星形成ゼミ 2015/ 2/15 Newsletter #255の26–30本目!

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26. “The dynamical evolution of molecular clouds near the Galactic Centre – I. Orbital
structure and evolutionary timeline” !
J. M. Diederik Kruijssen, James E. Dale and Steven N. Longmore!
http://adsabs.harvard.edu/abs/2015MNRAS.447.1059K!
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2015, MNRAS, 447, 1059-107!

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27. “Cross-sections for planetary systems interacting with passing stars and binaries”
Gongjie Li and Fred C. Adams!
http://mnras.oxfordjournals.org/cgi/doi/10.1093/mnras/stv012!
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2015, MNRAS 448, 344-363!

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28. “3D kinematics of the near-IR HH223 outflow in L723 R.”!
R. Lopez, J. A. Acosta-Pulido, R. Estalella, G. Gomez, and B. Garcia-Lorenzo (Spain)!
http://mnras.oxfordjournals.org/cgi/doi/10.1093/mnras/stu2607!
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2015, MNRAS 447, 2588-2602!

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29. “Study of morphology and stellar content of the Galactic H II region IRAS 16148−5011!
”!
K. K. Mallick, D. K. Ojha, M. Tamura, H. Linz, M. R. Samal and S. K. Ghosh!
http://mnras.oxfordjournals.org/cgi/doi/10.1093/mnras/stu2584!

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30. “SHADOWS CAST BY AWARP IN THE HD 142527 PROTOPLANETARY DISK” !
S. Marino, S. Perez, and S. Casassus!
http://stacks.iop.org/2041-8205/798/i=2/a=L44?
key=crossref.ce45c878a03a5d94cabe21b0186f7bac!
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2015, ApJL 798, L44!

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rbital structure and evolutionary timeline

M. Diederik Kruijssen,1‹ James E. Dale2,3 and Steven N. Longmore4-Planck Institut für Astrophysik, Karl-Schwarzschild-Straße 1, D-85748 Garching, Germany
ellence Cluster ‘Universe’, Boltzmannstraße 2, D-85748 Garching, Germany
26. “The dynamical evolution of molecular clouds near the Galactic Centre – I. Orbital
versitäts-Sternwarte München, Scheinerstraße 1, D-81679 München, Germany
andJohnevolutionary
ophysics Research structure
Institute, Liverpool
Moores University,timeline”
IC2, Liverpool!Science Park, 146 Brownlow Hill, Liverpool L3 5RF, UK

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ABSTRACTWe recently proposed that the star-forming potential of dense molecular clouds in the Central
Molecular Zone (CMZ, i.e. the central few 100 pc) of the Milky Way is intimately linked totheir orbital dynamics, potentially giving rise to an absolute-time sequence of star-forming
clouds. In this paper, we present an orbital model for the gas stream(s) observed in the
CMZ. The model is obtained by integrating orbits in the empirically constrained gravitational
2= 2.0) to the observed position–velocity distribution
potential and represents a good fit (χred3−3of dense (n > several10 cm ) gas, reproducing all of its key properties. The orbit is also
consistent with observational constraints not included in the fitting process, such as the 3D
space velocities of Sgr B2 and the Arches and Quintuplet clusters. It differs from previous,
parametric models in several respects: (1) the orbit is open rather than closed due to theextended mass distribution in the CMZ, (2) its orbital velocity (100–200 km s−1 ) is twice
as high as in previous models, and (3) Sgr A∗ coincides with the focus of the (eccentric)
orbit rather than being offset. Our orbital solution supports the recently proposed scenario in
which the dust ridge between G0.253+0.016 (‘the Brick’) and Sgr B2 represents an absolutetime sequence of star-forming clouds, of which the condensation was triggered by the tidal
compression during their most recent pericentre passage. We position the clouds on a common
timeline and find that their pericentre passages occurred 0.30–0.74 Myr ago. Given their shortfree-fall times (tff ∼ 0.34 Myr), the quiescent cloud G0.253+0.016 and the vigorously starforming complex Sgr B2 are separated by a single free-fall time of evolution, implying thatstar formation proceeds rapidly once collapse has been initiated. We provide the complete
orbital solution, as well as several quantitative predictions of our model (e.g. proper motions
and the positions of star formation ‘hotspots’). The paper is concluded with a discussion of theassumptions and possible caveats, as well as the position of the model in the Galactic context,
highlighting its relation to large-scale gas accretion, the dynamics of the bar, the x2 orbitalfamily, and the origin of the Arches and Quintuplet clusters.
Key words: stars: formation – ISM: clouds – ISM: kinematics and dynamics – Galaxy:
centre – galaxies: ISM.

tance, a fundamental physical understanding of star formation has
N T RO D U C T I銀河系の中心周りにある分子雲ガス、CentralON
not been achieved (McKee & Ostriker 2007; Kennicutt & Evansformation is one of the fundamental physical processes driving
2012; Krumholz 2014).
aryonic evolutionMolecularof the cosmos,Zone(CMZ)の運動の軌道を観測値からfrom reionizing the Universe
Several complicating factors are to blame for our limited unry high redshift to regulating galaxy evolution and depositing
derstanding of star formation. For instance, star formation is inmodelingした。!
ls in the interstellar medium (ISM), enabling the development
herently a multiscale process, of which the physics connecting theanetary systems and eventually life. Despite its critical impordifferent scales are highly complex. A wide range of recent work
has attempted to address this problem by connecting the empiriCMZでの星形成は朝夕力による影響を受けるた
cism of galactic star formation relations to the cloud-scale physics
of star formation (e.g. Bigiel et al. 2008; Heiderman et al. 2010;
ail: kruijssen@mpa-garching.mpg.de
め軌道決定することが重要らしい。今回、決定!

14 The Authors した軌道パラメータは先行研究の値と異なるとこshed by Oxford University Press on behalf of the Royal Astronomical Society

ろがある；!

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J. M. Diederik Kruijssen, James E. Dale and Steven N. Longmore!
http://adsabs.harvard.edu/abs/2015MNRAS.447.1059K!
pted 2014 November 27. Received 2014 November 6; in original form 2014 September 30!
2015, MNRAS, 447, 1059-107!

1. CMZで質量分布が広がっているため軌道は開いている(the orbit is open)!

2. 軌道速度(orbital velocity, 100-200 km/s) は以前のモデルに対して２倍大きい!
3. Sgr A* は(この研究で求めた)軌道の焦点にある!

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・fittingに利用していないSgr B2, Arches and Quintuplet の軌道も一貫している。!
・The Brick, G0.253+0.016の形成が、tidal compressionで誘発されたという我々の提案したシナリオをサポートする!

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この論文はシリーズ物でこれはIで、IIではシミュレーショ
ンについて出す予定。!

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J. M. D. Kruijssen, J. E. Dale and S. N. Longmore

27. “Cross-sections for planetary systems interacting with passing stars and binaries”
Gongjie Li and Fred C. Adams!
http://mnras.oxfordjournals.org/cgi/doi/10.1093/mnras/stv012!
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2015, MNRAS 448, 344-363!

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太陽系のような惑星系がクラスター内の他の恒星と遭遇したときの惑星軌道の変化をシ

ミュレーションで調べた。数値計算の結果から、cross-sectionを中心星質量、クラスター
速度分散、初期の惑星軌道半径、終状態での離心率の関数として求めた。!

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Motivation; 惑星系の多くはクラスター内で形成される。近年の研究で、太陽系の惑星はよりコンパクトな配列(軌道)で形成され外側に軌道が移動したと可能性が示唆されている。この研究では近接遭遇のcoros-sectionに着目して調べる。!

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method; 数値計算。（試行回数200万）。

この研究では、太陽系likeな系；木星、土星、天王星、海王星の４つの巨大惑星をもつ惑星系をターゲットとして基本的には考える(現在の太陽系、compact、離心率が異なるものとかを初期の惑星系としておく,Table 1)。
これにbinary(もしくはsingle star)が通過する。!

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↓Cross-section vs inclination (summation) . inclinationは初期ではどれ↓Cross-section vs 離心率も同じ公転面からスタートして、図は４線は木星、土星、天王星、海王星
つの惑星のinclinationを足したもの。実Blueは今の太陽系モデル、点線はこのFigure1ではCross-sections for planetary scattering 347
線、点線は衝突速度が1km/s - 16km/s
compact
配置。Figure2以降、他モデルを図示。
350
G.
Li and F. C. Adams
でstepはlog scaleでsqrt(2)倍。

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Cross-sections for planetary scattering

355

Results

G. Li and F. C. Adams
Figure 13. Scaled cross-sections versus eccentricity increase !e (equiva-

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Figure 14. Scaled cross-sections versus eccentricity increase !e (equiva-

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Analysis; cross-sectionの各変数への依存性をﬁtting

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lently, post-encounter eccentricity e) for a collection of different velocity
Even in the reduced case where we consider one planet at time,lently, the post-encounter eccentricity e) for the four giant planets in eachdispersions for the background. The starting state is taken to have four giof the starting architectures for the Solar system. The individual cases are
the interactions considered in this paper involve four bodies (theant planets with the current semimajor axes. Cross-sections are scaled by
shown
as
light
dottedcurves.
The
heavy
solid
blue
curve
depicts
theaverFigure 13.
host star, the planet, and two binary members). Unfortunately, four7/5
vb /a (see text). The individual cases are shown as light dotted curves,age, where the error bars depict the standard deviation. The straight red line
body interactions are rather difficult to describe analytically to any
with include curves for each of the planetsfor v b = 1–16 km s−1 , equallyshows the result for cross-sections with a purely exponential dependence on√reasonable degree of approximation. As a result, the goal of thisspaced logarithmically (by factors of 2). The heavy solid blue curve deeccentricity.
picts the average, where the error bars depict the standard deviation. The red
section is relatively modest: instead of building complicated anstraight line shows the result for cross-sections with a purely exponential
The curves in Fig. 13 are nearly straight lines on the semialytical models for four-body (and higher N-body) dynamics, we
dependence on eccentricity.
logarithmic plot, so that the dependence of the cross-sections on
consider here basic physical principles that can be used as motivaeccentricity
isnearlyexponential.
For
purposesofillustration,
wetion
combine these heuristic results with
354for scaling
G. Lilaws.
andWe
F. thenC. Adams
error bars depict one standard deviation. The heavy straight red line
use an exponential fitting function of the form
our detailed numerical determinations of the cross-sections. The
represents the same exponential dependence given in equation (7)⟨σ ⟩eresultis inphysicallymotivated
fittingcharacterize
Even
the reducedcase where
we formulaconsider thatone planet
at time,theand used in Fig. 13. The rms of the relative differences between= απ ℓ exp[ b (1 − e) ],
(7)
,
M
,
e).
cross-sections
overtheparameter
spaceofinterest(a,vab∗ (thethe interactions considered in this paper involve four bodies
the numerically determined cross-sections and the curve given by
Tostar,startthethe
discussion,
the simplest
case wherehostplanet,and two consider
binary members).
Unfortunately,
four-thewhere the first factor enforces consistency with the ejection crossequation (7) is ∼13 per cent. The cross-section curves are thus selfcross-section
for interactions
is the togeometricalcross-section
π a2
sections considered above. For the value b = 4/3, we obtain a good
similar to this level of accuracy. Furthermore, the dependence ofbody interactions
are rather difficult
describe analyticallyto anyfit to the calculated, scaled cross-sections, as shown by the heavy
the cross-sections on velocity dispersion is nearly independent ofprovideda planetin its initial orbit.
wegoalconsiderreasonablebydegreeof approximation.
As aFurther,result, theof thisthe
Figure 1. Cross-sections for eccentricity increase for the current Solar
Figure 2. Cross-sectionsforeccentricity
increaseforthecurrent
Solar
red line in Fig. 13. Except for first point (e = 0.1), the exponential
the dependence on starting semimajor axis a of the planet.
section tois berelatively
modest:insteadof duringbuilding
planets
independentof eachother
thecomplicatedencounters.anThis
system architecture and for a more compact configuration motivated by thesystem architecture and for a resonant configuration motivated by the Nicefit(straight
red
line)
agreeswiththeaveragevalues(solidblue
The dependence of the interaction cross-sections on the mass ofalyticalmodelsforfour-body(andhigherN-body)dynamics,
wecross-section
willbysystemgravitational
wecanNice model.
Forbetheenhanced
current Solar(solid bluefocusing,curves), theso
four
giant
model. For the current Solar system (solid blue curves), the four giant
curve)
towithin
about
3
per
cent,
i.e.
the
difference
is
much
less
the host star is somewhat more complicated than for the other variconsider
here
basic
physicalprinciples
that
can
be
used
as
motivawrite
down
anheuristic
expression
forthecross-section
in
the
form
planets are started with their current semimajor axes and zero eccentricity.
planets are started with their current semimajor axes and zero eccentricity.
thanthe widththedistributions
asare
measuredby theangles
standard
ables, as illustrated in Fig. 4. As the mass M∗ of the star increases,
Figure7. Cross-sections
forincreasing
the spreadof inclinationoftion forWe
theseheuristic
results
with where For
For scaling
thecompact
configuration(dashedred
curves),the planets
startedconfiguration(dashedofred
curves),the planets
started!laws.
"then combine
Figure6. Cross-sections
for eccentricity
increase
inare
systemsthethe resonant2
vesc
the cross-sections, considered as functions of eccentricity, become
deviations.
Another
measureof(forthequalityofthethesamefit isplane;
provided
the axes
planetary
orbits.
All10.38,of theandgiant
planets
arestartedinthewith giant
semimajor
axes
havinga fixedmratio
aj+1with semimajor
a = 5.88,
7.89,
12.01
au
the
analogues
where
Jupiter(jThe=cross-sections1)
j = 5/3,
planetsallhave(dashedred
curves).
The
our⟩0 detailed
determinationsofJ /athecross-sections.
P = 1m≈απ
a 2numerical
1+
,mass
(5)
⟨σ
2 location aJ = 5.2 au. For both sets of cross-sections,steeper. The spacing of the curves in Fig. 4 (for different stellar
therelativedifferences
betweenthenumericallydeterminedcrossof
Jupiterto
Neptune).
For
both
sets
of
cross-sections,thecurves,fromisstartedatitspresent
quantity
!i
is
the
total
range
of
inclination
anglesofthefourorbits
after
the
v
for
the
current
Solarsystemarchitecture
are
shown
for
comparison(solidresult is physically motivatedfitting formula that characterize the∞
top to bottom,
correspondto Jupiter(bottom),
Uranus,and the
Neptunethe curves,from top
to bottom,
correspondto Jupiter(bottom),
Saturn, semimajor
masses) grows with e, so that the curves are not self-similar. In
sections
used inconstructing
andform given
encounters.
Cross-sections
are Saturn,shownFig.
for 13a varietyofexponential
velocity dispersions,blue curves).
In bothcases, theplanets
arestartedM∗ ,currente). thecross-sections
theparameter
spaceof interest(a,withv b , theis overtheescapetarget
(at
where
v esc
−1
−1
(top).
Sincetheorbits
start
with
zero
eccentricity,theeccentricity
increase
Uranus,and
Neptune(top). speed
Since thefrom
orbits thestart withzerosystemeccentricity,
the
from
v
=
1
km
s
(top
curve)
to
v
=
16
km
s
(bottom
curve),
where
spite of this complication, we can still fit the cross-sections withbyequation
(7);
the
rms
of
these
relative
errors
is
∼12
per
cent.
Figure 12. Cross-sections
for planetary escapeversus the startingbbaxes
the giant planetsin ourSolar systemcase
and wherewith zeroeccentricity.
√ semiToeccentricity
start
theofincrease
discussion,considerthebe!e = e, whereis the Next
post
encountereccentricity.
!e
= e,
where
e isthethe simplest
post encountereccentricity.
isfortheasymptotic
relative
speed
locationoftheplanet),andv∞
theeThevaluesareevenlyspacedlogarithmically(by fourfactorsof dispersion
2).
a power-law function of stellar mass, although the accuracy of theweonconsidertheeffects ofvelocityof the
major
axis.
12
pointstheplotcorrespond
to
the
giant
planets
2
Cross-sections
are
shown
analoguesof
Jupiter(bottom
curves),Saturn,cross-sectionforinteractions
isthegeometricalcross-sectionπa
14. Scaled
cross-sections
versuseccentricity
increase!e (equivatween thetwosystems.
In ordertoversus
pass within
this cross-sectional
approximation is not expected to be as high as in the previous cases.
clusterenvironment.
Asarchitecture.
shownin Fig.
3 inthe previous
Figure13.
Scaledeccentricity
increase !e (equivain each of Figure
threebackgroundversionsof theinitialSolar system
The
symUranus,
andcross-sections
Neptune(top curves).
providedby a planet
in itsinitialorbit.beFurther,
we consider
the
area,
the
interacting
star(binary)
must
about
the
same
distance
lently,post-encountereccentricitye)foracollectionofdifferent
velocity
We thus consider a scaling of the form ⟨σ ⟩ ∝ M∗−µ , and vary the insection,
thecross-sections
vary successivewith
thepost-encountereccentricity
bolsofrepresent
the semimajoraxes of thesection).
The results areshownotheras theblueencounters.
curvesin Fig.
orbit
JupiteraJ = 5.2starting
andstates,
thenonincluding
letsufficiently
eachplanetlently,thepost-encountere)solid
for thefour
giant
planets
Solaratdifferent
systemisau,
unstablelongtime-scales(Batyginplanets
to be
independent
ofeccentricity
eachduringthe
This1, in eachfromtheplanet
as
its
host
star,
so
that
the
planethasachancedexµ to find the best fit. The choice µ = 1/3 provides the lowest rms
withapproximatelythesamefunctional
form
over
a
wide
range
of
present-day
Solarsystem(stars),acompact
configuration
with
5/3
semimadispersionsforthebackground.
The
starting
state
is
taken
to
have
four
giwhichalso
presents
thecross-sections
for
a
more
compact
starting
haveasemimajor
axis
that
is
largerthan
theprevious
one
by
a
of
the startingwill
architectures
forbytheSolar system.
The individual& Laughlin 2008; Laskar & Gastineau 2009).
cross-section
be enhanced
gravitational
focusing,
so we can cases are
jor axis
ratios(open
triangles),andan ultracompact
Solarsystemin uniform
the theleadingcoefficientchanges.
Moreover,thevevenly-spacedof being
ejectedfromits originalSolar
system.
This
expression
configuration(described
below).
The error
bars(notshown)
due
factorof 5/3.
This
Solar
systemthus
extendsout
tostartingb . Onlyantplanetswithcurrent
semimajor
axes.
Cross-sections
are scaled ofbythe relative error. Fig. 15 shows the result by plotting the scaled
shown
asconsiders
lightdottedcurves.
Thefor
heavy
solid
blue
curve
Another
way
in
which
planetary
orbitscanbealteredbyscattering
post-encountereccentricities
sufficiently
larger thethanaverthe
write down
an heuristic
expression
the cross-sectionin thedepicts
form
7/5 results,
multiple(open
squares).
The
red
shows
the
cross-section
spacingof theincurvesFig.line
3casesindicates
that
theascross-sections
cross-sections ⟨σ ⟩M∗1/3 /a (again including the linear dependence
to incompleteCarlosampling correspondto relativeonly
24 au.MMR
shown
Fig.individual
1,insolid
indicatethatthecrossthus where
represents
theMonte
escapecross-section.
The
parameter
α straightiserrorsa di- red linevThe(seeistext).
The
are
shown
age,
the!errorbars
the standard
deviation.
The
encountersbychanging
their
inclination
angles.
For light
all ofdotted
the curves,"depictb the/avalues).
2
indicated
limiting
formof equation(6); the blue
dashedcurve showson semimajor axis a) as a function of post-encounter eccentricity.
must
have
a power-law
on
the thanvelocity
dispersion −1
vb
withstartinga root-mean-square
(rms)
value
of is
∼4.4
per cent.
sections
for bythecompact
configuration
are dependence
somewhatsmallervesc
mensionlessof orderunity
andincluded
to encapsulate
2constantwithincludecurves
for
each
of
the
planets
for
v
=
1–16
km
s
,
equallyshowstheresultforcross-sections
withapurelyexponential
dependenceonbsimulations,
we
start
the
four
giant
planets
in
the
same
plane
(so
ofthe
planetary
mass onscattering
απFig.
aNext+ consider, the effects(5)⟨σ ⟩0 ≈ In
√ compact
the full
form. with(to theleadingorder).
We havescalings with velocity deThe light dotted lines show the individual (scaled) cross-sections
1,1 weandpaper,
cross-sections
fortheinthoseobtainedcurrent
semimajor
axes.explored
For this2
the uncertainties
inherent
in thisthis
approximation.
After insertingthevthroughout−γof encounters,∞ The
spacedlogarithmically
solid
blue
=
iS = (notiof
iNform=(by0).
thebestinclinationthateccentricity.
cross-sections.
shown inallFig.
6 forthethatusual SolarU =and 2).
findThethatheavythefit occurs
pendencethe⟨σfactors⟩After∝vb theand the heavy blue curve shows the average. The error bars depict
creasing
the eccentricity
toresults
e = 1 areincorporateof theways
Solar system,
theiJrms
errorsshown)
duetoincomplete
sampling
Cross-sectionsforplanetary
scattering
355curve deexpressionforescapespeed,
we
obtainpicts
theaverage,
where
the
error
bars
depictthestandard
deviation.
The
red
anglesofthefourplanetsare,ingeneral,non-zero.
Wedefinetheis thetheescape
speed
from
the
targetsystem(at thewherethevplanet
systemparameters
andfor
anSolaranalogueSolarsystemwhere allFigureoftoisolatethedependenceofthecross-sectionson
initialsemimaesc
canberemovedfromits
system.
These
channels
are
∼4.6
per
cent.
for γ ≈for
7/5.
Using escapethis choice
power-lawindex, we plot thethe corresponding standard deviation, which is larger than for the
12. Cross-sections
planetary
versusofthe
starting semi!
"
straight
line
shows
the
resultforcross-sections
withapurelyexponential
post-encounterspread!ioftheinclination
anglesaccording
to
the
is
the
asymptotic
relative
speed
belocationofplanet),
and
v
GM
=
1m
).
Both
classesthethegiant
planets
have
the
mass
of
Jupiter(m
include(1)
actually
increasingtheeccentricity
to
e
≥
1,
which
inNextweconsider
an
even
more
compact
orbitalconfiguration,The
curves
in
Fig.
13
are
nearlystraight
lines
on
the
semi∞
jor
axis
from
planet–planet
scattering
effects,
thislattercaseusesscaledcross-sections
versuspost-encountereccentricity
in
Fig.
14,
cases considered previously (compare Fig. 15 with Figs 13 and 14).
∗PJ
major axis. The 12 points on the plot correspond to the four giant planets
απtwo
a 2 systems.
1hyperbolic+→ toαπpass
ℓa,
(6)
⟨σ
⟩0 =
2 with
tween
the
In
order
withinthiscross-sectional
expression
dependenceoneccentricity.
cludesorbitsandplanetary
orbitsthatintersecttheagainmotivatedbythe ofNicemodel,
where
the a-dependence
four
giant
planets
where
weincludethelinearfound
previously(i.e.
The heavy red straight line shows the same result as before (fromof bothsystems
start
the
same
semimajor
axes(thepresent-day
smallerplanetmasses(byafactor
of
10),
so
they
act
more
like
test
in
each
of
three
versions
theinitialSolarsystemarchitecture.
The
symlogarithmic
plot,
so
the
dependenceofthecross-sections
on
avthatb
7/5"#$!
star, (2)inejection
from theSolarsystemthe ki-shows bols
aremasses;
in mutual
MMR.
Instarting
thiswidercase,
wechoosestarting
area, hostthe values
interactingstarsystem)(binary)must
beaboutbytheincreasing
same
distanceEachlightdottedcurve
theaxesfigureshows the result⟨σ ⟩v
equation 7). The rms error between the exponential line and thestates,includingsemimajor
ofcompact
the
our
and
zero
eccentricity.
The
figurethat
todifferent
explorerange
a,thethewealso insemimajor
take
this
b a /a).
eccentricity
isreplacednearlyexponential.
For
purposesofillustration,
we represent
JauJin
j · (Jupiter),kthevelocity
where
wethe
have
theforasymptotic
v ∞(3)withenergyso itsthathost
theorbit
issounbound,capturebyaxes
to havevalues
of
acos
=−1
7.89
au (Saturn),from netictheplanetasstar,
thatspeedtheplanethasa chanceforone
planetandone
choiceof velocity
dispersion.
heavy
numerically determined cross-sections is about 20 per cent. This
Solar
(stars),a factorcompactwith
5/3, compared(4) red line!i
≡systemmaxcross-sections
increasing
theandof oneNeptune present-day
andsystemtobesmallerby
a5.88
ofconfiguration
1.35
to semimathat considerror
bars
depictone
deviation.
The
heavyThe
straight
useanfittingfunctionofsystem.
theareeccentricities
form
Jultracompact
Jstandardthecluster(orotherbackgroundstellardispersionv(two)
jthe
k averageof exponential
thestars.
These
channelsnotmutually
exclu10.38
au(Uranus),
andcurve
12.01andau an(Neptune).
WithSolar
thesesemimajor
b of passing
jor
axis
ratios(open
triangles),systemstarting
in
blue
shows
over
all
of
the
curves,wherethelarger error measure results from fitting the cross-sections with aof being
ejectedfromlargely
itsoriginal
Solar
This
expression
Uranusareunaffectedby theincreaseinsystem),planetarymass,ered in theprevious section.
The exponential
error bars delineate
the uncertainty
represents
thesamedependence
givenin
equation
(7)
2
sive,
but thesimulations
are stopped after
one of theseℓeventsaxes, Jupiter
Saturnare inThe
a 3:2
SaturnandtheUranusareandhave
definedthe corresponding
length-scale≡GMmultiple
MMRand(open
redMMR,
solid lineshowscross-section∗ /vb larger.
thus
represents
theescapecross-section.
The
α isatakesdiangular
momentum
vectorsof crossthe planetary
where
thesquares).
JMontebut thecross-sections
for Jupiterandparameter
Saturn are
somewhat
⟨σ
⟩ewe
due
toMMR,
incompleteCarlosampling.
Not
onlydo(forthei are the−1include ejection processes that
and
used
in
Fig.
13.
The
rms
of
the
relativedifferences
betweenplace.
However,thesechannels
only
in
a
3:2
while
Uranusand
Neptune
are
in
a
5:4
MMR
indicatedbythelimiting
formofequation
(6);
the
blue
dashedcurve
shows(where
ℓ
∼
890
au
for
v
=
1
km
s
).
The
final
expression
rep= Note
απ constantℓ thatexp[thebof(1−], and(7)b e)unity
mensionlessorder
isplotted
included
to forencapsulate
orbitswhere
the indices
run throughall
fourof&the ofgiant
cross-sectionsareonly
eccentricity values
sections
show andnearlydependence
on a,the slope
theplanets.
MNRAS 448, 344–363 (2015)
a happenthefull form.
duringor form,
immediately
after
the encounter(weinserting
denote thesefurther
discussionof this linear
initial determinedstate,and others,seebutBatygin
theresulting
numericallycross-sections
andofthecurve given by
resentsthewhich
is applicablewhengravitational
the uncertainties
inherent
insuchthisaapproximation.
After
the
cross-sections
for Liincreasing
the2014).
spreadinclinatione ≥limiting
0.20.
Withinmassivesystem,small
curve2010;is Thepredicted
by&theabove analysis.
The
solid
(blue
dashed)processes
as prompt
ejections).
Inlinearother planetarycases,
the planetsarescat-eccentricBrown
Nesvorný
Morbidelli
2012;
& red
Batyginfocusing
dominates,
andimpliesadependenceofthecrosswhere
the
first
factorenforcesconsistency
withtheejectioncrossequation
(7)
is
∼13
per
cent.
The
cross-section
curvesareathus selfexpressionfortheescapespeed,
we
obtainanglesis shows
shown
in
Fig.
7.not
Theimply
figureshowsthesemimacross-sectionsfor(e∼ 0.10) are easily
excitedbyorbits
planet–planet
interactions;
as
curve
in
Fig.
12
the
cross-section
predictedbyequation(6)
toisolatethedependenceofthecross-sectionson
initialteredities
into !
high-eccentricity
orbits,
so
that
the
cross
each
other.
Note
that
the
semimajor
axis
ratiosdoperiodratioswith"section on
a. Given
thisform,For
theforcross-section
another
sections
above.
valueb = requires
4/3,
weplanets
a goodsimilarto thislevel
offor
accuracy.
Furthermore,
the dependenceofrange
of casevelocity
dispersions
ofboththeAllbackgroundcluster,
result,theGMcross-sections
eccentricity
increase–obtainas determinedjor
axis
planet–planet
effects,
thiscaseuses from
Withaconsidered
these
configurations,in the theabsenceof resonance,
the
exact
integervalues(althoughthey areclose).
oflattertheforthefrom
limiting(fullscatteringform);
cases,
theorbital
characteristic∗√ vb = 1απa 2 In1 this
+
→ απ
ℓa,
⟨σ ⟩tolength-scale.
problem,theorbit –speedofbinary,
the
0 =
2
fit
the
calculated,scaledcross-sections,as theshown
by(6)andthe areheavy
onvelocity
dispersion
isα testnearly
of
,au
where
the
valuesspacedfactorsof 1 independent
2. The
tothe
16
km
s−1
throughournumerical
schemeare whichextremely
notwill
eventuallyexperienceclose
encounters,in turnlarge
leadtoelements
must
beproperlyto
put sothearesystemin bymutualsmallerplanetmasses(byaandfactor
of
10),
they
act
more
like
length-scaleℓcross-sections=chosen890
thedimensionlessparameter
=
7/5.
av
b
of(we
thedenote
encounter,
the asspeed
theasymptotic
speed
v ∞ figure.
ejectionsor collisions
this point
processdelayedejection).
MMR,
and
this
requirement
displacestheperiodratiosarealmostinsomewhat.
the axissemi-logarithmicplotplottedin13.
themasses;to cross-sections
explorea widerrange
in evenly
a, we
also
take
this
compact
red
line
inhaveFig.
Exceptforfirst(eorbit=
0.1),
theofexponential
theconsider
dependence
on starting
semimajor
a of
Next
we
the
dependence
ofspacedthecross-sections
on thethe planet.
withthevelocitywhere
wereplacedtheasymptotic
speed
v
∞
arecross-sections
all beroughly
comparable
planet,the velocity
Theandcross-sections
fordispersion
delayedconsideredNonethetoless,theresonancethesystemall tothreeandsmallerhave
nearlythe sameshapeasequivalenta(forfunctionof!i.
properties
The
numerical
resultsejections
forv bthe
canlater.
be(solid
understoodsystembeby a angles
factorof
1.35
comparedthat
considpost-encountereccentricity
eof(whichistoplanet%eThesebecausefit(straight
red
line)
agreeswiththewillaveragevaluesblue
Thedependenceoftheinteraction
cross-sections
on the mass of−1
thecluster(orotherstellardispersion
vcomparison,
b).of
For
example,theorder,orbitbackgroundoftheNeptune
in ourSolarsys(1–10km
For
we alsopresent
resultsfroma system),seriesof to acteredpairs)
insection.
the theinitialstate,
requiredWith
indicatethatcross-section
haspower-lawdependence
on votherass within
follows:
to leading
we
often
the
planetsasthelibrating
previous
errorasbars
delineate
theuncertaintybtheinareorbits
start
with
zero
eccentricity).
ForaforallMMR.
four
planetsin all2
curve)
toabout
3using
per
cent,
i.e.
the expect
difference
less
the
host
star
isThe
somewhat
morecomplicatedthanfor thevari−1/vmuchand haswehavedefinedthe corresponding
length-scaleℓ ≡architectureGM∗isnumerical
experimentsamorecompact
orbitalthisinitialstate,
theSolarsystemismuch
more
compact
than
at
the
,
whereas
the
orbit
speed
tem
an
escape
speed
of
∼5.5
km
s
b
to incompleteMontesampling.
Not only
dothe cross(seeofSection
4).Carlo thetest particles, so that the cross-sections
should not be sensitivedue
to
−1
three
types
Solar
system,e-dependence
issimilar.
Since
the−1standard(where∼890
forv bthe= curves
1 km sindistribution).
The
expression
repof
the
star
increases,
thantheℓthe
width
of
the
distributionsasmeasuredbytheables,
asillustrated
in
Fig.
4.
As
the
mass
M
(showntheau
reddashedFig.
1),final
which
is motivatedbypresent
epoch,
andthecross-sectionsforinteractionsaresmaller.
∗.
ofa binary
atas
the
peak
of
period
is
also
∼5
km
s
show
lineardependencea, buttheangles
slope axisof theInnearly
general,increasesin theon
inclinationarepositively
corplanetary masses. For sufficiently massive planets, however,sections
anejectioncross-sections
scale
linearly
withsemimajor
a(seeresents
thelimiting
form,
which
is
applicablewhengravitational
Nice
model
of Solarsystemformation(Gomesetal.fit2005).
This trend isthe
illustrated
in Fig. 2, which
compares the
cross-sections
cross-sections,consideredasfunctions
eccentricity,
deviations.
Anotherofof
thequality
ofleadtheisimprovided
If wethe
only
have
ain
singlevelocityV,
thenanalysis
curveis12),
predicted
bytheincreases
above
analysis.
The reddividing
solid
(blue
relatedwithineccentricity.
This
resultis ofnotpower
unexpected, become
increase
the measureeccentricity
onedimensional
planet can
to significant
perFig.those
we scalecross-sections
outdashed)one
focusing
dominates,
andimpliesdependence
theAlthough
the Nicemodel
has aa linear
numberone
featurewith
obtained
forthesolar
systems with by
the standard starting
2 cross, as
given
plies
thatturbations
thedifferences
relevant
length-scalemust
beofℓvariations,=therebyGM∗of/V
spacingof
thecurvesininFig.
4(6)(for different stellar
therelative
betweentheplanets,
numericallydeterminedcrosscurveFig.
12showsThethecross-section
predicted
by equation
assteeper.
changesinbothorbitalelements
to disruption
oncould
the otherleading
to increasedof a.inTheresultingscaledcross-sections
arecorrespond
shown
Fig.
13 as a of theis thattheGiven
giantacting
planetshaveformed witha morecompact
configuration.
To leading
order,
thesmallercross-sections
obtainedsection
on
a.
this
form,
the
cross-section
requiresanother
in equation(6);inadditional
uncertainties
can thebe absorbed
into
the
for
the limiting
case
(full
forthisthecharacteristic
masses)grows
e,bothsothatindividual
curvesnot self-similar.
In
sections
used
constructing
Fig.
13day.
andexponential
form
given
initialstates.
Toform);illustrate
trend,
in theFig.
8 weplotaretheincreasesin
eccentricity
By
increasing
massof
the planetsofeccentricitye.
Inawithaddition
tocases,thecases(shownconfiguration
thanexcitation.
that of thethe
present
Forthe
this
case,
weall
fixof
the
forfunction
the resonantarchitecturearedirectconsequenceof thesmallerlength-scale.
In thisproblem,
orbit
speed
of
the
binary,the1dimensionlessparameterα.
Finally,
note
that
for
ais
velocity
dis-cent.length-scale
ℓ=
890
and
the dimensionlessparameter
α
=
7/5.
the
spreadofinclination
angle !iversuschangein cross-sections
eccentricity
to thatof Jupiter,theresultingsolar
systems
are
more
The
spite
of au
this
complication,
can thestill
fitheavy
the
with
by
equation
(7);
the
rms
of
thesewerelative
errors∼12excitable.
peras the light
dottedcurves),the average
iswe
shown
as
the
blue
asymptotic
speedv∞−1 of the encounter, the orbit speed of the
=
1
km
s
,
the
gravitational
focusing
term
dominates
persion vlargest
Next
weconsiderthedependenceofthecross-sectionsontheb(equivalently,
thepost-encountereccentricity
since
= e).
The of theincreasethein thecross-sections,which occurs
for Jupitera power-law
functionof stellar
mass,althoughtheaccuracy
Nextandweconsidereffects
velocity
dispersionof andthecurve, where
the error bars
depict thestandard
deviation.
This!elatteraretheall roughly
comparable
planet,thevelocity
dispersionv b of
MNRAS
448, 344–363
post-encountereccentricity(whichis equivalent
tobut%e(2015)
becauseby a factorof 30.eccentricities, is only
variables
areeinfact
wellcorrelated,
theoffactor of∼2;
mostcross-sections
quantitytwo
providesa measurethe
spreadintothebevaluesof rangethecross).cluster
For example,
the orbit As
of aNeptune
inour(1–10 kmfor
s−1low
approximation
is of
not
expected
as highasallin
thepossible
previous cases.
backgroundenvironment.
shown
instarting
Fig.Solar
3 semimainsysthe previous
theorbits
start
with
zero
eccentricity).
For
all
four
planetsin
The limitingformofequation(6)
islinearinthe
Figure14.
Scaledcross-sections
versuseccentricity
increase!e(equiva−1!ivaluesforagiven
eccentricity
e
=
!e
is
large.
Asaresult,intheexperiencesmallerchanges.
These
resultsaregenerally
consistent
13.
Scaledcross-sections
versus eccentricity!e (equiva−µwhereasthe orbit increase
speed
tem has an Figure
escapespeedof ∼5.5kms ,thesection overthevariouscases.
The standard
deviation
fromandvary the inWe
thus
considerae-dependence
scalingofe)theform⟨σ⟩varies∝theM
section,cross-sections
witheccentricity
threetypesfigureofSolarsystem,thesimilar.
Since∗ in ,!elently,eccentricity
a collection
different
velocity
jor axis athe
oflently,
theTo ourseevary
howeccentricity
wellthisispost-encounter
expression
works,
we planthee)‘full’,
the
four
giant
in
each
wepost-encounterplot
thecross-section
meanvalues atof low!iisforaveraged
overeof
a=
bin
with
with
theplanet.
ideapost-encounter
that
Solar
systemi.e.
noadditional.
of a binary
at thepeakof theperioddistribution
isforalso∼5km
s−1planets
about
17
per
cent
of
the
eccentricity
0.10
dex
µ
to
find
the
best
fit.
The
choiceµ=1/3
providesthelowestrms
withapproximatelythesamefunctional
form
over
a
wide
range
of
dispersions
for0.05.
thelinearlybackground.
The
starting
state
to have four
giejection
scalewithshowsemimajor
axisisataken
(see
plot
ejectionofV,theplanets
versussemimajor
axis planets
ofand
theacross-sections
startingarchitectures
for
the
Solar
system.
cases
are cross-sections
a
width
of
δ
=
The
data
a
well-definedcorrelation;foretshavelittle
mass
can
be
added
toThe
theindividual
extantIf wetheonly
singleadditional
velocitythendimensional
analysis
imto 12),
only we
about
9 per
e=
1.0. byplanets
withatthecurrent
semimajor
axes.
are scaledthe
Fig.
scale
thecent
dividing
out
one
power
shown
as light thedottedThe
heavy solid
depicts
theaverOnly
theleadingcoefficientchanges.
Moreover,
the
ofant
the
relative
error.
Fig.
15
shows
theCross-sections
resultby plotting
vin
12.
As
expected,ejectioncross-section
is/V
ablue
nearlylinear
2 curveb .Fig.
thischoiceofcross-sections
binning,thespreadin theinclination
anglesgrows toby scaled
without
rendering
thecurves.
systemunstable.
In
fact,
theuniform
current, aseven
given
plies
that
the
relevant
length-scalemustbeℓ = GM
7/5∗
1/3 are cases
vb /a scaled(see text).
The
as as
light
of a.
cross-sections
shownareinshown
Fig. 13a dotted curves,the error
bars
depictthe standard
deviation.
The straight red
lineThe resulting
functionofthesemimajor
axis.
This
for
solar
systems
dependence
spacing
ofage,
thewherecurvesinFig.
3 trendindicatesthat
theintocross-sections
cross-sections⟨σ⟩Mindividualin equation(6);
additional
uncertainties
canholds
be absorbed
the∗ /a (again including the linearwith includee.curves
for eachof theplanetsforv b (shown= 1–16 km s−1 , equally
the result for semimajorcross-sections
with(stara purely
exponential
dependence
on of eccentricity
function
In addition
to theindividual
cases√ parameter,
starting
withshowsthepresent-day
axes
symbols)andthe
1 In orderonsemimajor
axis
a)
as
a
functionofpost-encountereccentricity.
must
have
a
power-lawdependenceonthevelocity
dispersionv
dimensionlessparameterα.
Finally,
we
note
that
for
a
velocity
disto
set
the
value
for
this
dimensionlessandothersbspacedlogarithmically(byfactors
of
2).
The
heavy
solid
eccentricity.
as the light dotted curves), the average is shown as the heavy−2
blue blue curve deMNRAS
448,(2015)−1 344–363
morecompact
configuration
where
the
semimajor
axes
arespacedspecified inThethis section,
we generally
searchin theincrementsof
10 (scaled), findthe cross-sections
1 km s ,We
the havegravitational
focusing
term dominates
persionv b = order).
(toleading
exploredscalings
with velocity curve,delight
dottedlines
showbars
individual
the average,
where
thestandard
error
depicttheThisstandard
deviation.
The redwhere picts
the error
bars
depictthedeviation.
latterby
5/3 (open triangles).−γWe also plot results for ultravalue that gives
the line
minimumrms
error,
then choose thewith
nearestroundexponential
by afactors
factor ofof30.
straightshows
the spread
result andforcross-sections
a purely
andnearly
find thatthe lines
best onfit theoccurspendenceof Thethe curves
form ⟨σ
∝ v13
and
the
heavy
blue
curve
shows
the average.
Theerror bars depictin ⟩Fig.
arestraight
semiquantity
measureofthebMMR
numberprovides(ratioof arelatively
small
integers).in the values of the crosscompact
solar systems
in multiple(opensquares).
In orderThe limitingformofequation(6)
islinearinthestarting
semimadependenceoneccentricity.
logarithmic
so that thethe cross-sectionsfor γ ≈ 7/5.
Using plot,
this choiceof dependence
power-law ofindex,
we plot section
theon overthe
which
the corresponding
various cases. Thestandard
standard deviation,
deviation varies
fromis larger than for thejor axis a ofeccentricity
the planet. To
see howwell this expression
works,
isnearlyexponential.
For eccentricity
purposesof weillustration,cent ofconsidered
the cross-section
at low (compare
eccentricityFig.
e = 150.10with Figs 13 and 14).
scaledversuspost-encounterin Fig. about
14,we17 percases
previously
plot thecross-sections
ejectioncross-sectionsofthe planetsversussemimajor
axis
error
bars
depictone
standard
deviation.
The
heavy
straight red line
useanexponential
fittingfunction
of
the
form
to only about
9 percent atrede =straight
1.0.
MNRAS
448,
344–363(2015)
where
we
includethethelineara-dependence
(i.e.
The
heavy
line shows
the sameas before
in Fig. 12.
As
expected,ejectioncross-section isfounda nearlypreviously
linearrepresents
the same
exponential
dependence
givenresultin equation(7) (from
7/5⟩esemimajor
Each
light
dottedcurve
inthefigureshows
theresult⟨σ⟩vb /a).
equation
7).
The
rms
error
betweentheexponential
line and thefunction
of ⟨σ
the
axis.
This
trend
holds
for
solar
systems
and used in Fig. 13. The rms of the relative differences between= απ ℓ exp[ b (1 − e) ],(7)
a present-day
startingthesemimajor(star symbols)
and the1 In order to numericallyfor
onewithplanetand one choiceof axesvelocity
dispersion.
The heavycross-sections
isothers
about
20givenper cent.
thethenumericallydeterminedcross-sections
and
the
curve
by Thissetvalue fordetermined
thisdimensionlessparameter, andmore curve
compact
configuration
where
the
semimajor
are spaced
section,wesearchin incrementsof
10−2 , find
the
where
the thefirst averagefactorenforcesconsistency
the ejection
crossequation(7)generallyis∼13 per
cent.
Thefromcross-section
curvesare thus self- with a
blue
shows
over
all of axes
the with
curves,where specified
the in thislargererror
measureresults
fitting
thecross-sections
by factors of
5/3 (open
triangles).
WeFor
alsotheplot
results
the minimum
error,
then choose
the nearest roundsections
considered
above.
valueb = for
4/3,ultrawe obtain avalue
goodthat givessimilar
to thisrms
level
ofand
accuracy.
Furthermore,
the dependence ofsmall integers).
compact solar
systemsin multiple
MMR(open squares).
orderby the number
fit tothe calculated,scaledcross-sections,
as In
shown
heavy (ratio of
therelatively
cross-sections
on velocity dispersion is nearly independent of
red line in Fig. 13. Except for first point (e = 0.1), the exponential
the dependence on starting semimajor axis
a of the 448,
planet.
MNRAS
344–363 (2015)
fit 344–363(straight (2015)
red line) agrees with the average values (solid blue
The dependence of the interaction cross-sections on the mass of
MNRAS 448,
curve) to within about 3 per cent, i.e. the difference is much lessthe host star is somewhat more complicated than for the other varithan the width of the distributions as measured by the standard
ables, as illustrated in Fig. 4. As the mass M∗ of the star increases,
the cross-sections, considered as functions of eccentricity, become
deviations. Another measure of the quality of the fit is provided
steeper. The spacing of the curves in Fig. 4 (for different stellar
the relative differences between the numerically determined crossmasses) grows with e, so that the curves are not self-similar. In
sections used in constructing Fig. 13 and the exponential form givenspite of this complication, we can still fit the cross-sections withby equation (7); the rms of these relative errors is ∼12 per cent.
a power-law function of stellar mass, although the accuracy of the
Next we consider the effects of the velocity dispersion of theapproximation is not expected to be as high as in the previous cases.
background cluster environment. As shown in Fig. 3 in the previous
We thus consider a scaling of the form ⟨σ ⟩ ∝ M∗−µ , and vary the insection, the cross-sections vary with the post-encounter eccentricity
dex µ to find the best fit. The choice µ = 1/3 provides the lowest rms
with approximately the same functional form over a wide range ofof the relative error. Fig. 15 shows the result by plotting the scaledv b . Only the leading coefficient changes. Moreover, the uniform
spacing of the curves in Fig. 3 indicates that the cross-sections
cross-sections ⟨σ ⟩M∗1/3 /a (again including the linear dependence
on semimajor axis a) as a function of post-encounter eccentricity.
must have a power-law dependence on the velocity dispersion v b(to leading order). We have explored scalings with velocity deThe light dotted lines show the individual (scaled) cross-sections−γ
pendence of the form ⟨σ ⟩ ∝ vb and find that the best fit occurs
and the heavy blue curve shows the average. The error bars depictfor γ ≈ 7/5. Using this choice of power-law index, we plot thethe corresponding standard deviation, which is larger than for the
scaled cross-sections versus post-encounter eccentricity in Fig. 14,
cases considered previously (compare Fig. 15 with Figs 13 and 14).
where we include the linear a-dependence found previously (i.e.
The heavy red straight line shows the same result as before (from7/5⟨σ ⟩vb /a). Each light dotted curve in the figure shows the result
equation 7). The rms error between the exponential line and thefor one planet and one choice of velocity dispersion. The heavy
numerically determined cross-sections is about 20 per cent. This
blue curve shows the average over all of the curves, where thelarger error measure results from fitting the cross-sections with a

Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of Japan on February 12, 2015354

28. “3D kinematics of the near-IR HH223 outflow in L723 R.”!
R. Lopez, J. A. Acosta-Pulido, R. Estalella, G. Gomez, and B. Garcia-Lorenzo (Spain)!
http://mnras.oxfordjournals.org/cgi/doi/10.1093/mnras/stu2607!
!
2015, MNRAS 447, 2588-2602!

!
!
!
!
!
!
!
!
!
!
!
!
!
!
!
!
!
!

近赤外にあるH2 line 2.122 um (他 1.7umや[FII]1.26 um)の輝線をWilliam Herschel
Telescope 4.2m, MOS and LIRIS用いて、L723にあるoutflowを観測した。!
H2 lineでは複数のHH天体 (HH233)が見られ、shockで励起されたガスのknotsがある。!
観測ではLine分光で視線速度を測定した。また複数期間での観測をし固有運動を調べた。!
CO outflowが知られており、H2 lineでの結果と比較。!
titleの3D kinematics は including proper motions and radial velocitiesでIR H2 line (2.122um)のHH233 knotsの駆動源を同定した。!
L723; !
distance 300±150 pc!
CO outflow (2 separate pairs of red-blue lobes)!
中心ではVLAの観測で VLA2が4つのYSOs(VLA 2A, 2B, 2C, 2D)があることが知られている。VLA 2Aが東西のCO outflowの駆動源である。!
SMAの観測1.35mmで二つの成分が知られている(SMA 1 and SMA 2)!
!

!
ページ 6/10

↓P-V diagram;→H2 2.122 umのL723領域のimage.
2590

LIRIS MOS of HH 223

長方形はMOSでの分光のスリット。円形(eg, J22)はmaskしたところ。

2593

R. López et al.

SMA 2はK1,K2付近にいる。

!

!

!
!
!
!
!
!
!
!
!
!
!
!
!

Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of Japan on February 12, 2015

Slitlet

Knot

Offset b(arcsec)

VLSR c(km s−1 )

s1
s2
s3
s4
s5
s6
s7
s8
s9
s11
s12
s13
s14
s15
s16

N1
A
B
C
D-E
F
N2
K2
K1
H2-NWabH2-NW3H2-NW2ab
H2-NW2d
H2-NW2eH2-NW2fg

−83.9
−67.0
−64.0
−59.9
−55.4
−51.4
−37.7
−8.3
−3.9
40.3
82.8
101.9
112.8
118.2
123.9

−72.1 ± 12.4
−26.3 ± 9.7−21.3 ± 8.2−7.0 ± 6.2−33.2 ± 12.4
−39.2 ± 14.0
−30.3 ± 6.5−24.0 ± 9.8+7.3 ± 8.0+1.5 ± 8.3+43.3 ± 11.5
+49.0 ± 10.6
+36.0 ± 6.1+39.9 ± 12.6
+25.6 ± 7.9

Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of Japan on February 12, 2015

Table 3. Radial velocities (VLSR )a of the H2 2.122 µm emission within the full slitlet aperture.

Notes. a From Gaussian fits to the H2 2.122 µm line of thespectra, averaged for the full slitlet apertures (see the text).
b Offsets from the position of SMA 2 to the centre of the
Figure 6. Position-velocity maps of the H 2.122 µm line through the slitlet labelled in each panel (black contours). The PV maps of the [Fe ] 1.644 µm aresuperposed (red contours) for the slitlets where [Fe ] emission was detected. The PV maps of the H 1.748 µm line are also displayed (green contours) for
slitlet.
Typical values of the error are ≤ 0.5 arcsec.
these slitlets to check the reliability of the spatial displacement between the H and [Fe ] emissions. The knots intersected by each slitlet have been labelledc The error values correspond to the error of the Gaussian fit
to help with the identification.
added in quadrature with the rms residual of the wavelength(dome flats, hereafter) spectra were used to determine2600
the correc-R. López et al.
calibration.

Figure 1. Image of the L723 field obtained with LIRIS through the H2 2.122 µm line filter (the continuum has been subtracted). The rectangles represent theposition of the slitlets of the mask designed for the multi-object spectroscopy (MOS) observations. The knots of HH 223 enclosed in the corresponding slitletof the mask have been labelled according to the nomenclature of López et al. (2010). An additional slitlet, labelled J23 was included to check the individual
frame quality. The four circular appertures, labelled J10, J14, J22 and FJ1, used for mask centring, are also drawn. North is up and east is to the left.

!

Table 2. 2596
Configuration
the mask.et al.
R.ofLópez

II

2

II

2

2

II

López et al. 2
each slitlet. T
each of these
slitlet encomp
1 arcsec wide
the velocity ofas the offsetof the knot, isthe knots proj
outflow are plwill be usednear-infraredcomplex kinem
which confirm
HH 223 optica
additional H2in the previous
Finally, andpattern foundof the knots fr
in most of th
and 10 (green
good agreeme
Concerning thwith those derbrightness disµm line are als
the peaks of th
general dtrend fVTthe(km1.748s−1 ) µm∼40–60 km s−
25.3
± SNR
2.0 lev
same

SlitletLengthCenter Position→CO
outﬂowαID (arcsec)

2000

δ 2000

PA
(deg)

!

・H2 lineによる大局的に視線速度は双曲流パタン(bipolar pattern)を示した。SMA2に対して、南東方向にnegative velocity(blue shifted)
で北西方向にpositive (redshifted)。この結果

は、CO lobeとよく一致し、H2とCO outflow
は同じ駆動源であることを強く支持する!

・固有運動は他の低質量原始星で知られる典型的なH2 outflowの値の範囲内であった。

blueshifted と redshiftedのH2 knotsの固有運動は同程度の値である。一般的な傾向として、↓固有運動

Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of Japan on February 12, 2015

MNRAS 447, 2588–2602 (2015)tion. Next, we used the a-priori mask design information to trace
Table 6. Positions and proper motions of H2 knots.
the slitlet positions and their limits from a distortion-free dome–flatframe. Then the 2D spectra from each slitlet were extracted followKnot
xa
ya
µx b
µy b
ϵxc
ϵyc
ing the limits determined previously and were calibrated indepenN1
s1
6.0
19 17 59.034 +19 11 42.10−95
HH 223(arcsec)(arcsec)(arcsec yr−1 )(arcsec yr−1 )(arcsec)(arcsec)
dently.
The
wavelength
calibration
wasdetermined
from
argon
and
A
s2
3.0
19 17 57.970 +19 11 50.02−70
xenon lamp exposures obtained with the same instrument configuA−60.6−29.8 −0.0177 ± 0.0014 −0.0015 ± 0.0015
0.00590.0057B
s3
2.5
19 17 57.717 +19 11 49.51−70
C-North−54.2−27.0 −0.0080 ± 0.0003 −0.0029 ± 0.0006
0.00120.0027 12.2 ± 0.5
1
ration through the MOS mask. The arc spectra for each slitlet wereC
s4
3.5
19 17 57.501 +19 11 52.71−70
C-South
−52.4−29.2−0.0126±0.00130.0005±0.00200.00530.009717.9 ± 1.8
extracted in the way described before. The initial wavelength caliD-E
s5
4.5
19 17 57.087 +19 11 51.33
−90
E
−47.3
−28.2
−0.0094±0.0018−0.0015±0.00050.00810.0022
13.5
±
2.5
F
s6
4.5
19 17 56.697 +19 11 50.04−70
bration was determined as in the ‘long-slit case’ following the usual
3.2.2 Kinema
F0−46.4−31.6 −0.0076 ± 0.0003 −0.0057 ± 0.0040
0.00140.0192 13.5 ± 3.5
12
N2
s7
13.0
19 17 55.660 +19 11 53.80−70
tasks in IRAF (i.e. combining identify, reidentify and fitcoords). The
F1−43.0−29.9 −0.0122 ± 0.0026 −0.0005 ± 0.0057
0.01140.0261 17.4
3.7
We ±derived
th9
K2
s8
2.0
19 17 54.100 +19 12 13.60−70
accuracy of the wavelength calibration was then checked using a set
K2−5.5−5.9 −0.0076 ± 0.0024 −0.0110 ± 0.0007
0.01000.0036 19.0
± 2.1
1
locity
from
th
K1
s9
5.0
19 17 53.950 +19 12 18.00−87
of bright, isolated OH sky lines present in the observed wavelength
K1-West−3.4−1.6 −0.0008 ± 0.0004 −0.0075 ± 0.0015
0.00200.0071 10.7 ± 2.1
11.257 µm andstar J23
s10
4.0
19 17 53.170 +19 12 33.20−75
H2-NWc31.7
16.5
0.0063 ± 0.0014 −0.0003 ± 0.0001
0.00620.00029.0 ± 1.9−
ranges. It was noted that wavelength offsets were present, which is
given in Table
H2-NWabs11
15.0
19 17 51.235 +19 12 40.67−40
H2-NWb
34.5
20.3
0.0131 ± 0.0008
0.0026 ± 0.0005
0.00320.0023 19.1 ± 1.1−
expecteddue
tospectrographflexures.
The
determinationofradial
H2-NW3s12
5.0
19 17 48.440 +19 12 57.50−70valuesappearH2-NWa36.3
22.5
0.0078 ± 0.0001 −0.0005 ± 0.0003
0.00010.0012 11.1
± 0.1−
velocities from emission lines is a crucial aspect of this work, thereH2-NW2abc
s13
16.7
19 17 47.560 +19 13 14.40
−70
These
velociti
H2-NW374.7
37.7
0.0098 ± 0.0031
0.0061 ± 0.0036
0.01400.0171 16.4
± 4.6
−5
fore we needed to refine the wavelength calibration by applying an
H2-NW2ds14
4.5
19 17 46.643 +19 13 13.81−70eby[Fe II]−6H2-NW2b1
86.7
52.9
0.0028 ± 0.0021
0.0011 ± 0.0016
0.00890.0076
4.3the± 2.8
offset determined from the wavelength difference observed in the
H2-NW2es15
4.0
19 17 46.101 +19 13 10.85−70
H2-NW2c
94.7
56.3
0.0059 ± 0.0006
0.0101 ± 0.0002
0.00280.0008 16.7
± 0.5−
(tracedby theFigure 7. Radial velocities obtained from the line centroid of a Gaussian
mentioned OH sky lines. After this correction was applied, the acH2-NW2fg
s16
8.0
19 17 45.771 +19 13 14.40 −110
H2-NW2d98.0
54.1
0.0018 ± 0.0007
0.0141 ± 0.0014
0.00280.0056 20.2 ± 2.0−
fits to the H2 2.122 µm line of the spectra obtained by averaging the signal
curacy reached for the wavelength calibration was better than 0.35
H2-NW2e107.5
51.2
0.0173 ± 0.0023
0.0145 ± 0.0011
0.01040.0048 32.1 ± 2.7−
Figure 8. H2 image of the L723 field (grey-scale) superposed to the contours of the redshifted and blueshifted CO outflows from Lee−1
et al. (2002). The red
within the full aperture of the corresponding slitlet. Offsets are relative to
107.5
56.2
0.0068 ± 0.0007
0.0031 ± 0.0019
0.00320.0074 10.7
−6
for theHH2-NW2f
Å in all the observed wavelength
range (i.e. 4.9 bykmthes crosses.
2scale of
3.2.3± 1.5
Individu
and blue open circles of different sizes represent the radial velocity obtained from the 2.122 µm line at −1
the knot positions marked(ThethepositionofSMA
2.
2.122 µm, and 6.3 km s and 8.3 km s−1 for the 1.644 µm andexposuresof 600s to avoid
the bright sky lines andNotes. a Offsets from the position of SMA 2, α(J 2000) = 19h 17m 53.s 694, δ(J2000) = 19◦ 12′ 19.′′ 68, of the intensity peak in t
velocityis shown
at thesaturation
bottom rightbycorner).
In the followi
1.257 µm [Fe II] emission lines).
to compensate for their time variability. The observations wereimage. For the more extended knots, Cs and H2NW2f+g, the position of the centre of the box used for the proper motion determ
is worth andnotingthat theHH 223b Error
H2 outflow
lies
A check of the reliability of the wavelength calibration in theis given. The x-axis pointsItwestwards,
the y-axis
northwards.
values are√the projected
error in the properfoundmotionatfitsmal
to thperformed repeating the sequence Object-Sky (OSOS). The sky
on to thepairlobes ofoneproperof thetwofit,bipolar
MOS
thr
the Hαand close[S II] linesFig.
10 displays
close-upof Fig.
more detail thewavelength
long-slit
through
HH 223-A
show
epochsobserved. c rms residualof theknotofpositionsin themotionϵ = σ CO1 − outflows
r 2 , where σ is thewith
standard
deviatrange
to thein line
usedspectra
to derive
the kinematics
exposures
were
obtainedain
an empty
field,9 showing
offset bywith∼1 arcmin,d Proper(thelarger,east–westone)
detected
in
L723.
As
a
generaltrend,
HH
223-A
t
line:
a
changeinvelocbehaviourofthevelocityasafunctionoftheangulardistance
to
the
same
behaviour
as
in
the
case
of
the
H
r
is
the
correlation
coefficient.
motionvelocity,
assumingadistance
of
300
pc.
2
(H2 2.122 µm) can be seen in Fig. 2, which shows the line profiles
roughly perpendicularly to the outflow axis, since the observed field−1
the H2 knots with blueshifted velocities are projected on to theFive slitlets
betweentwo condensationsof HH 223-A,
ity of(top∼60
km s and
SMA
Thetargets
spectrumof each knot
averaging theof this line
panels)of the theseOH 2.073µm sky line (bottom
is crowded
and2.ourare extended
and wascloseobtained
to each byother.
blueshifted CO outflow lobe, while the H2 knots with redshifted
ering the knot
easterncondensation(HH the223-A1)
mosttheblueshifted.
signalwithin theaperturesquotedTable
and thecomplex
velocities werepanels) beingof the thespectraobtainedby averaging
signalthe
withinfull
The data
reduction
processof thisdatainset
was4,rather
velocities are projected on to the redshifted CO outflow lobe, as∼22 arcsec in
1.644µm [Fe II]apertureFurthermore,
velocity
trendofshown
by the 2.122µm lineof
throughderivedLIRISDR
from the1.748µm andprocedures
2.122 µm H
2 , andof the slitletthe
labelled
in eachthe panels.
The position/IRAF
standardwith
more
dediand combinedcanbeeasilyvisualized
in
Fig.
8.
This
spatialcoincidence
betweenthat this struct
)
and
the
ionized
lines
(e.g.
the
2.224
µm
and
lines.
Bothemissions,theneutral(traced
by
H
HH
223-A
is
also
found
in
other
H
2
the corresponding rest wavelength has been2marked with the dashedcated correction steps developed in IDL. For the standard
procethe CO and H2 velocity signs gives support to the existence of asion feature’ rII]) appear blueshifted. As expected, the velocities
substructures
found(traced2.406
theappears
two H2 shiftedverticaltheline.
Noteµm).
thatInthecontrast,
emissionfrom therestwithin
dures, weused bythe[Feroutines
available for the LIRIS MOS mode
physical relationship between the H2 and CO outflows.
Herbig–Haro oII] are more blueshifted than those derived for bothderivedfrom [FeTheHH 223-A
were
not case
identified
PV map of the [Fe II] 1.644
wavelength
positionin theof theinHthe
package.
first step was to determine and correct
in the LIRISDR
2 2.122 µm line, while it
Furthermore, we searched for changes in velocity at spatial scales
Fig. 6 (leftwhich in turnaretheconsistent
betweenthem.
II] emission peak coincides with
H2 lines,distortions
line. atTheof the [Fefor
appearsµm
centredtheposition
rest wavelength
the sky line, as should
the geometricalalongspatial axis
in orderto Accordingly
have
sion
lines at 2
smallerthanthelengthoftheslitlets.
In
Fig.
6(blackcontours),
wecounterpart
is foundfor the otherwith what was found in the PV maps (see Fig. 6) and in the H- andbe expected.
HH 223-A1,
while
no clear [FeWe thenconcluded
thatII]thewavelength
calibration
the spectra aligned within the CCD rows. Dome white lamp
2.122µm
emission
1.644µm (redplot
position-velocity
(PV)
maps
of
the
H
2
H2 substructure (HH 223-A2). Hence, we confirm that there are two
K-narrow-band images (López et al. 2010), there is some shift bethrough all the slitlets. From these maps, and comparing with theby the slitlets,
unresolved clumps within HH 223-A, with different excitation contween the spatial distribution of the [Fe II] and H2 emissions: offset
uum substract
H2 K narrow-band continuum-subtracted image of the field (see≥ 1 arcsec(2015)
are found between the peaks of the [Fe II] andditions, which may be originated in shocks with different strengths.
MNRASvalues
447, 2588–2602
H2 emissions in knots HH 223-B, -C (that corresponds to -C NorthIn spite of the lack of detection of H2 emission from the knotin H2 ) and -D. Furthermore, there are some substructures that were
HH 223-B in the narrow-band images, we were able to detect H2emission through the slitlet positioned on this knot (S3). From thedetected in H2 but did not appear in [Fe II] (e.g. HH 223-A2) whilesome [Fe II] substructures were not detected in H2 (e.g. a third conH2 2.122 µm line, we derived the velocity at the intensity peak of thedensation lying between HH 223-F0 and -F1). As it was noted, theseemission, and at a secondary weaker substructure lying to its southfacts may be indicative of shocks of different degree of excitationeast (named HH 223-B-peak and -B-east, respectively; see Table 4).
coexisting along the HH 223-A to -F emission region.
At both positions the emission appears blueshifted, although no
A trend for the velocity can be found, consisting in an increasesignificant variation in velocity along the knot is found. In addition,of the absolute value of the velocity, suggestive of an acceleration,
we detected emission from other H2 lines, both in the H and K bands,
the behaviour being consistent with that of the 2.122 µm emission
from east to west, towards the location of SMA 2. The general trend(see e.g. the PV map of the 1.748 µm emission in Fig. 6). Note,
seems to be broken at a few positions (e.g. at A1 and D knots).
in contrast, the spatial displacement found between the [Fe II] andHowever, this might be expected because of the highly complex
H2 emissions, inferred from their PV maps (Fig. 4). The position ofkinematics of HH 223-A-F already found at optical wavelengths,
Figure 12. Image of the L723 field through the H2 2.122 µm line filter (continuum is not subtracted). The arrows indicate the proper marcsec) with thethe [Fe II] emission peak is coincident (within 0.2
where a fully spatial sampling of the emission was obtained with
derived for each knot. The position of SMA 2 is indicated by a × sign. The scale of the arrows (in km s−1 ) is indicated by the double
it isoffsetpeak of HH 223-B-east, the H2 weaker substructure,
IFS (López et al. 2012). One indication of such a highly complex
bottombut
rightcornerof the map.
by 1.4 arcsec from HH 223-B-peak.
kinematics can be found by examining the PV maps of the 2.122
The slitlet S4 included emission from the intensity peak of theµm line of HH 223-A-F at spatial scales of ∼1 arcsec: as can beto the west of
inhomogeneous environment, have been observe
im-millimetre sources, are also misaligned with rebow-shaped feature labelled HH 223-C in the narrow-band
H2theseen from Fig. 6, a change in velocity can be appreciated along
spectofto itthe(labelled
east–west CO outflow axis. The CO velocity maps of
HH 110, see Kajdic et al. 2012).
age, and from a weaker condensation, south-east
HH 223-A (i.e. in S2). Moving from east to west through the slitlet,
Leework).
et al. (2002)
show the presence of a clump of blueshifted CO
HH 223-C-North and C-East, respectively, in thisA spatial
we found an intensity peak of emission, labelled as knot HH 223-A1
emission,superposedontheredshiftedCO
outlfow
lobe,
aroundshift between the [Fe II] and H2 emissions was already found in thein Table 4, and another secondary emission enhancement (HH 2233.4 Full spatial velocities
location ofinthese
H and K narrow-band images, and also can be theappreciatedthe H2 knots. It might be indicative of a complex
A2 in Table 4), their peaks being separated ∼1 arcsec. From the
The full spatial velocity (Vtot ) and the angle (φ)
emission isbetween
offset molecular clumps at these positions that dePV maps of this work. The peak of the [Fe II]interaction
spectra obtained for these two HH 223-A knots, we found a differviates theproper motions away from the outflow axis. Such
motion and the sky (with φ ≥ 0 towards the o(HHknot
223-Cby ∼1.2 arcsec west from the peak of the H2 emission
ence in velocity of ∼50 km s−1 between them, being the eastern
deviations, attributed to complex interactions of the outflow with antained for the knots for which the two velocity
knot more blueshifted. It is worth noting that HH 223-A presents
North).
two emission peaks separated ∼1 arcsec, partially resolved in the
The slitlet S5 included emission coming from the bow-shaped
Hα images (López et al. 2009). The radial velocities derived fromfeature HH 223-E, and from HH 223-D. We foundan 447,
offsetof
MNRAS
2588–2602(2015)Knot
HH 223-

CO outflow軸の方向にknotsは固有運動を持っている。固有運動は双曲流パタンに従い、knotsは東西で逆方向の運動を持つ。!

!

Summary;!
HH223 H2 outflowの3D力学構造を導出した。その駆動源はYSOのSMA 2であり、最も近いH2 knotsはHH223-K1である。力学構造はH2とCo outflowは関係していることが示され

MNRAS 447, 2588–2602 (2015)た。
ページ 7/10

ck, D. K. Ojha, M. Tamura, H. Linz, M. R. Samalhosh1,5

nomy and Astrophysics, Tata Institute of Fundamental Research, Homi Bhabha Road, Colaba, Mumbai 400 005, India
al Observatory of Japan, Mitaka, Tokyo 181-8588, Japan
for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany
rsité, CNRS, LAM (Laboratoire d’Astrophysique de Marseille) UMR 7326, F-13388 Marseille, France
29.
morphology
and411stellar
content of the Galactic H II region IRAS 16148−5011!
Radio Astrophysics,
Tata“Study
Institute of of
Fundamental
Research, Pune
007, India”!
K. K. Mallick, D. K. Ojha, M. Tamura, H. Linz, M. R. Samal and S. K. Ghosh!
mber 4. Received 2014 December 4; in original form 2014 September 18
http://mnras.oxfordjournals.org/cgi/doi/10.1093/mnras/stu2584!

TION

ABSTRACTAn investigation of the IRAS 16148−5011 region – a cluster at a distance of 3.6 kpc – ispresented here, carried out using multiwavelength data in near-infrared (NIR) from the 1.4 m
Infrared Survey Facility telescope, mid-infrared (MIR) from the archival Spitzer GLIMPSE(Galactic Legacy Infrared Midplane Survey Extraordinaire) survey, far-infrared (FIR) fromthe Herschel archive, and low-frequency radio continuum observations at 1280 and 843 MHzfrom the Giant Metrewave Radio Telescope and Molonglo Survey archive, respectively. A
combination of NIR and MIR data is used to identify 7 Class I and 133 Class II sources in theregion. Spectral energy distribution (SED) analysis of selected sources reveals a 9.6 M⊙ highmass source embedded in nebulosity. However, Lyman continuum luminosity calculation
using radio emission – which shows a compact H II region – indicates the spectral type of
the ionizing source to be earlier than B0-O9.5. Free–free emission SED modelling yields
the electron density as 138 cm−3 , and thus the mass of the ionized hydrogen as ∼16.4 M⊙ .
Thermal dust emission modelling, using the FIR data from Herschel and performing modified
blackbody fits, helped us construct the temperature and column density maps of the region,
which show peak values of 30 K and 3.3 × 1022 cm−2 , respectively. The column density maps
reveal an AV > 20 mag extinction associated with the nebular emission, and weak filamentary
structures connecting dense clumps. The clump associated with this IRAS object is found tohave dimensions of ∼ 1.1 pc × 0.8 pc, and a mass of 1023 M⊙ .
Key words: stars: formation – dust, extinction – H II regions – ISM: individual objects:
IRAS 16148−5011 – infrared: ISM – radio continuum: ISM.

!

対象；IRAS 16148-5011, more
3.6kpc,!
so as high-mass star formation is not very well understood(Zinnecker & Yorke 2007).
観測；NIR (Infrared Survey
Facility telescope 1.4m)!
n activity is known to take place in clusters
IRAS 16148−5011 is an infrared nebula (Fig. 1) in the south3), and as such, !observational
studiesof young
ern sky (α2000 = 16h 18m 35.s 2, δ 2000 = −50◦ 18′ 53′′ ) associated with
MIR
(Spitzer
GLIMPSE)!
which is an IR cluster found using the Two Micron All-Sky Survey
uster regions are imperative, because they serve!
FIR (Herschel archive)!
her investigate various associated processes and(2MASS) data by Dutra et al. (2003). It is located at the Galactic!
low-frequency843
ue to the youth of such regions
and the fact that 1280,
plane(l ∼MHz!
333.◦ 047, b ∼ +0.◦ 037) and is in the vicinity of the wellknown star-forming region RCW 106. Though other star-forming
has still not been dispersed, these clusters can beregions are present nearby, IRAS 16148−5011 appeared to be a
arious theories related to star formation, stellar
relatively isolated region in past mappings (Karnik et al. 2001;
as well as stellar and cloud evolution. Of even
<Conclusion>!
Mookerjea et al. 2004). Kinematic distance estimates to this rere the cluster regions
which harbour high-massClass
I (7),
Class
gion 同定し、SEDから星雲に埋もれた9.6vary from ∼3.3 to 11.9 kpc (near- and far-distance estimates;
se such regions 1.
are few
and farbetween,and II (133)
Molinari et al. 2008), and we adopt the distance of 3.6 kpc from
Msunの大質量星を見つけた。!
Lumsden et al. (2013, based on the spectrophotometric distance of aprominent source – ‘G333.0494+00.0324B’ in their nomenclature.res.in!
!
!

2. 低周波電波放射はcompact HII regionを明らかにし、その放射は南University Press on behalf of the Royal Astronomical Society

側に伸びている。

Lyman continuum luminosityの計算は、ionizing sourceとして(単一のソースを考えたとき)Spectral typeのlower limitとしてB0-O9.5 となった。
HII regionの力学年齢は0.3–0.5 Myr
SEDのmodelingによるFree-free emissionは電子密度138/ccが得ら!

れ、電離されたHIIガス〜１６．４Msun 質量である。!

ページ 8/10

3. 星雲の中心領域で(PDRをトレースする)PAHのring状の構造がcompact HII regionの境界近くで見られた。
３つのNIR,MIR,visible 中心天体がある。6.2, 9.60, 8.5 Msun.

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電波のfluxとMIR SED fittingの不一致から、この領域に埋もれた大質量星が存在し得る!
IRAS 16148−5011

2317

Figure 11. 1 arcmin × 1 arcmin central region of IRAS 16148−5011. The images (left-to-right, top-to-bottom) are J, Ks , 3.6, 4.5, 8.0, and 24 µm. North isup and east is to the left. The symbols on the Ks band image are same as Fig. 7 (green cross is #16 from Table 2). The numbered sources on the 3.6 µm image
are from Table 2. The cyan arrows on the 8.0 µm image mark the ‘ring-like’ morphology. The white patch on the 24 µm image is the saturated region.

effects. Volk et al. (1991), on the basis of low-resolution spectra, hadclassified IRAS 16148−5011 as a PAH source. Bulk of the radio
emission (see Fig. 8), as well as the 24 µm emission from hot dust,
is confined within this ring-like feature as expected (Watson et al.
2008). Some of the emission from this ring-like structure could also
be due to swept-up material by the expanding ionization front of
the H II region.

rejecting the pixel values outside ± 2σ in each iteration, till thefit converged. The same patch of the sky was used for each band.
The background flux level, Ibg , was thus determined as 0.29, 2.67,
1.27, and 0.45 Jy pixel−1 for the 160, 250, 350, and 500 µm images,
respectively. We note that the 70 µm image, though available, was
not used here as the optically thin assumption might not hold trueat this wavelength (Lombardi et al. 2014).
Modified blackbody fitting was subsequently carried out on a
pixel-by-pixel basis using the following formulation (Battersby
et al. 2011; Sadavoy et al. 2012; Nielbock et al. 2012; Launhardt
et al. 2013):

4. Thermal dust emission modeling (Herschel dataより、黒体放射fit)で、温度と柱密度が得られる。peakは30K, 3.3 x 10^22 /cm^2。

柱密度マップは、星雲放射に付随したAv>20 magの減光を明らかにし、またIRAS
4.4 Herschel results

16148-5011とdense
clumpsに接続したweakfilamentary structureがあった。
Thermal emission from cold
dust lies in the FIR wavelength range,
and thus its analysis can be used to obtain the physical parameters

!
!
!
!
!
!
!
!
!
!
!
!

Sν (ν) − Ibg (ν) = Bν (ν, Td )#(1 − e−τ (ν) ),(8)τ (ν) = µH2 mH κν N (H2 ),(9)like dust temperature and column density of a region (Battersby
IRAS天体に付随したクランプは~1.1pc
x 0.8pc,
1023Msunである!
with
et al. 2011; Launhardt et al. 2013). This was carried out by the
K. K. dust
Mallick
et al. whose Rayleigh–SED modelling of 2318
the thermal
emission,
Jeans regime is covered by the Herschel FIR bands (160–500 µm),
using the following order of steps.
First of all, the surface brightness unit for all the images was converted to Jy pixel−1 . Since the PACS image is already in Jy pixel−1 ,this step was only required for the SPIRE images (whose units are in
MJy sr−1 ), and was carried out using the pixel scales for the respective SPIRE bands. Next, the 160–350 µm images were convolved
to the resolution of the 500 µm image (∼36 arcsec, lowest among
all images) using the convolution kernels of Aniano et al. (2011),
and regridded to a pixel scale of 14 arcsec (same as 500 µm image).
Using these final reworked images with same resolution and pixel
scale, a background flux level, Ibg , was determined from a ‘smooth’(i.e. no abrupt clumpy regions) and relatively ‘dark’ patch of thesky. The distribution of individual pixel values in the dark patch of
the sky, for each of the bands, was fitted with a Gaussian iteratively,

where, ν is the frequency, Sν (ν) is the observed flux density, Ibg (ν)
is the background flux in that particular band (estimated above),Bν (ν, Td ) is the Planck function, Td is the dust temperature, # is thesolid angle (in steradians) from where the flux is obtained (just thesolid angle subtended by a 14 arcsec × 14 arcsec pixel here), τ (ν) is
the optical depth, µH2 is the mean molecular weight (adopted as 2.8
here), mH is the mass of hydrogen, κ ν is the dust opacity, and N(H2 )is the column density. For opacity, we adopt a functional form ofκ ν = 0.1 (ν/1000 GHz)β cm2 g−1 , with β = 2 (see Hildebrand 1983,
Beckwith et al. 1990, André et al. 2010). For each pixel, equation (8)was fitted using the four data points, keeping Td and N(H2 ) as freeparameters. Pixels for which the fit did not converge, or the error
was larger than 10 per cent, had their values taken as the median ofeight immediate-neighbour pixels. The final obtained temperature
MNRAS 447, 2307–2321 (2015)

ページ 9/10
Figure 12. (a) Dust temperature map, and (b) column density map of the region around IRAS 16148−5011, derived using the SED fitting to the thermal dust
emission.

!

The Astrophysical Journal Letters, 798:L44 (4pp), 2015 January 10

doi:10.1088/2041-8205/798/2/L44⃝
C 2015. The American Astronomical Society. All rights reserved.

30. “SHADOWS CAST BY AWARP IN THE HD 142527 PROTOPLANETARY DISK” !
S. Marino, S. Perez, and S. Casassus!
SHADOWS CAST BY A WARP IN THE
HD 142527 PROTOPLANETARY DISK
http://stacks.iop.org/2041-8205/798/i=2/a=L44?
key=crossref.ce45c878a03a5d94cabe21b0186f7bac!
S. Marino1,2 , S. Perez1,2 , and S. Casassus1,2!
2015,1ApJL
798,deL44!
Departamento
Astronomı́a, Universidad de Chile, Casilla 36-D Santiago, Chile; smarino@das.uchile.cl
2 Millenium Nucleus “Protoplanetary Disks in ALMA Early Science,” Universidad de Chile, Casilla 36-D Santiago, Chile（チリ大の人たち）!
Received 2014 December 14; accepted 2014 December 15; published 2015 January 7
ABSTRACT
Detailed observations of gaps in protoplanetary disks have revealed structures that drive current research on
circumstellar disks. One such feature is the two intensity nulls seen along the outer disk of the HD 142527 system,
which are particularly well traced in polarized differential imaging. Here we propose that these are shadows castby the inner disk. The inner and outer disk are thick, in terms of the unit-opacity surface in the H band, so that theshape and orientation of the shadows inform on the three-dimensional structure of the system. Radiative transfer
predictions on a parametric disk model allow us to conclude that the relative inclination between the inner andouter disks is 70◦ ± 5◦ . This finding taps the potential of high-contrast imaging of circumstellar disks, and bears
consequences on the gas dynamics of gapped disks, as well as on the physical conditions in the shadowed regions.
Key words: planet–disk interactions – protoplanetary disks – stars: individual (HD 142527)

天体HD142527の原始惑星系円盤のinner (とouter disk)の位置関係等のmodelingをした。!

Motivated by our research on the intra-cavity gas kinematics(Casassus et al. 2013a; S. Casassus et al. 2015a, in preparation),
we considered the possibility of a continuous warp linking
different orientations between the outer and inner disks. The
by observational breakthroughs, such as the direct images ofunderlying parametric disk model is documented in Section 2,
protoplanetarygaps in near-infrared scattered
light. A particu・H-bandのPDI(Polarized
differential
imaging)で二箇所のnullが北(11.5時方向)と南(7時along with the radiative transfer setup we implemented tolarly interesting target is the Herbig Ae/Be star HD 142527, at
calculate the emergent specific intensity fields. Our results are方向)に存在する。Figure
2 (a)!
a distanceof ∼140 pc. The disk surrounding
HD 142527 boasts
presented in Section 3. While the inner disk intercepts stellar
the largest inner cavity known, with a 1′′ radius and seen at a・連続波の観測ではそのPDIで抜けている場所でdust
densitylight, shadowsare cast peakがあり、物がない訳
on the outer disk such that their shape andclose to face-on
orientation, with an inclination of about 24◦
orientationaresensitive
measures of the inner disk orientation.
(Fukagawa et al. 2006; Fujiwara et al. 2006). Improvements in
ではない(sub-milli波観測？)!
We briefly consider the consequences of this finding in Section 4.
Ks -bandcoronography revealed fine structures in the outer disk,
including＜内側円盤＞!
two intriguing intensity nulls (Casassus et al. 2012)
2. PARAMETRIC MODELING
that break the continuity of the outer disk. In this Letter, we
propose that
these
are
shadows
cast
byawarpedinner
disk.
・VLTIの観測で内側円盤があることが分かっている。が配置関係などは分かっていない!
2.1. Underlying physical structure
Recent polarized differential imaging (PDI) of HD 142527
We constructed a synthetic model of HD 142527 inspiredat the H and Ks bands revealed a striking view of its dustby multiwavelength observations and gas kinematics, and dedepleted
gap, outer disk, and spiral structures (Canovas et 1のような状況を主張している。内側円盤が外側
al.
結果として、彼らの主張としてはFigure
tailed enough to provide a comparison point with the near-IR
2013; Avenhaus et al. 2014; Rodigas et al. 2014). These H-band
scattered light images as well as the radio continuum emispolarized
intensity images outline the outer-disk nulls with the degでP.A.=-8±5degをもち、内側円盤が外側円盤に対して相対的なinclinationとして70±5
sion and gas kinematics. As explained in a companion paper
best available detail thus far. One null is found ∼11.5 hr (north(S. Casassus 2015b, in preparation), this synthetic modelof the 円盤に影を落としている。
star), and the other is found at 7 hr (south). The northern
matches the observed spectral energy distribution. Here we docnull is particularly puzzling as it coincides with the location ofument the aspects of the model that are relevant to this report
the dust density peak, whose spatial distribution is shaped intoon the H-band images.
a large crescent (Ohashi 2008; Casassus et al. 2013b; Fukagawa
The model consists of a star at the center of a dusty disk
et al. 2013). Thus, the northern null cannot be interpreted as aseparated into three zones with different density structures,lack of material.
dust grain size distributions, and compositions: an inner disk,
As witnessed by a bright thermal IR central point source,
a gap, and an outer disk. A schematic render of the model iswell in excess of the photospheric emission, the large dustprovided in Figure 1. We model the star using a Kurucz template
depleted cavity of the HD 142527 disk is in fact a gap (Fujiwara
spectrum (Castelli & Kurucz 2003) at a temperature of 7750 K
et al. 2006; Verhoeff et al. 2011; Rameau et al. 2012). Whileand with a radius of 3.3 R⊙ to be consistent with the extinction
knowledge of the outer edge of the gap is limited by theory ratherof AR = 1.3 ± 0.3 reported in Close et al. (2014).
than observational detail, the inner regions of the gap are poorly
The inner disk starts at 0.3 AU and cuts off at 10.0 AU. It is
resolved. A ∼10−7 M⊙ yr−1 stellar accretion rate (Garcia Lopez
described by a surface density proportional to r −1 and a scale
et al. 2006) would quickly deplete the inner disk, whose dust
height of 1.0 AU at 10 AU with a flaring exponent of 1.1. It is
mass is estimated at ∼10−9 M⊙ (Verhoeff et al. 2011). This innercomposed of amorphous carbon and silicate grains with sizes
disk can be thought of as a transient feature of stellar accretion:
between 0.1 to 2.5 µm, with a total dust mass of 5.0 × 10−9 M⊙ .
it is the convergence point of matter being accreted from theWe introduce an inclination of α = 70◦ to this region with respectmass reservoir in the outer disk (Casassus et al. 2013a, 2013b;
to the outer disk midplane.
S. Casassus et al. 2015a, in preparation). Long-baseline optical
The gap spans 120 AU from 10.0 AU. It is described by ainterferometry (VLTI) has resolved this inner disk, and found it
surface density proportional to r −1 and a scale height of 18 AU
to be highly chrystalline (van Boekel et al. 2004). However, its
at 100 AU with a flaring exponent of 1.6. It is composed oforientation and extent remain ellusive.
1. INTRODUCTION
観測事実として、!
The studyof planet formation is rapidly being revolutionized＜外側円盤＞!

!

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1