THE STAR FORMATIONNEWSLETTER
No. 289 (#26-30)
塚越崇(茨城大)3 Mar. 2017

1. Spiral-driven accretion in protoplanetary discs - III tri-dimensional
simulations
2. Probes of turbulent driving mechanisms in molecular clouds fromfluctuations in synchrotron intensity
3. Q+: Characterising the structure of young star clusters
4. Gaps in Protoplanetary Disks as Signatures of Planets: III.
Polarization
5. Binaries among low-mass stars in nearby young moving groups12Hennebelle et al.: Spiral-driven ac

Spiral-driven accretion in protoplanetary discs - III tri-dimensional simulations
P. Hennebelle et al.

• 星周infallがあるケースの3D計算• これまでの2Dを拡張[Lesur+2015]• RAMSESコードを使用(Cartesian coord.)• 特にspiral armの駆動に着目• 軸対称/非軸対称のexternal accretionを導入• ケプラー回転相当の比角運動量を持つ
質量降着率は同一になるように
Spiral-driven• accretionin protoplanetary discsy(7)he gas is inhickness rd .
lues of Uacc
etion rate is(8)label
NAASR2ASR2h
ASR0.5h
SR2ASNRaccretnoaaasaU00220.5
22

ρacc
00.1
0.1
0.1
0.1
0.1

P00.04
0.04
0.01
0.04
0.04V0111
10resolution
standard
standard
highhigh
standard
standard

Table 1. Summary of the diﬀerent runs performed in thepaper. accret ’a’ means asymmetrical accretion while ’s’means symmetrical. U0 is the value of the accretion speedand ρacc the accretion density (see Eq. 7). P is the accretion parameter (see Eq. 9). V0 is the rotation velocity ofthe accretion flow. The standard and high resolution aredescribed in the text.

分解能限界8• 非軸対称ケース

Hennebelle et al.: S• spiral armの生成 <- incoming flowによるもの• 円盤半径より外側でcumulated mass増加• log(r)<-1.2では時間変化なし=>質量が通過していない• α分布は変化なし=> no accretionと同じ振る舞い

initial disk edge

• infall rateを減らすとcumulated mass減少• 2D計算との比較 [Lesur+2015]• 3Dではαが大きめに出る=>αZのcontribution(Z軸の速度構造)

Fig. 4. Simulations ASR2 (left) and ASR2h (right). Top panels
mass diﬀerence over total mass (see text). Solid lines at r = 0.

• 軸対称accretion -> high-mのspiral armも対称• αは外側は非対称とcomparableだが、内側では減少• mass fluxも非対称と比べると小さくなる• accretion of non-rotating gas-> 円盤が早くshrink, spiral arm生成されず• 観測のspiral armは惑星よりinfalling motionでは

Probes of turbulent driving mechanisms in molecular clouds from fluctuations
in synchrotron
Synchrotron probes of turbulent driving.
intensity
C.A. Herron et al.

• 分子雲乱流の起源は様々• driving parameter ζで切り分ける• compressive, curl-free driving (ζ=0)• solenoidal, divergence-free driving(ζ=1)• シンクロトロン放射強度によってζの導出が可能か数値計算で調査• 3D MHD simulation• FLASH v4コードを使用, 256^3 cells• 初期はuniform gas density,
rho=6.56E-21 g/cm-3• L=2pc, B=10μG, T=10K, cs=0.2km/s• Herron+2016に倣って擬似観測

58Herron et al.

8Herron et al.

• Sol.は平均値が大きくSole. ζ=1Comp.ζ=0Comp.は幅が広い傾向• 平均値と標準偏差(の平均)はζの関数=>起源の切り分け可能• z方向(ほぼ磁場と並行)でsensitive

Figure 2. Histograms of the time-averaged, log-normalised synchrotron intensity for lines of sight along a) the x-axis, b) the y-axis andc) the z-axis. The bottom row depicts the same PDFs as the top row, but with a logarithmic y-axis. The error bars denote the standard
error of the mean number of counts in that bin, calculated over the di↵erent snapshots of the simulation. The simulation with solenoidal
driving (⇣ = 1) is shown in blue, with blue error bars, and the simulation with compressive driving (⇣ = 0) is shown in orange, withblack error bars (overlapping areas appear purple).

deviation of the Gaussian used to smooth the synchrotron
finite angular resolution. Signal-to-noise ratios above 5, andintensity map as a fraction of the width of the image. Thisangular resolutions that are smaller than 5% of the total
expresses the angular resolution as a fraction of the total
cloud diameter, should be sufficient to permit ⇣ to be detersize of the molecular cloud.
mined from measured statistics.
Finally, we note that changes in the noise level or anguWe find that for all noise levels and angular resolutions
lar resolution introduce quantitative changes in the relations
studied, there is always a correlation between the mean andbetween ⇣ and µI , as well as between ⇣ and I , and so thesestandard deviation
log-normalised
synchrotron inten10 of theHerronet al.
observational e↵ects must be taken into account before atsity and the turbulent driving parameter, for ⇣ 2 [0, 0.5]. For
tempting to use measured statistics to determine ⇣.
larger noise levels, or poorer angular resolution, the meanand standard deviation become less sensitive to ⇣, but thisis not significant enough to erase all dependence of these
5.2
Observabilityof Synchrotron
Emissionfrom intensity
Figure 3. The time-averaged a) mean, b) standard deviation, c)skewness,and d) kurtosisof the log-normalised
synchrotron
statistics
on the turbulent driving.
maps as a function of the turbulent driving parameter ⇣. The error bars
denote theClouds
standard error of the mean of the statistic for that
Molecular
Insimulation,
Figure 5 weplot theover
time-averaged
mean andofstancalculated
the di↵erent snapshots
the simulation. The statistics calculated for lines of sight along the x-axis are shownAs z-axis
discussed
in squares.
Section The
1, Protheroe
et al.
andasJones
dard deviationof thealong
log-normalised
intensity
as blue circles,
the y-axis assynchrotron
green triangles,
and along theas redlines of bestfit (2008)
are drawn
dashed andet al.to(2011)
expected
to field,see synchrotron
againstdot-dashedthe turbulent⇣ forcase where or parallellines,driving
for linesparameterof sight thataretheperpendicular
the meanmagnetic
respectively.emission from theSagittarius B molecular cloud, but did not find any detecthere is both noise and finite angular resolution. The syntions in their observations. This implies that the degree tochrotron maps were smoothed to have a FWHM of 3 pixels,
which cosmic rays penetrate this molecular cloud is lower
and the noise levels are specified as the noise-to-signal ratios
within molecular clouds of di↵erent sizes, masses, and magSagittarius B molecular cloud, and synchrotron emission isthan expected.
in the image after smoothing. We find that the relationshiponly emitted by secondary electrons that are produced in thenetic field strengths. In particular, Gabici et al. (2009) foundbetween the mean or standard deviation and the turbulent
Gabici et al. (2007) and Gabici et al. (2009) investithat cosmic ray electrons with energies between 100 MeV
molecular cloud by collisions between very high energy cosdriving parameter persists even in the presence of noise andgate the propagation of cosmic ray protons and electrons
and a few hundred TeV are capable of travelling throughmic ray protons and the molecular gas (Brown & Marscher
giant molecular clouds, because the propagation time for
1977; Marscher & Brown 1978). As discussed by Protheroe
MNRAS 000, 1–14 (2016)
these electrons to travel through the cloud is shorter thanet al. (2008) and Jones et al. (2011), their non-detection of
the timescale for energy loss. This is supported by Boettcher
synchrotron emission from secondary electrons may indicate
Figure 3. The time-averaged a) mean, b) standard deviation, c) skewness, and d) kurtosis of the log-normalised synchrotron intensity
etal. as(2013),
who offoundthat cosmic
raysparametercan pass ⇣.throughthat
it isthemoredifficultforofcosmic
ray ofprotons
to penetratemaps
a function
the turbulent
driving
The error bars
denotestandarderror
the meanthe statistic
for thatmolecular
cloudsseveraltimes
beforelosingasignificant
SagittariusBthanexpected.
simulation, calculated over the di↵erent snapshots of the simulation. The statistics calculated for lines of sight along the x-axis are shown
amountof energy.
ray aselectronsin the energy
range
as blue circles,
alongCosmic
the y-axis
green triangles,
and alongthe z-axis as red squares. The lines of best fit are drawn as dashed andof100 MeV
to 100TeVofemit
emission at
fredot-dashedlines,
for lines
sightsynchrotronthat are perpendicularor parallelto the mean magnetic field, respectively.

• 観測ノイズと分解能の影響はあるがそれでも平均と標準偏差で判別可能• SNR>5および分子雲直径の5%相当の分解能があれば良い

LoS毎• Perseus分子雲をGBTで観測すればシンクロトロンが受かるかも• 9 arcmin分解能で1.7mKが受かればいい• confusion回避にVLAも活用する

Q+: Characterising the structure of young star clusters
racterising
S.E. Jaffa ettheal. structure of young star clusters• いままで用いられて来たQパラメータから発展させ、3つのパラメータを用いてA. P. Whitworth,
O. Lomax
クラスターのフラクタル形状を調べる新しいアルゴリズムを紹介

nd Astronomy, Cardi↵ University, Cardi↵ CF24 3AA, Wales, UK• the fractal dimension D; 分布がclumpyかsmoothかを決める• the number of levels L; cloud中のsubstructureのスケールを決める• the density scaling exponent C; 異なるスケール間の質量の相対分布を決める
ABSTRACT

Many young star clusters appear to be fractal, i.e. they appear to be concentrated in anested hierarchy of clusters within clusters. We present a new algorithm for statisticallyanalysing the distribution of stars to quantify the level of sub-structure. We suggest
that, even at the simplest level, the internal structure of a fractal cluster requires
the specification of three parameters. (i) The 3D fractal dimension, D, measures theextent to which the clusters on one level of the nested hierarchy fill the volume of
their parent cluster. (ii) The number of levels, L, reflects the finite ratio between thelinear size of the large root-cluster at the top of the hierarchy, and the smallest leafclusters at the bottom of the hierarchy. (iii) The volume-density scaling exponent, C =d ln[ n]/d ln[L] measures the factor by which the excess density, n, in a structure ofscale L, exceeds that of the background formed by larger structures; it is similar, but
not exactly equivalent, to the exponent in Larson’s scaling relation between density
and size for molecular clouds. We describe an algorithm which can be used to constrain
the values of (D, L, C) and apply this method to artificial and observed clusters. We
show that this algorithm is able to reliably describe the three dimensional structure ofan artificial star cluster from the two dimensional projection, and quantify the variedstructures observed in real and simulated clusters.

Gaps in Protoplanetary Disks as Signatures of Planets: III. Polarization
H. Jang-Condell• Gap有り/無し円盤の見え方をinclination angleを変えて計算• total intensityとpolarized intensity (at 1µm)の違いに着目
Table 1. Gap parameters for given• 円盤モデル詳細はPaper 1 [Jang-Condell 2008,2009]planet masses.

• M*=1M⦿, R*=2.6R⦿, Teff=4280K• Macc=10-8 M⦿/yr, α=0.01, N(a)∝a-3.5 (0.005µm-1mm)• ガウシアン型ギャップ
RT for density & temperature structures• self-consistent
We calculate the brightness of total scattered light and• face-on diskの振る舞い
models described. Simulated images generated at J, H,

polarized light at 1 µm from the surface of the three disk•••

and K bands produce qualitatively identical results, withslight di↵erences arising from the variation of extinction
ギャップなしはfractionがほぼ一定(rに伴いやや減少)
coefficients and albedos across wavelengths. Compositional and grain size variations can edgeでピーク
also result in variギャップ位置でIもPIも減少し、outer
ations in extinction and albedo, which suggest that the=>scattering
angleではなく、一回/多重散乱のoverall
results presented here apply to a wide range of
強度の比率で決まる(多重散乱のほうが効く)
grain properties as well. The opacities used to generate
the models here correspond to a well-mixed Pollack et al.
outer edgeではtotalbrightness増大(1994) dust composition
with a Dohnanyi (1969) collisional
size
distributionwith
minimum/maximum grain
fractional pol.は減少
size of amax = 1 mm and gas-to-dust ratio of 0.138 (seealso Jang-Condell (2009)). For the record, the extinction
opacities at 1 µm is 10.86 cm2 per gram of gas, with analbedo of 0.9099.

Planet mass1

gap width (w/a)

gap depth (d)72 M (0.23 MJup )210 M (0.65 MJup )

0.11
0.17

0.56
0.84

1 Actual masses used for this work. In the text, the masses
have been rounded to 70 and 200 M

for convenience.

5gap parameters are shown in Table 1 and are calculatedbased on viscous gap opening simulations by Bate et al.
(2003) and Crida et al. (2006). The gaps are partiallycleared, in that the vertically integrated surface densityof the disk is decreased by no more than 60% for themost massive planet modeled. The gaps are modeled as
Gaussian perturbations of the form⌃(r) = ⌃0 (r) 1

d exp[ (ra)2 /(2w2 )] .

(1)

where ⌃0 is the unperturbed surface density profile, d isthe depth of the gap, a is the gap position, and w is thegap width. For a full discussion of the relationship between planet mass, disk properties, and gap parameters,
see Paper I.
2.2. Polarized Scattered Light=>IでもPIでもギャップ位置は決まるが、PIのみ As shown in Paper II, the scattered light brightness ofではギャップのコントラストを過小評価する
3.1. Face-on Disks in Polarized Scattered Lightthe disk surface, assuming single isotropic scattering, isI1s (⌫) =!⌫ µR⇤2 B⌫ (T⇤ )2,(2)• 傾いた惑星なし円盤• 等輝度線は楕円でよくフィットするが、中心は星位置とずれる
no planet• 円盤上層の散乱光を見るため• PIのほうが顕著• anisotropy in P fractionのため• 傾いた惑星ギャップ円盤
1270Mearth• 見かけのギャップ幅が実際とは異なる• minor axisでギャップ幅が狭く見えるは

11. Model images of a gap-less disk observed in scattered total and polarized light at 1 µm at varying inclinations
ずだが、PIでは逆になる
200M
earth
Columns are arranged the same as in Figure 1. From top to bottom: 15 , 30 , 45 , 60 , and 75 inclination. The
are tilted so that the top side of the image is the far side of the disk. In the I and P images, white contour=>円盤がwarpしているように見えるdisks
show isophotes and black dashed lines show best-fit ellipses to those isophotes. Isophotes are chosen to have maximum
Figure 3θ1 = π/2 − β + ifar

near

radius of 5, 10, 15, and 20 AU. In the f images, solid, dotted, and dashed contours show f values of 0.5, < 0.5, and> 0.5, respectively. The values of the contours are indicated on the color scale bar.
• ギャップのminimum & imaximaは楕円でよくβフィットできる(IおよびPIどちらも). Same as Figure 3, for a disk with a partial gap cleared by a 70 M planet. The isophotes here trace the=>ギャップの存在を調べるのに有用
to observerlevel equal to half the brightness of the brightened gap edge along the minor axis toward the far edge obrightness
riFigure 4dthe disk. The two inner contours bracket the gap trough, while the two other ones bracket the brighten far edge othe gap.

• total intensityはnear sideで常に強いが、pol.
i12βrdθ2 = π/2 − β − ifractionはnear sideで低くなる=>near/farはPIではなくpol. fractionで見ることが大事

Figure 9.

Figure 5.

Illustration of scattering angles from opposite
side of an inclined disk. The disk is inclined at an angle
i and the opening angle of the disk isat a distancerd from the star. Light is scattered o↵ the disk surface
toward the observer in a non-symmetric fashion because

Same as Figure 4, for a disk with a partial gap cleared by a 200 Mplanet.

Binaries among low-mass stars in nearby young moving groups
M. Janson et al.

• moving groupにあるlate-typeはまだpre-main sequenceであり、isochronによるクラスター年齢推定に有用• isochronに使われるMstar(Teff)-L*はモデル依存するため不定性大• late-typeは特にモデルの不定性が大きい• バイナリであれば、モデルを介さず星質量の直接導出が可能=> 近傍moving groupにある低質量星の高分解能観測を実行
M. Janson et al.: Young• 61/181天体で伴星を検出, 38天体は新検出=>カタログ化
詳細な物理量導出などはbeyond the scope′thgi
diof
temm
it
inK
us
tiosh
thinbuw
50
raimofcl