measu
d iscs
For each of the 27 filaments, we identify aAnalysis of Zeeman
measurementsof the magnetic
H=p(2)fieldthe length (spine) of the filament. Then
2⇡G⇢ by Crutcher (1999) showsin molecular clouds presented
ament we measure the distance to the nthat the turbulent
line
width
is
comparable
totheAlfvén
L0は観測されたフィラメント間の距離と仮定(~1.0pc)• if the ambipolar di↵usion time istzer 1942). Now
point of another filament. We obtain a se
speed (Basu 2000). If we make the plausible estimate
distance measurements (eliminating douber thantime(Kudohto&theBasu2014),
thatthethe•compression
flow
speed iscomparable
Alfvénspeed,
Viewing
angleに対して3つのケース(φは無視)
cases where two filaments areAuddy
mutuallyetfreezing
compression,i.e., 6i.e., vist0valid' vA0during, the filament
width becomes
nearest neighbor). We ignore e↵ects of aB
BL0' L0 /3.
angle i in this analysis, which could mean t
1.00=.
(3) (9)sured distances are less than the actual dt⌃⌃0viewing
angleをランダム• illustrates the fact that the final width of a
The resultw
factor sin i. Our measured filament spacin
the surface⌃ = 2⇢H,
equation(3)thecanmedium.
befilamentdensity
is independentof thedensity ofequate with L0 ) have a minimum value 0sにとり、観測される柱密
it is a fraction of the initial length scale L0 of
ittenInsteadasimum value 2.2 pc, a median value 0.9 pcg
0.10
the compressedregion.B0
度v.s.幅の関係を調査
Brvalue 1.0 pc.
=
.
(4)
1
1
l
2
2
⇢
2.3.2.
Estimatefromstarformationt⇢=> 観測とよく合う
0Mean Ll
θ=0
Another way to constrain L0 is throught
B =B
θ= 90
g equation (2) and equation (4) in equation (1) andmation timescale in molecular clouds. Sincp
0.01
some simplifications we get
is widely accepted to be in the range 1-3 s
23
10 21
10 22
✓ νt0 ◆1/2✓
◆2
Stahler 2000,
2002;−2Hartmann10 2001),
and seObserved Column Density [cm ]⇢vt0
is often coordinated along a filament over tc
=2
+ 1,
(5)
⇢0
vA0
Figure103. Apparent ribbon
col- empirS
we canwidth
placeLobs
an versus
upper observed
limit (again
Z

2H

L

obs

°

°

exceeds 20% [22]. At distances exceeding 1’1, SFE | 5%,star formation regions [2]. In the JH/HK diagram the cen
Herbig stars.
Search for Compact Stellar Groups in theAe/Be
Vicinityof IRAS Sources
N.M. Azatyan et al.
2.2. Data. We have used images, coordinates, and photometric data from the GPS UKIDSS• データベースを元にしたクラスター探査

survey [10]. The astrometric accuracy and resolution of the survey is ~0.1 arcsec/pix. For purity• IRAS source近傍

18m.05. In addition, we have excluded those objects that have a probability greater than 90%• 19天体of being the result of various kinds of noise patches or image defects, as well as objects withof the sample with respect to the photometric parameters we have chosen objects with K <coordinates that match densifications of directed outflows revealed in H2 lines [8].

• GPS UKIDSS survey + Spitzer

We have also used mid IR data from the Spitzer GLIMPSE survey [11]. There theresolution of the images is 0.6 and 1.2 arsec/pix. The photometric completeness of the sample
GLIMPSE surveyのデータを使用
for the different ranges varies from 12m.0 to 14m.0.

• クラスター同定• color-color diagram作成• luminosity function• 15/19でクラスター検出• 4天体は初検出• サイズや星密度はバラバラ• “initial stage of research”
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Fig. 1. Radial distribution of the stellar density of a cluster.