星形成ゼミ
2016/11/18
SFN #286 6-10
工藤哲洋（長崎大）

Late veneer and late accretion to the terrestrial planets
R. Brassera,1,∗, S. J. Mojzsisb,c,1,∗, S. C. Wernerd,1 , S. Matsumurae,2, S. Idaaa Earth Life Science Institute, Tokyo Institute of Technology, Meguro-ku, Tokyo 152-8550, Japanb Department of Geological Sciences, University of Colorado, UCB 399, 2200 Colorado Avenue, Boulder, Colorado 80309-0399, USAr Geological and Geochemical Research, Research Center for Astronomy and Earth Sciences, Hungarian Academy of Sciences, 45 Budaö
H-1112 Budapest, Hungary
d The Centre for Earth Evolution and Dynamics, University of Oslo, Sem Saelandsvei 24, 0371 Oslo, Norwaye School of Science and Engineering, Division of Physics, Fulton Building, University of Dundee, Dundee DD1 4HN, UK

ally accepted that silicate-metal (‘rocky’) planet formation relies on coagulation from a mixture of sub-Mars siz
os and (smaller) planetesimals that dynamically emerge from the evolving circum-solar disc in the first few millar System. Once the planets have, for the most part, assembled after a giant impact phase, they continue to be botude of planetesimals left over from accretion. Here we place limits on the mass and evolution of these plan
onstraints from the highly siderophile element (HSE) budget of the Moon. Outcomes from a combination of Nlo simulations of planet formation lead us to four key conclusions about the nature of this early epoch. First, mato lunar HSE ratio requires either that the late veneer on Earth consisted of a single lunar-size impactor striking
5 Ga, or that it originated from the impact that created the Moon. An added complication is that analysis of luna
he Moon does not preserve convincing evidence for a late veneer like Earth. Second, the expected chondrit
t on Mars is 0.06 weight percent. Third, the flux of terrestrial impactors must have been low (! 10−6 M⊕ Myr−1
melting of Earth’s crust after 4.4 Ga, and to simultaneously match the number of observed lunar basins. This con Hadean eon which is more clement than assumed previously. Last, after the terrestrial planets had fully forPreprint typeset using LTEX style emulateapj v. 5/2/11

THE MASSIVE PROTOSTELLAR CLUSTER NGC6334I AT 220 AU RESOLUTION: DISCOVERY OF FURTHER
MULTIPLICITY, DIVERSITY AND A HOT MULTI-CORE
C. L. B ROGAN 1 , T. R. H UNTER 1 , C. J. C YGANOWSKI 2 , C. J. C HANDLER 3 , R. F RIESEN 4 , R. I NDEBETOUW 1,5
Accepted for publication in ApJ

ABSTRACT
arXiv:1609.07470
We present VLA and ALMA imaging of the deeply-embedded protostellar cluster NGC6334I from 5 cm
•to 原始星の星団NGC6334IをALMAとVLAを使って高分解能で観測1.3 mm at angular resolutions as fine as 0.′′ 17 (220 AU). The dominant hot core MM1 is resolved into
seven components at 1.3 mm, clustered within a radius of 1000 AU. Four of the components have brightness
した（分解能0.17秒：220AU）．
temperatures> 200 K, radii∼ 300 AU, minimum luminosities ∼ 104 L⊙ , and must be centrally heated. We
term this new phenomenon a "hot multi-core". Two of these objects also exhibit compact free-free emission! at longer wavelengths, consistent with a hypercompact HII region (MM1B) and a jet (MM1D). The spatial
kinematics of the water maser emission centered on MM1D are consistent with it being the origin of the high•velocity
bipolar molecular outflow seen in CO. The close proximity of MM1B and MM1D (440 AU) suggests a
MM1というコアは７つの構造に分かれる(1000AUの範囲にある)．
proto-binary or a transient bound system. Several components of MM1 exhibit steep millimeter SEDs indicativeof 内４つは中心に熱せられている（”hoteither unusual dust spectral properties or time variability. In addition
to resolving MM1 and the other hot coremulti-core”）．さらにそ(MM2) into multiple components, we detect five new millimeter and two new centimeter sources. Water masers
areの内の２つ（B，D）にはジェットとHII領域が付随し，protodetected for the first time toward MM4A, confirming its membership in the protocluster. With a 1.3 mm
brightness temperature of 97 K coupled with a lack of thermal molecular line emission, MM4A appears to bea highly
optically-thick 240 L⊙ dust core, possibly tracing a transient stage of massive protostellar evolution.
binaryが示唆される．
The nature of the strongest water maser source CM2 remains unclear due to its combination of non-thermal
continuum and lack of dust emission.
! radio
Subject headings: stars: formation — infrared: stars — ISM: individual (NGC6334I) — radio continuum: ISM– submillimeter: ISM•

その他，MM2などもいくつかの構造に分解される．MM4Aから
1. INTRODUCTION
disks, particularly in the context of binary and multiple sys水メーザーが観測されたので，原始星の星団のメンバーと確認さ
tems. These observations must be undertaken at wavelengths
Massive star formation is a phenomenon of fundamental
long enough to penetrate the high column of obscuring dust.
importanceれた．
in astrophysics, but our understanding of thisprocess is hampered by the heavy extinction and large distances to the nearest sites. Over the past decade, several

The NGC6334 region, known as the Cat’s Paw Nebula, contains multiple sites of high mass (M∗ >8 M⊙ ) star formation(Straw & Hyland 1989; Persi & Tapia 2008; Russeil et al.

The Massive Protocluster NGC6334I

5

Ambipolar diffusion regulated collapse of filaments threaded byperpendicular magnetic fields
C. A. Burge1 , S. Van Loo1 , S. A. E. G. Falle2 , and T. W. Hartquist1
16879v1 [astro-ph.GA] 22 Sep 2016

2

School of Physics and Astronomy, University of Leeds, Leeds LS2 9JT, UK
Department of Applied Mathematics, University of Leeds, Leeds LS2 9JT, UK

Received date; accepted date
ABSTRACT

Context. In giant molecular clouds (GMCs), the fractional ionisation is low enough that the neutral and charged particles are weakly
coupled. A consequence of this is that the magnetic flux redistributes within the cloud, allowing an initially magnetically supported
region to collapse.
Aims. We aim to elucidate the eﬀects of ambipolar diﬀusion on the evolution of infinitely long filaments and the eﬀect of decaying
turbulence on that evolution.
Methods. First, in ideal magnetohydrodynamics (MHD), a two-dimensional cylinder of an isothermal magnetised plasma with initially
uniform density was allowed to evolve to an equilibrium state. Then, the response of the filament to ambipolar diﬀusion was followed
using an adaptive mesh refinement multifluid MHD code. Various ambipolar resistivities were chosen to reflect diﬀerent ratios of
Jeans length to ambipolar diﬀusion length scale. To study the eﬀect of turbulence on the ambipolar diﬀusion rate, we perturbed theequilibrium filament with a turbulent velocity field quantified by a rms sonic Mach number, Mrms , of 10, 3 or 1.
Results. We numerically reproduce the density profiles for filaments that are in magnetohydrostatic and pressure equilibrium with
their surroundings obtained in a published model and show that these equilibria are dynamically stable. If the eﬀect of ambipolar
diﬀusion is considered, these filaments lose magnetic support initiating cloud collapse. The filaments do not lose magnetic flux.
Rather the magnetic flux is redistributed within the filament from the dense centre towards the diﬀuse envelope.
The rate of the collapse is inversely proportional to the fractional ionisation and two gravitationally-driven ambipolar diﬀusion regimes
for the collapse are observed as predicted in a published model. For high values of the ionisation coeﬃcient, that is X ≥ 10−7 , thegas is strongly coupled to the magnetic field and the Jeans length is larger than the ambipolar diﬀusion length scale. Then thecollapse is governed by magnetically-regulated ambipolar diﬀusion. The gas collapses at velocities much lower than the sound speed.
For X ! 10−8 , the gas is weakly coupled to the magnetic field and the magnetic support is removed by gravitationally-dominatedambipolar diﬀusion. Here, neutrals and ions only collide sporadically, that is the ambipolar diﬀusion length scale is larger than theJeans length, and the gas can attain high collapse velocities.
When decaying turbulence is included, additional support is provided to the filament. This slows down the collapse of the filament
even in the absence of a magnetic field. When a magnetic field is present, the collapse rate increases by a ratio smaller than for the
non-magnetic case. This is because of a speed-up of the ambipolar diﬀusion due to larger magnetic field gradients generated by the
turbulence and because the ambipolar diﬀusion aids the dissipation of turbulence below the ambipolar diﬀusion length scale. The

初期条件の密度（色）と磁力線(Tomisaka 2014)

A&A proofs: manuscript no. AA201629039

is given bytφ =A&A proofs: manuscript no. AA201629039

hmic density, ρ/ρS , of the equilibrium configuration for the models listed in Table 1. The density range is from 0.1 tond C3c which have a maximum density of 1000 and 300, respectively. The dashed lines show contour lines for density
30, 50, 100, 200, 300 and 500. The magnetic field lines are shown by the solid lines.

0ρS ≈ 32 with ρS = pext /T 0 the surface
ithin ≈ 5tﬀ,S (where tﬀ,S = (1/4πGρS )1/2
uated with the surface density). Contrary
ion, the filament does not oscillate as it

causes the cloud to lose magnetic flux Fig.
and leads
a denser fila5. (a)toNormalised
central density, ρ/ρS as a function of time forment centre than expected from its initialconditions.
ambipolar diﬀusion regulated collapse of model C3b. (b) Ratio of line
Figure 2 shows the equilibrium density
and magnetic
fieldwith the black dot the value of the equilibrium
mass
line
mass
Fig. to10.critical
Similarfigureas Fig. 5, but with a turbulent velocity field ofconfiguration for the models listed in profile.
Table 1.TheComparingthesesquaresshow
the ratio when
the Thecentraldensityis 1000
Mrms = 10 added to the equilibrium
filament.
purpleline shows

νﬀ tﬀ
χK
≈ 2.141.4
πG

where νﬀ = tﬀ /tni isthe average collision
ionisation fraction ation is given by χ =proportional to X if
free-fall time. Using
we find tφ ≈ 125tﬀ,S≈ 1.25tﬀ,S for X = 1
factor of two with tplain the near-propo
filament. However,break down at somecannot exceed the fr
mum collapse times
the magnetic flux su
ues to zero and we fi
Eq. 12, we expect th
Bailey & Basu (201tion fraction for thebility analysis shows
to-flux ratios equiva
the average ion-neu
the free-fall collapse
This break is alstion of gas and the m
7 show the density p
X = 10−6 , 10−7 andof 250ρS and 1000ρ
at diﬀerent times fo
to represent an insta(1000ρS ) phase of th
the profiles and fieldent collapse rates, inA gas density drop in the inner 6 AU of the transition disk aroundthe Herbig Ae star HD 139614
Further evidence for a giant planet inside the disk?
A. Carmona1,2,3,?,?? , W.F. Thi4 , I. Kamp5 , C. Baruteau1 , A. Matter6 , M. van den Ancker7 , C. Pinte8,9 , A. Kóspál2 , M.
Audard10 , A. Liebhart11 , A. Sicilia-Aguilar12 , P. Pinilla13 , Zs. Regály2 , M. Güdel11 , Th. Henning14 , L.A. Cieza15 , C.
Baldovin-Saavedra11 , G. Meeus3 , and C. Eiroa3,16
Université de Toulouse, UPS-OMP, IRAP, 14 avenue E. Belin, Toulouse, F-31400,
France e-mail:arXiv:1609.06708
andres.carmona@irap.omp.eu
2
Konkoly Observatory, Research Centre for Astronomy and Earth Sciences, Hungarian Academy of Sciences, P.O. Box 67, H-1525
Budapest, Hungary
3
Departamento de Fı́sica Teórica, Universidad Autónoma de Madrid, Campus Cantoblanco, 28049, Madrid, Spain
4
Max Planck Institute for Extraterrestrial Physics, Giessenbachstrasse 1, D-85748 Garching bei München, Germany
5
Kapteyn Astronomical Institute, Postbus 800, 9700 AV, Groningen, The Netherlands
6
Laboratoire Lagrange, Université Côte d’Azur, Observatoire de la Côte d’Azur, CNRS, Boulevard de l’Observatoire, CS 34229,
06304 Nice Cedex 4, France.
7
European Southern Observatory, Karl-Schwarzschild-Str. 2, D-85748 Garching bei München, Germany
8
UMI-FCA, CNRS/INSU, France (UMI 3386), and Dept. de Astronomı́a, Universidad de Chile, Santiago, Chile
9
Univ. Grenoble Alpes, IPAG, 38000, Grenoble, France; CNRS, IPAG, 38000, Grenoble, France
10
Department of Astronomy, University of Geneva, Ch. d’Ecogia 16, 1290 Versoix, Switzerland
11
University of Vienna, Department of Astronomy, Türkenschanzstrasse 17, 1180, Vienna, Austria
12
SUPA, School of Physics and Astronomy, University of St Andrews, North Haugh, St Andrews KY16 9SS, UK
13
Leiden Observatory, Leiden University, P.O. Box 9513, 2300RA Leiden, The Netherlands
14
Max-Planck-Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany
15
Núcleo de Astronomı́a, Facultad de Ingenierı́a, Universidad Diego Portales, Av. Ejercito 441, Santiago, Chile
16
Unidad Asociada Astro-UAM/CSIC

6708v2 [astro-ph.SR] 17 Oct 2016

1

•

!
•

!
•

Herbig Ae 星 HD 139614 の円盤をCOの回転振動輝線で観測した（ESO/VLT CRIRES）．
HD 139614 の円盤の内側（2.3-5.3AU）にはダストの分布にギャップがあることが知られているが，ガス(CO)にはギャップはなかった．ただし，ガス表面密度は内側（<6AU）で極端に少な
く，表面密度の分布は半径の増加関数であった．
ABSTRACT
Context. Quantifying the gas surface density inside the dust cavities and gaps of transition disks is important to establish their origin.
Aims. We seek to constrain the surface density of warm gas in the inner disk of HD 139614, an accreting 9 Myr Herbig Ae star witha (pre-)transition disk exhibiting a dust gap from 2.3±0.1 to 5.6±0.3 AU.
その結果，4AUあたりに木星質量程度の惑星が示唆される．
Methods. We observed HD 139614 with ESO/VLT CRIRES and obtained high-resolution (R⇠90 000) spectra of CO ro-vibrationalemission at 4.7 µm. We derived constraints on the disk’s structure by modeling the CO isotopolog line-profiles, the spectroastrometricsignal, and the rotational diagrams using grids of flat Keplerian disk models.
Results. We detected = 1 ! 0 12 CO, 2!1 12 CO, 1!0 13 CO, 1!0 C18 O, and 1!0 C17 O ro-vibrational lines. Lines are consistentgap from 2.5 to 6 AU

CO p=1-0

26

1024

102

10 3

10
3

10

1022

1022
10 3

1020

1018

1018

1016
100

1016
100

4

1020

10 2

10

-30

-20

-10 0
10 20
Velocity [km s-1]

30
10

a)
NH [cm-2]_ = -1.8
1.2 to 4.0 AU
5.0 to 15 AU
mergedno gap

1.2

500

600

500

600

C18O p=1-0

26

-10 0
10 20
Velocity [km s-1]

1022
1020

1016
100

1.0
-20

400

1018

1.1

-30

300

1024

gap from 4 to 5 AU
1.3

200

10 3

1.1
1.0

Normalized Flux

102
104

1024
NH [cm-2]

1.2

13

CO p=1-0

103

10_ = -1.8
1.2 to 2.5 AU
6.0 to 15 AU
mergedno gap1.3
Normalized Flux

12
26

30

b)
Fig. 5. Line-profiles predicted for models with a gap in the intensity distribution: a) line-profiles expected for a intensity distribution with a gap devoid of gas from 2.5 AU to 6 AU, in orange thecontribution from 1.2 to 2.5 AU, in blue the contribution from 6
to 15 AU, and in red the total line-profile; b) similar plot but fora gap devoid of gas of width 1 AU extending from 5 to 6 AU.
Error bars on the spectrum are 1 .

200

300 400
Tgas [K]

200

300

400

500

600

TRANSVERSE VELOCITY SHIFTS IN PROTOSTELLAR JETS: ROTATION OR VELOCITY ASYMMETRIES?
Fabio De Colle 1 , Adriano H. Cerqueira2 , Angels Riera3
Draft version September 19, 2016

ABSTRACT
Observations of several protostellar jets show systematic di↵erences in radial velocity transverseto the jet propagation direction, which have been interpreted as evidence of rotation in the jets. In
this paper we discuss the origin of these velocity shifts, and show that they could be originated byrotation in the flow, or by side to side asymmetries in the shock velocity, which could be due toasymmetries in the jet ejection velocity/density or in the ambient medium. For typical poloidal jet
velocities (⇠ 100 200 km s 1 ), an asymmetry & 10% can produce velocity shifts comparable tothose observed. We also present three dimensional numerical simulations of rotating, precessing andasymmetric jets, and show that, even though for a given jet there is a clear degeneracy between thesee↵ects, a statistical analysis of jets with di↵erent inclination angles can help to distinguish between
the alternative origins of transverse velocity shifts. Our analysis indicate that side to side velocities
asymmetries could represent an important contribution to transverse velocity shifts, being the mostimportant contributor for large jet inclination angles (with respect the the plane of the sky), and can
not be neglected when interpreting the observations.
Subject headings: hydrodynamics – shock waves – methods: numerical – stars: formation – ISM:
Herbig-Haro objects – ISM: jets and outflows – ISM: individual (HH 30, DG Tau,CW Tau, RW Aur, Th 28)
1. INTRODUCTION

Astrophysical jets can be found in a variety of physicalscales, energies and environments. For instance, planetary nebulae, cataclysmic variables, neutron stars and supernovae sometimes show bipolar jets (e.g., Livio 1999).
In active galactic nuclei, we can see accretion powered
jets emanating from the central engine (Yuan & Narayan
2014). In star forming regions, jets are ubiquitous andcan be found associated with young stellar objects ranging from brown dwarfs to high mass stars (e.g., Li et al.
2014).
Star forming regions like Orion and Taurus show atimation of the amount of angular momentum extracted
from the star-disk system. This is a key parameter thatcan be used to understand the origin of protostellar jets,and might also help to understand the general mechanism responsible for the production of astrophysical jetsat di↵erent scales.
Recent observations of HH jets show transverse velocity shifts in several emission lines. Davis et al. (2000)
first observed a shift in the H2 lines for the molecular
jet HH 212. Bacciotti et al. (2002) observed the DG Tau
micro-jet, much closer to the central source and in atomic
lines, finding a similar shift in several emission lines.

ジェットの傾き角による輝線の見え方の違い
3

shifts in protostellar jets

Rotationy, v),
, can
t per
lving

1

0.2
0

he in2flow ,
magscopeo(4)

sivity
ocity
< r0
vz is
a despectetriesn theby a/Rj ),
nsity
ults.
tione exh reransed as

Normalized Flux

(3)

=10o

0.8
0.6
0.4

(2)

t ver) by

Asymmetry1-100

傾き角が小さいと回転の効果が重要
で，傾き角が小さいと速度の非対称
性が重要

=45o

0.8
0.6
0.4
0.2
0
1
0.8
0.6
0.4
0.2
0

=80o

0

100

200

0

100
-1

200

Radial velocity [km s ]Fig. 2.— Intensity profiles of synthetic emission lines computed inpositions symmetric with respect to the jet axis (black, full line andred, dashed lines), for a jet with rotation (left panels) and with velocity asymmetries (right), for inclination angles ✓ = 10 , 45 , 80
(upper, middle, lower panels respectively) with respect to the planeof the sky. The radial velocity di↵erences, given by the shifts in thesynthetic emission lines, depend strongly on the jet inclination angle. In particular, rotation (asymmetry) produces a larger velocity
shift for small (large) jet inclination angles.

観測結果と合わせるためには，v ⇠ 5 km s

流体の数値シミュレーションで境界

1

Figure 2 also shows that the TVS produced by rotation and velocity asymmetry strongly depend on the jet
inclination angle. In fact, the velocity gradient between
two fluid elements located at positions symmetric withrespect to the jet axis (i.e. at = 0, y = y0 and y = y0 )is given (see equation 1) by:vz ⇠ 15 km svr = vr (y0 , 0) vr ( y0 , 0) == 2v (r0 ) cos ✓ + vz sin ✓からジェットを噴出させ，そのシミュ
1(5)（ジェットの速度
~
100
km/s）
for a

orrewith

where vz = vz (y0 , 0) vz ( y0 , 0) is the transverse gradient of the jet velocity.
If the jet is nearly axisymmetric, vz ⌧ v . In this

レーション結果から輝線を求め，同様の結論を得ている．