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

global average
was decreasing
cal size ofto thewhichdust-to-gas
solids grow.ratio
Plotting
these tracks
ure 1. Upper panel: the radial dependence of the particle−2
outside
snow
line,
fall
in
the
Epsteindragregime.
withthetime,
presumably
early
on.
mass accretion rate. Lower panel: the radial dependence ofin Figure 1 for α = 10
(at the high end of obserFor particles
ofmaterial
densityρandradiuss,theplanetesimal formation. Solid curves refer to models with
A steady-statemodel willonly
be valid, eteven
approxal. 1998),
wevationally estimatedvalues,m Hartmann
2 dimensionlessṀpthe/Ṁ =snow
1, dashedline
curves models with Ṁp /Ṁ = 0.1. The
seethat
if ṀtimeṀ =10−2 thetoo matelow
atis,grow
p /canimately,
ifstopping
particles
to metallicity
aformation
size setisby
Promptplanetesimal
beyondassumed planetesimal formation time for the mm-sized parany τsbefore
to allowthey
promptplanetesimalformation.
Planπ ρmrial barriers
driftsignificantly(otherwise
ticles cases increases from right to left (K = 102 , 103 , 104 ).
in
the
rapid
formationofplanetesimals.
τ
=
s.
(3)
, 2005) results
etesimalformationisonlypossible
in
the
outer
disk
if
s
P.J.willArmitage
they
be in the regimeof “drift-limited growth”,
2et
Σ al.
−2
For cm-sized particles only the case with K = 102 is plotted.
Ṁp /Ṁ !
3 × 10, streaming
i.e. for relative
accretion rates thatfromtheinstabilityn Planetesimals formed
et al.drift2012).
et ratioal.
(2012)−2 estimate
The Birnstiel
aerodynamic
rate Birnstiel
isdust-to-gas(Takeuchi& Lin2002),
exceedthefiducialof2012;10. This imet al. 2007,
large(Johansen- are characteristically
微惑星形成とradial
driftの関係
that a particle
grows
from
radius
s
to
s
on
a
time
scale,
0pliesand
thattoplanetesimalformation wouldthey
occur whilethea−1goodapproximations Simon et al. 2016),
Figure. The radial
the modelsshownFig-−1 . For
!
"
τ
v
−
ηv
r,gas
K
scale 2within
the projection
allowed ofregion,
tform
=inKΩ
s
globalvraverageof thedust-to-gas
ratio was decreasing
,
(4)
=
1
Σ
s
ure
1.
Upper
panel:
theradialdependenceoftheparticlewillbeimmunetoaerodynamic
drift
and
stay
in
place.
(タイムスケール)をモデル化
ln on. .
(7)
τ≈s + τs−1
Z <accretion
Zcrit rate.
we smoothly
the rate of
by a facwith time,tgrow
presumably
early
mass
Lower panel:suppress
the radial dependenceΩΣs
The above sketchA defines
a
model
for
“prompt”
planp
0
planetesimalformation.
Solid
curvesrefer
tomodelswithsteady-state model will only be valid, even approxtor exp[(Zcrit − Z)/0.1Zcrit]. This roll-oﬀ in the planwhere
vr,gas is thattheradialvelocity
ofhelpfulthe gas,vless
K is theformationdoes
notinvokebutṀp /Ṁ = 1, dashed curves models with Ṁp /Ṁ = 0.1. The- etesimal
imately,
if particles
can andgrowsizewesetfindby mateAdopting
s0 = µm,
s = mm,= a0.5,that
2 βtoetesimalformation
rate istime
imposedfor numericalconassumed
planetesimalformation
for the mm-sized
parKeplerianvelocity,andη=−(h/r)
(β/2
−
3)
is
a
paprocesses
asthey“luckygrowth”
be-driftbarriers(suchbeforedrift
significantly(otherwise- well-understood
2
3
4
the ratiorial
betweenthegrowthtime
and thetime
ticlescases increases
to left (K =plausible
10 , 10 , 10that).
venience,
thoughfrom
it right
is physicallysomemeasuring
theofradialofpressure
support2
they willand
bedegreein theconcentration
regime“drift-limited
growth”,material
barriers,
local
in
zonal
- yondrameterFor
cm-sizedparticles
only
the
case
with
K
=
10
is
plotted.
tdrift = r/|vr | is,planetesimal formation persists at metallicities just befor
the
gasor
disk.
We
take
the limit
whereτs(2012)
is& smallBirnstiel
et al.
2012).
Birnstiel
et al.
estimate‘simpleなモデル’
flows,
vorticespersistentparticletraps,
forreviews%!from radius"2 sradial
low the nominal threshold. We consider icy particles$−1
that
asparticle
grows
tosobviona time scale,and |vr,gas|≪
v#K |.αFora steady-stateparticle
0 is
tgrow
h/r
/"Ṁ
et
al.|ητ
2014;
Armitage2015). !
ItṀpl see Johansen
2≈ 87 rateτps, . (8)scaleallowedtform Ṁ= KΩ
(ρplanetesimal= 1 g thecm−3) in aregion,
disk with= −1
10.−8ForM⊙ yr−1 ,m withinflow,
withaccretionandsurfacedensity
Σ
−2 Ṁ
1
Σ
s
p
ously
incomplete.
Ourintentistohighlight
in
the
simt
10
0.05
10
drift• metallicallyとstoping−2 we smoothly suppress the−0.5
tgrow ≈
lntimeの関係.
(7)
Zα formation不可
< 10
Zcrit
=
, and T = 150(r/3 AU) rateK.by a facΩΣsp0
plest possible modelΣthe critical
role
of
the
competition
tor Figureexp[(Zcrit− Z)/0.1Z
roll-oﬀ
in theTheplan-key point
3αthe trackṀp −1
crit]. Thisp
1 shows
solutionsto thismodel.
theseparameters,by0.5,
equation(5)
τsgiven. β=(5) that
between For
radialdrift
and=planetesimalformation.
Adopting
s0 = µm,
s = mm,and
we findetesimal formation rate is imposed for numerical conΣ2(3 −streaming
β/2)Ṁ regionis that, within the allowed region, the time scale forintersects
theallowedatτ≃
0.035,
at
s
the
ratio
betweenthegrowth
time
and
the
drift
time
)
venience, though it is physically plausible that some
2. A 2MODELFOR
PLANETESIMAL
Figure. whichTheradialthemodelsshown
in Fig- of theprompt planetesimalformation
scales withradius as
point
thegrowthtime
isFORMATIONa smallfractiontdrifttprojection=r/|v| is,ofthe
rdriftThis
expression
shows
how
solid-phase“metallicity”
planetesimalformation persists
at metallicities
just bet•
growure 1. Upper
panel:
theradialdependenceoftheparticle
−3/2
%
&
FROM
DRIFTING
SOLIDS!properties"2
r the, nominal
which canbe comparedto icytheparticles
time scale fortime.
This ordering
is−1not,
however,robust, andlowthreshold.
We consider# the$Z ≡driftΣp /Σ
depends
disk
g mass accretion
tgrow uponh/r dependenceṀ /andṀp ratio
αthe
rate.
Lower
panel:
radialof2≈around
87 may(8)(ρg cm−3) in a=disk
with Ṁ= 10−8M⊙ yr−1 , a critiradialtdriftconst.
Thereis therefore
for formation.
larger
growthslowerConsidera gas τaccretion
a for
star
of
m = 1drift
5 planetesimalof solid
torates,
given
τmass
stdisk∗ τs .
s . drift.
refer
tothanmodels10−2be0.05
10MwithdriftSolid curves
−2
−0.5
= 10
, and
T =(dashed
150(r/3
AU)
Figure 1. αcal
Analytic
tracksin K.
Z-τs spaceforradiuswithinwhichlines)
planetesimalformation
domiisWriting
parameterized
the ofaccretion, thean-TheTomodel
the byimpactplanetesimalformationon thee Ṁpthatthe familiar
physicsof radial
drift
in0.1.
above
/Ṁ =
1,
dashedcurvesmodelswithṀp /rate
Ṁ =Ṁ
−2
1
showssolutionstothismodel.
The
key
point
FigureFortheseparameters,
the
track
given
by
equation
(5)
a
disk
with
α
=
10
,
β
=
1/2,
and
variousratiosofthenates, whereas outside radial drift leads to particle pilegularformmomentumtransport
eﬃciency
parameterα
trackin Z-τ
space,thewetimeintegratethecontinuity
equationassumedplanetesimalfor themm-sizedparis useful
becauseconditions
forthe streaming
sformation
is gasthat,accretion
within theallowedregion,the time
2
3 τs ≃
4 0.035,
intersects
theallowedstreaming
regionat
atparticletorate.
Particles
evolvefrom scale
lowerforcases
increases
from
right
to left(K collapse=temperature
10 , 10and, 10plan).
連続の式から定常解を調べる(Shakura& theSunyaev
1973),
andcentral& radial
Chiang2004).
Ifred
thecurve
streaming
instabilityup (Youdin•whichforparticles,- ticlesinstability
to
lead
to
gravitational
right
to
upper
left
under
drift.
The
defines
prompt planetesimal formation scales with radius as
point the growth time is a 2small fraction of the−β
cm-sized
onlycase with
10 is plotted.
T (r)particles∝ canr drift
,also
with
βthe
a 1constant.
In
steady-state,gravitational
collapseessentially−3/2 tol Forprofile
etesimalsbeexpressed
asK
an=allowedin andthe approximate
regionwithinwhich
promptplanetesimalrleads, whichcanbe compared
to on
thean
time
scale for dynam∂ordering
Σp region
∂Σ
time.
is
not,
however,robust,p This2+(rΣp vr ) = − than,drift.
(9)
formationradialis possible.
The
particles)
ical
time
(K
∼ 10(for,Thereasmm-sized
is isfoundin simulations,e.g.
drift scale
tdrift= greenconst.
thereforea critifor larger∂t τs growthrṀ∂r may be slowertform
andblue(cm-sized)
curvesshow
tracksthat
include
a
sink
,
(1)
νΣ =
cal
radiuswhich
planetesimalformation
domi- skirt theSimon
et withinal. 2016),
then
the derived
Z-τs tracks
To modelthe impactof planetesimalformationon the−13πterm for planetesimal formation, for K = 102 (solid curves),
to
find
a
steady-statesolution.
We
assumethatplanscale within the trackallowedregion,
t
=
KΩ
.
For
form
nates,
whereasoutsideradialdrift leads
to particle
pilein Z-τs space, we integrate the continuity equation
lowerboundaryof theallowedSlower
K = 103 andK=
104 (dashedcurves).
Theregion.
disk model
has planetesiwheretheeﬀectiveviscosity
is
given
in
terms
of
the
etesimals
form
onamultipleKofthedynamical
time&
Chiang2004).
Ifthestreaming
instability
up(Youdin
Z < Zcrit we smoothly
suppress the rate by a facfor the particles,malscales AU)
lead−0.5
to clearly
defined regions2Ṁ = 10−8leadsM⊙ formationyr−1
and Ttime= 150(r/3
K.
sound
speed
c
and
angularvelocity
Ω
by
ν
=
αc
/Ω,
to
gravitational
collapseonanessentiallydynams Z)/0.1Zs
tor exp[(Z
−
].
This
roll-oﬀ
intheplan1 ∂
Σpcrit
crit∂Σpwhere first radial drift
and then planetesimal formation
2
(rΣ
,
(9)
p vr ) = − conical
time
scale
(K
∼
10
,
as
is found in simulations, e.g.
and c2s = kB T /µmH . Here ∂t
kB +isr the
Boltzmann∂r for numericaltform conetesimal formation
rate
is
imposeddominate.
We
remarkthatmodelswithskirt
large
5
Simon et al. 2016), then the derived
Z-τs tracks
thevalues ofstant, µ ≃1 2.2
istothemeanmolecular
weight,andmCarreraet al.
(2015)
model diskswith negligibly
weak intrinH
macroscopic
particles,
that
have
grown
by
coagulation
aphysicallysteady-statesolution.
We
assumethat planvenience, sic
thoughitfindisSignificant
that
some
K havean alternate
interpretationin termsof stochasturbulence.
levelsplausibleof turbulence
would
impede setlower
boundaryof the allowed
region. Slower
planetesi-

f
y

•

•

3
Z-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 /4Ṁ ! 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 imZ-τs needed
space.
Aprompt
polynomial
fitwhile
to formathe
pliesthethatplanetesimalformation
would
occur
the simulation results ofmetallicity
forplanetesimal
Figure 2. The 1radial projection of the models shown in Fig2
tionaverage(withCarrera
Kof= the10 )dust-to-gas
1%al.
of thetime.
globalratiogives
was decreasinget(2015)
the critical metallicity
as panel:,ure 1. Upper
the radial dependence of the particle
2 shows theearly
solutionsFigurepresumably
with time,on. as a function of radiusmassaccretion
rate.
Lower
panel: the radial dependence of−2
log Zwill=
2.84
× 10even(logτs )4be+ 0.11(logτs )3 formation. Solid curves refer to models withplanetesimal(notionallyextendedintoonly1 AU,though
this
would
crit
A
steady-statemodel
bevalid,
approxsimple
modelの場合、promptṀp /Ṁ = 1, dashed curves models with Ṁp /Ṁ = 0.1. The
inwardif ofparticles
the physicalsnow toline).
ForsetṀpby
0.1
2 mateimately,
can growa size(6) formation time for the mm-sized par+0.38(logτs/)Ṁ
+!0.6(logτassumeds ) − 1.52.
planetesimal
mm- tocm-sizedplanetesimalformation
rial andbarriers
beforetheyparticles,
drift significantly(otherwise
ticles cases increases from right to left (K = 102 , 103 , 104 ).
planetesimalformation(>snowline)はか
in
the
outer
disk
is
an
eﬃcient
sink
for
the
radialflow
they will be Equationin the regime“drift-limited(5) ofthendefines agrowth”,
line in that For
space
that
ei- only the case with K = 102 is plotted.
cm-sizedparticlesofsolids.
The
mass
accretion
rateofsolidsreaching
Birnstiel et al.
2012).
et al.
estimate
therdoesBirnstielor doesnot(2012)
admit
planetesimal formation.
the
snow
line
is
substantiallyreduced,byafactorthat
なり限られた物理条件が必要
that a particle
grows
from
radiuss
0 to s on a time scale,
Whetherplanetesimalformationis possible depends on
ranges from several to more !
than"an order of magniscale within the allowed region, tform = KΩ−1 . For
1
Σ
s
and
the
particle
flux
butformation
not(7)on theZ actual
tude. As αa result,
the rateof planetesimaltgrow ≈
lnformationは円盤初.
< Zcritphysiwe smoothly suppress the rate by a fac=>このplanetesimalΩΣspat any given
on 0whether
cal radius
size depends
to which
solids planetesimals
grow. Plotting
trackstortheseexp[(Zcrit− Z)/0.1Zcrit]. This roll-oﬀ in the planhave
formedfurther
out
in
the
disk.
For
most
of
the
−2
Adopting s0 in= µm,
s = mm,andαβ =
= 0.5,
Figure1 for
10 we find(at that
the highetesimalend offormation
obser- rate is imposed for numerical con期に起きる必要がある(?)
modelsconsidered
the
sweet
spot
for
planetesimal
forthe ratio between the growth time and the drift time
thoughHartmannetvenience,al. 1998),
we it is physically plausible that somevationallyestimatedvalues,mationiswelloutside
the
snow
line,
in
some
cases
as
tdrift = r/|vr | is,planetesimal formation persists at metallicities just bethatṀp /Ṁ10−2
is too low at
far out assee
30-40
AU. ifPlanetesimals
predicted
to be%= are&the metallicity
!
"
2
low
the nominal threshold. We consider icy particles#$−1
tgrowdown any
h/rṀ /Ṁ
ατ a broad
laidacross
range
of
radii
for
high
particle
p
2
to
allow
promptplanetesimalformation.
s≈ 87
τs . (8)6(ρm = 1 g Plancm−3 ) in a disk with Ṁ = 10−8 M⊙ yr−1 ,fluxes,lower fluxes
distributdrift whereas
10−20.05 yield a narrower
10−0.5
etesimal formation is only possible in the
αFigure=outer10−2
andTif= steady-state
150(r/3 AU)
3., disk
ExamplesurfaceK.
density profiles ofk [parsec ]10

k [parsec ]100

10[NII] 6583 - Thin

100

[NII] 6583 - Thick

II
III
IV
TurbulenceI in the
IonizedGas of the Orion Nebula
S.J. Arthur et al.
k3P(k)

1.00

10• Orion領域のionized gasの乱流構造を調査

S. J. Arthur et al.

Second-orderl [parsec]
0.01

0.10

•

γt2 = −2.30±0.08
γ = −2.58±0.09
理論モデル論文(Paper
I)の検証 γT2T3 = −4.14±0.13
γt3 = −4.14±0.11• 可視光スペクトル(Garcia-Diaz+2008)から
Hα 6563 - Thin

1.00

Hα 6563 - Thickl [par0.01

0.10[NII] 6583
σ2vc = 5.62 km2s-2

Hα 6563σ2vc = 9.37 km2s

1.00

k3P(k)

8

S2 (l )

速度構造を抽出 (Paper I提唱)

0.10

• Velocity Channel Analysis(空間構造)

m2D,m = 0.82±0.11
m2D,u = 0.75±0.07
m2D,l = 0.93±0.04

0.10

S. J. Arthur et al.

γ = −2.70±0.12
γ = −2.79±0.15
Structure Functions(速度構造含む)• Second-orderγ = −4.40±0.12
γ = −4.26±0.16
k [parsec ]k [parsec ]
VCA
10
100
10
100
t2

T2

-1

t3

T3

[OIII]
5007--Thin
Thin[NII] 6583
I

0.01

-1

[NII] 6583

[OIII]
5007--Thick
Thick[NII] 6583
IIIII

Hα 6563

1.00

IV

k3P(k)

S2 (l )

1.00

m

0.10

0.10
−2.30±0.08
γt2 = −2.53±0.08−4.14±0.11
γt3 = −4.03±0.11−2.58±0.09
γT2 = −2.82±0.10−4.14±0.13
γT3 = −3.68±0.25

0.01
1

1.00

0.01
0.10
Hα 6563 - Think [arcsec-1]

0.01
0.10
Hα 6563 - Thickk [arcsec-1]

10
l [arcsec]

100

1l [arcFigure8. velocity
Second-order structure functions for the velocity centroid i
Figure 5. Compensated power spectra k3 P (k) of the velocity channels for the [N ii], H↵ and [O iii] 7emission lines
for thin (32slices) velocity slices and thick (one velocity slice) in the velocity range 40 to 70 km s 1 . Four di↵erent slope
regimesandconsidered in their
this work. Top panel (pdf-selected pixels): the structure fu

Turbulence in the Ionized Gas of the Orion Nebula13• (特に)Region II(8~22asec)に着目し、
3000• 0.02-0.05pc相当• thick/thinどちらでもVCAのslope
breakは変化なし

Line corecomponents

2000

Line corecomponents
A + B only

A

B

1000

C

0
−40

=>速度ではなくemissivity fluctuation
14[O III] λ5007 Å
ONC systemic velocity=>Kolmogorov slopeとconsistentLine intensity

VCA/Soのslopeを比較

4000

−20

0

20

Back-scatteredcomponent

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 scat8
tered component, but also to any other weak component
3.4 Line-of-sightthat isvelocity
not blendeddispersions
with the line core. For the low-ionization

Figure
11. Typical [O iii] line proile, showing Gaussian decompoTable 5. Comparison of structure function slopes
sition into three components, following the labeling of Castañeda
2
Range(1988). The red-shifted component C
is dueto back-scattering ofReference
Ion
Method (km2 s 2 )
(00 )
Slope
the blue-shifted components A and B by dust that lies
behind the
O’Dell
&
Wen
1992
[O
i]
Mean
3:
6–85
0.68:
emitting gas. Inset figure is same as lower right panel of Fig. 10,
This paper[S ii]
Mean
5.4
7–32
0.80 ± 0.12
butThisafter
removing Component
profile.
paper[N ii] Mean C from
5.6 the8–22
0.82 ± 0.11

• 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

Wen & O’Dell 1993
This paper
This paper
Castañeda 1988
This paper[S iii]
H↵[O iii][O iii][O iii]Comp A
Mean
Mean
Comp A
Comp A

13.8
9.4
10.2
13.7
15.3

5–20
8–22
8–22
3–15
3–15

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 Mtion oof thZ

2H sec θ• ~0.1pcスケールの分子雲フィラメ

θ < θcrit

To observer

L cos θントの構造解釈L• 磁場があるとシリンダーではなくシート

2H sin θ⊗y2H状に収縮するはず。また観測での0.1pcにx(a)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2H sec θ = L csc θZviewing angleによる解釈

L cos θθ

2H sin θ• semi-analytic modelθ = θcritx(b)• local pressure balance in a subcriticalZ

L

cloud is assumed

To observer

2H sin θ• thermal pressure is neglected

To observer

2H

L(Kudoh&Basu2014)y⊗θ > θcrit
2H

L cos θL csc θ

⊗

y

x

(c)

9

Figure 2. Di↵erent orientations of the magnetic ribbon with

Sin
revea
and c
make
we co
of 10
unifo
Furth
the r
a pai
Fig.
using
✓crit ,
locus
value
value
of ta
of co
repla
mean
to ⌃.
show
black
order
the a
dotte
wher

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”
11

Fig. 1. Radial distribution of the stellar density of a cluster.