南極30m級テラヘルツ望遠鏡 (極地研究所, 12Sep. 2018)

星・惑星形成領域におけるサブミリ波偏光観測の展望

Munetake MOMOSE (Ibaraki University)

11. Magnetic fields in Star-forming
regions
2. Polarization observations ofProtoplanetary (& Debris) Disks

2

Taurus@250μm + Optical/IR Pol.
(Palmeirim+ 2013; Heyer+ 2008)
Pattle et al.

10

only, soresults is
ned from
hence dointensity)

The Astrophysical Journal, 845:32 (10pp), 2017 August 10

mJy/beamKandori et al.

Prestellar Core, FeSt1-457
Protobinary, L1333 IRS4A@1.6μm(IRSF; Kandori+2017) @880μm(SMA; Girart+2006)

Figure 5. Polarization vectors after subtraction of the off-core component. The ﬁeld of view is 288 or 0.19 pc at a distance of 130 pc, which is equal to the diameter ofthe core. The background image for this ﬁgure is a zoomed-in version of that used in Figures 1 and 4. The white lines indicate the direction of the magnetic ﬁeld
inferred from parabolic ﬁtting. The scale of the 5% polarization degree is shown at the top.

that the collapse of the core from the current stage could not
create the steep curvature of the magnetic ﬁeld lines in the outerregions of the core, because the collapse proceeds in a runaway
fashion and the free-fall time is long in the diffuse outer region(see, e.g., Fiedler & Mouschovias 1993). The steep magnetic
curvature should thus be created mainly during the formation of
the dense core.
It is noteworthy that the present observational data
successfully trace the dust polarization into the dense part ofthe core because of the linearity of the relationship between the
dust extinction (i.e., reddening of the stellar color) and thepolarization degree (Figure 6). It was found that there is no
turnover in the relationship up to AV » 15mag, with a slope ofPH (H - Ks ) » 4.8% mag−1. This demonstrates that thepolarization of background starlight can be used to trace themagnetic ﬁeld inside cold and dense clouds, despite reports tothe contrary for some dark clouds (e.g., Goodman et al. 1995;
Arce et al. 1998).
In contrast to the present result, Alves et al. (2014a, 2014b)
reported the existence of a kink at AV » 10 mag on the AV
versus PH AV diagram for FeSt 1–457. They also reported thatPH AV is a decreasing function of AV. Their results, however,consider the superposition of polarization both from the core
and the off-core medium, which have different position angles.
Note that the sum of vectors with different position angles can
produce a depolarization effect. We obtained a AV versus
PH AV diagram similar to that obtained by Alves et al.

E-ve~ 0.001pc (200au)

1pc

C1 usingmi 1953)
magnetic
measuredmeasureﬁeld bystrength
od given

B-fi
eld
300au

~ 0.01pc (2000au)

0.19 pc

B-fi
eld

FIG
arro
redsh
prob
main
show
positi
show
Cont
super
inten
fracti
angle
9,…3
show
imag
meas
corre
mode
the
sym
α0(J2
δ0(J2
parab
along
centr

28. M. N. Machida, T. Matsumoto, K. Tomisaka, T.
(2005)
29. M. N. Machida, T. Matsumoto, T. Hanawa, K. T
30. The Submillimeter Array is a joint project betw
and the Academia Sinica Institute of Astron
Smithsonian Institution and the Academia Sinica
and Ministerio de Educación y Ciencia (Spain)
AYA2005-08523-C03-02.

1pc~ 0.1pc

ctor

145au

Figure 6. Relationship between H - Ks color and polarization degree toward
background stars with PH taken after the subtraction of the off-core component.
Only stars with PH dPH . 10 are plotted.

(2014a, 2014b) using the present polarization data before thesubtraction of off-core components. On the basis of the“pre-subtraction” data, PH AV clearly decreases with increasing
5(1 )ispersionangles; ρersion in
iation in
numbere μ is theof orders smaller
Stokes Q
Ostriker

100 AU

Orion @ 880μm(JCMT; Pattle+ 2017)

HD169142, PI @1.6μm
HD142527, PI @ 880μm
Figure 1.
In(ALMA;the left panel,
the colorscale2016)
represents the p(Subaru; Momose+ 2015)
Kataoka+gray contours show the continuum emission, and the white3

the polarization vectors are set to be the same. The levels o1(1) Magnetic fields inStar-forming regions

4Formation of low-mass stars

5

Importance of magnetic (B-)field information of stars and planets
Transportation of angular momentum in a core
inevitable during star formation（Lcore/Mcore ≫ L★/M★）
formation of disks, outflows & jets

Turbulence by MRI in a disk
provide viscosity in an accretion disk
hinder the growth of dust grains

Dissipation of B-field should occur during starformation
BcoreRcore2 ≫ B★ R★2
Ambipolar diﬀusion (Low ρ)→ Ohmic Dissipation (High ρ)
6Observational studies on B-field
B-strength: Zeeman effect
OH, CN, (HI)CCS with large SDs and ALMA (at 40 GHz) in near future ?

B-direction: polarization due to extinction/emission
by aligned dichroic dust particles
Opt. & nIR: extinction in background stars(B ‖ E-vector)
fIR - mm: thermal emission of dust particles (B ⟂ E-vector)

Millimeter & sub-millimeter wavelengths are uniqueB-fields in densest & coldest regions
Ground-based telescope → high resolution + wide FOVs
7

nearbyDebris Disksa rotating disk(including PPDs, DDs)

Taurus/Lupus etc.
↓

Orion/
Serpens↓outerenvelope1”

dense Core

diameter ofindividuala star-forming
clump
filamentclusterforming
clump

3”

JCMT,
SPICA/Safari,
etc.

log
(Size [pc])-3.0 (=200au)

SMA,
IRSF/SIRPOL
etc.

inner envelope(infall + rotation)

toroidmassive star
forming regions

0.1”
ALMA,
Subaru,JWST-2.5-2.0 (=2000au)-1.5-1.0 (=2x104au)

30”

-0.5

log (Spatial DR) ~ 2.5
for most telescopeswith modest observing time

molecular clouds

10pc

100pc

1kpc
8

0.0
10kpc

Distance

1.3mm Survey of Dust Polarization by CARMA

FIGURE 1: (A) Sketch of the axis directions: red/bluearrowsshow
thedirectionoftheredshifted/blueshifted lobes of the molecular outflow,
The Astrophysical Journal Supplement Series, 213:13 (48pp), 2014 July
Hull et al.
probably driven
by IRAS 4B (8), solid lines show themain axis of the magnetic field, and dashed lines
show the envelope axes. The solid triangles show theposition of IRAS 4A1 and 4A2. The small crossshows the centre of(a)the magnetic field symmetry. (B)40
5”=(Stokes I) (b)Contour map of the 877 µm dust emission
NGC 1333-IRAS
4A
B-fields
vs. Cont.
superposed with the color image of the1600au
polarized flux
35
intensity. Red vectors: Length is proportional tofractional polarization and the direction is position
30
angle of linear polarization.
Contour levels are 1, 3, 6,
9,…30×65 mJy Beam-1. The synthesized beam is25
shown in the bottom
left corner. (C) Contour andimage map of the dust emission. Red bars show the+31 13 20
measured magnetic
field vectors. Grey bars
10.5
3 29 10.0
correspond to the best 11.0
fit Right
parabolicmagnetic
field
Ascension (J2000)
model. The fit parameters 10
are the position angle of(c)the magnetic field axis, θPA
the centre
05 =61°±6°,
B-fields:
Smallofvs. Large scales
symmetryof
the14 00
magnetic
field,hmssα0(J2000)=3 29 10.55 ±0.0655
and50
δ0(J2000)=31°13’31.8"±0.4" and C=0.12±0.06 for the45
parabolic form y=g+gCx2, where
the x is the distance40
along the magnetic field axis
35 of symmetry from the1’=
centre of symmetry.
30

Hull et al. (2014); “TADPOL”-survey

L1333 IRS 4A (Class 0); d=320pc

Pol. towards 30 cores and 8
regions forming stars at 2.5”40

35Compare with ≳20” B-fields
with JCMT etc. as well as
small-scale outflow directions

Declination (J2000)

Declination (J2000)

including low-mass Class 0 & I

30

25

◦

+31◦13 20

B field

s

s

h

m

B field

s

11.0s

10.5s

3h29 m10.0s

Right Ascension (J2000)

Declination (J2000)

Outflow∦
Bsmall‖Blarge300au

0.093pc

25
20
15
10
05

13 00

c.f. ) B-vectors derived from
λ=877μm Pol. with SMA(red);Girart et al. (2006)◦+31 12 55

B field
s

13.0

s

12.0

s

11.0

s

10.0

s

09 .0

h

m

3 29 08 .0s

s
n
o
i
t
a
bserv

Right Ascension (J2000)

O

Figure 12. NGC 1333-IRAS 4A. Same as Figure 4. (a) The velocity ranges of the CO(J = 2 → 1) line wing emission are 11.2 to 3.8 km s−1 (redshifted) and −4.6
to −14.2 km s−1 (blueshifted). σSL = 2.41 K km s−1 . (b) σI = 10.9 mJy beam−1 .
(A color version of this figure and associated FITS images and machine-readable tables are available in the online journal.)

91.3mm Survey of Dust Polarization by CARMA
Hull et al. (2014); “TADPOL”-survey

The Astrophysical Journal Supplement Series, 213:13 (48pp), 2014 July

Hull et al.

The Astrophysical Journal Supplement Series, 213:13 (48pp), 2014 July

L1448 IRS 2 (Class 0); d=230pc(b)
25

B-fields vs. Cont.

20

Declination (J2000)

15

10

small-scale
B-fields vs. Outflows

(b)

L1527

B-fields vs. Cont.

15

15

20

20
Declination (J2000)(a)

L1448 IRS 2

15

10Declination (J2000)

25

small-scale
B-fields vs. Outflows(a)L1527 (Class 0); d=140pc

Declination (J2000)

30

Hull et al.

10

10

05

05

+26 ◦03 05

5”=1150au

45 00+26 ◦03 00

23.0s

22.5s

22.0s

21.5s

3h25m21.0s

23.5s

23.0s

Right Ascension (J2000)

50

22.5s

22.0s

3h25m21.5s

(c)

53.5s

B field

4h39 m53.0s

54.4s

Right Ascension (J2000)

Right Ascension (J2000)

B-fields: Small vs. Large scales

40

(c)

54.2s

54.0s

53.8 s53.6 s4h39 m53.4s

Right Ascension (J2000)

B-fields: Small vs. Large scales

35

40

30

35

25

Outflow⟂
Bsmall∦Blarge30

Declination (J2000)

54.0s

45

45

Outflow‖
Bsmall‖Blarge54.5s

B field25
20

1’=0.067pc

15
1020Declination (J2000)

23.5s

B field+30◦45 00

B field+30◦44 55

5”=700au

05

05

45 00
55
50

10
05
03 00
55
50
45
40
35
◦

45

+26 02 30

40
25.0s

24.0s

23.0s

22.0s

21.0s

56 .0s

55.0s

54.0s

53.0s

52.0s

4h39 m51.0s

Figure 17. L1527. Same as Figure 4. (a) The velocity ranges of the CO(J = 2 → 1) line wing emission are 12.9 to 9.7 km s−1 (redshifted) and 2.3 to −0.9 km s−1(blueshifted). σSL = 0.32 K km s−1 . (b) σI = 1.0 mJy beam−1 .
(A color version of this figure and associated FITS images and machine-readable tables are available in the online journal.)

3h25m20.0s

Right Ascension (J2000)
Figure 6. L1448 IRS 2. Same as Figure 4. (a) The velocity ranges of the CO(J = 2 → 1) line wing emission are 18.0 to 7.5 km s−1 (redshifted) and 1.1 to −6.3 km s−1(blueshifted). σSL = 1.39 K km s−1 . (b) σI = 2.0 mJy beam−1 .
(A color version of this figure and associated FITS images and machine-readable tables are available in the online journal.)

B field

Right Ascension (J2000)

B field+30◦44 35

1’=0.041pc

15

10

O

s
n
o
i
t
a
bserv

1.3mm Survey of Dust Polarization by CARMA
Hull et al. (2014); “TADPOL”-survey
90

A subset of objects (high pol.)have consistent B-directions
in both size scales, butothers do not.

|χlarge - χsmall| (deg)(Results)

Δ(B-field angles) vs. Pol. fraction

The Astrophysical Journal Supplement Series, 213:13 (48pp), 2014 July

Hull et al.

90

75

The AstrophysicalJournal Supplement Series, 213:13 (48pp), 2014 July75

morphologydesert
such as NGC? 7538-IRS
becauseof the widely
scatter
in
χ
sm
◦
60
acrossthesource.
Nevertheless,an20
0−in the map has an uncertainty ofdetections where Pc > 2σP .
0.8
45The value |χlg − χsm | was used
tency
between large- and small-sca
m
do
n
0.6
dispersion in |χlg − χsm | is equal toRaχlg added in quadrature.
◦30
high-P
90
Generallyathe
outflow fracangle χo i−
70
0.4
the centersourceof the continuum
source aresides
red
and
blue
outflow
lobes,and tak
15
P.A.s. Of course,
thisissomewhatain uniform
0.2
on the selected velocity ranges for
B-fields
because outflows
can have comple
High-polarization0reportoutfloworientations
in sourc
0
1
2
3
4
5
6
7
8
Low-polarization
0.0
extremely complex. The outflow or
P frac (%)
0
15
30
45
60
75
90
first panel of mostplots in Appendi
fracFigure 1. Large- vs. small-scale B-field orientation|χlg −angle
χsm | between
as a functionof polarization
Projected
outflow
and B-fieldfraction(deg) P frac . Sources are included if they have (1) B-field detections at
Note
that
as
a
test,
we perf
both scales, (2) CARMA polarization detections Pc > 3.5σP , and (3) distances d ! 400 pc. The plotted uncertainty in |χlg − χsm | is equal to theuncertainties
inweighted(asopposedtototal-inten
χsm and χlg added in quadrature, where those uncertainties
reflect
dispersion of the
B-field orientations in each source. The fractional polarization P frac = P / I ,Outflows
vs.thesmall-scale
B-fields◦ line
and
found
that
our m
where P and I are the unweighted
the χ
45sm
exhibitoverallalignmentlg and
1.0 averages of the polarized◦ and total intensities in locations where Pc > 3.5σP . Pointsχbelow
between large- and small-scale fields.
0changed.
Forthelow-mass
cores
p
2
−
(A color version of this figure is available in the online0journal.)
the two weighting schemes resulted
consistency angle |χlg − χsm |.
0.860|χlg − χsm| (deg)Cumulative distribution function (CDF)

0−

45 ◦

45
30
15

Cumulative
Distribution Function
Cumulative distribution function (CDF)

Outflows seem randomlyaligned with B-fields at least
for high-Pfrac sources

Outflows vs. large-scale B-fields

1.0

0

0

1

2

3

4

5

6

7

8 P

(%)

Outflows vs.
Small-scale B
the “polarization hole” in each source (see Section 5.3). The• B-directions (small & large)• Outflows• AM (the axis of rotating disk)
are not always parallel

0−

45 ◦

All maps from the TADPOL survey are publicly available
as FITS images and machine readable tables for each figure in
uncertainty
in the fractional polarization is calculated ratheromddifferently:
Appendix A. For each figure
itis the median of the uncertainties
in the fractional
5. ANALYSIS
AND D
n0.6 we include maps of Stokes I, Q, andaU; bias-corrected polarization intensity Pc ; polarization fraction R polarizationineach
pixel.
◦
Pfrac = Pc /I ; and inferred B-field orientation χsm . Additionally,
paper,
we donottheattempt
to90 that when calculatingInNote
P thisaverageonlymagnifrac we−0
we include FITS cubes of0.4
total intensity (Stokes I) spectral-line
morphology
object.
Rather
7tude
of Pc (and not the orientation
χ of oftheeachB-field)
acrossthedata, as well as machine readable tables listing the R.A., decl., I,
source, which makes our measurements
sensitivetoonly
to depoB-field orientations
deriveconcl
Pc , Pfrac , χsm , and associated uncertainties of each line segment
larization along the line of fromsight (LOS)
or
in
the
plane
of
the
sky l
the ensembleofsources.
The
fracplotted in the figures. These
at scales smaller than the resolutionof ourCARMA
maps.dispe
χsm in Table
1 reflect
the large
0.2files are available in a tar.gz package
available via the link in the figure caption.
We should note that interferometricmeasurementsof frac-T
tions fracacross eachof these objects.
The results for each source are summarized in Table 1. We
tionalpolarization can be problematic
becauseaninterferometer
High-polarizationis necessarily determined by detect
give fitted coordinates of0.0the dust emission peaks, maximum
acts
as a spatial filter, and isinsensitive
to large
scale structure.
Low-polarization
tions
where theobservations
have su
total intensity Ipk , maximum bias-corrected polarized intensity
This makes direct comparisonsoffractionalpolarization
results
may not reflect the B-field orientati
0
15
30
45
60
75
90
from single dish telescopessource.
and interferometers
dif- msmall-scalePc,pk , average polarization fraction P
frac , average
Projected
angle between outflow and B-field (deg)Furthermore,extremely
the B-fields
ficult. For example, in cases where polarized emission (Stokes
B-field orientation χsm , outflow orientation χo , source type, discollapse,
pinching, or outflows, andFigure2. Thick, steppedsize
curvesthe cumulativeQdistributionfunctions but totalor U) is localized,
intensity (Stokes I) is extended,tance d to the source, andsynthesized-beamθbmshow(resolution
wheninterpretingthe source-avera(CDF) of the (projected) angles between the bipolar outflows and the mean
it is possible to overestimate the polarization fraction
with anelement) of the maps. large-scale (top) and small-scale (bottom) B-field orientationsin the low-mass
in
Table
1.
interferometer.
The
comparisonofpolarization
anglesshouldbe
We also 11
tabulate theaverage cores
large-scale
B-field
orientation
protostellar
listed in Table
1. Sources
included in the plot have an asterisk
lessproblematic,
however,asitisunlikelythat
StokesQwouldχlg from the SCUBA, Hertz,
and
SHARP
data.
We
averagedχlg Large-scale B-fields are from(*) next to their outflow orientation in the table.

High-P : random
Low-P : outflows⟂B ?

projected angle between B and outflow

Recent progress (1): New large-scale maps
Ward-Thompson+ (2017); Pattle+ (2017); “BISTRO”-team
JCMT + SCUBA-2/POL-2,
14”-beamat λ=850μm
However, in this work we use P for half-vector selection only, so
The Astrophysical Journal, 846:122 (21pp), 2017 September 10

the effect of our choice of debiasing method on our results is
minimal. The measured magnetic ﬁeld angles are determined fromthe relative values of the Stokes Q and U parameters, and hence do
not depend on the absolute calibration (i.e., the polarized intensity)of the data.

B ⟂ filament vs. B ‖ filament

B-field strength3.estimatedby
Results
We determined the magnetic ﬁeld strength
in OMC1 using
Chandrasekhar-Fermi
methodthe Chandrasekhar–Fermi (CF; Chandrasekhar & Fermi 1953)

B-field map in Orion based on
λ=850μm Pol. imagePattle et al.
North-eastern
filament
OMC1/integral
filament

method. The CF method assumes that the underlying magnetic
equipartitionof energy
betweenﬁeld geometry is uniform,
and that the dispersionof measured
polarization angles (after any necessary correction for measureB-field& turbulence
ment errors)
represents the distortion of the magnetic ﬁeld byturbulent and other motions in the gas.
We determined the plane-of-sky magnetic ﬁeld strength(Bpos) in OMC1 using the formulation of the CF method given
by Crutcher et al. (2004):
Bpos = Q ¢

4pr

sv
Dv
» 9.3 n (H2)
mG,
sq
ásqñ

(1 )

where s is the one-dimensional nonthermal velocity dispersion
a systematic
methodto derivein the gas; s is the dispersion
in polarization
position angles; ρis the gas density; Dv is the FWHM velocity dispersion in<σθ>
Dv = s employed
8 ln (2) ); ás ñ is the typical deviation in
km sis (alsopolarization position angle in degrees; n (H ) is the number(Hildebrand+2009;
Pattle+ 2017)
density of molecular hydrogen (r = m m n (H ), where μ is the
v

q

-1

v

q

2

H

2

mean molecular weight of the gas); and Q¢ is a factor of orderunity accounting for variation in ﬁeld strength on scales smaller
than the beam (labelled Q¢ to distinguish it from the Stokes Q
12parameter). Crutcher et al. (2004) take Q¢ = 0.5 (see Ostriker

Figure 1. A map of the polarization half-vectors in the center of OMC1, with

Recent progress (2): ALMA Pol. maps
Hull+ (2017)

B-fields around Ser-emb8

The Astrophysical Journal Letters, 842:L9 (10pp), 2017 June 20

Hull et al.
d=436±9pc

λ=
870μm
λ=
870μm

λ=
1.3mm
2.89”×2.43”

0.35”×0.32” beam

No hour-glass morphology (weakly magnetized cloud ?)Figure 1. Multi-scale view of the magnetic ﬁeld around Ser-emb8 (aJ2000 =18:29:48.089, dJ2000 =+1:16:43.32). Line segments represent the magnetic ﬁeld
orientation, rotated by 90 from the dust polarization (the length of each segment is identical and does not represent any other quantity). Grayscale is total intensity(Stokes I) thermal dust emission. Panel (a) shows 870 μm JCMT observations (Matthews et al. 2009), (b) shows 1.3 mm CARMA observations (Hull et al. 2014), and(c) shows 870 μm ALMA observations, revealing the magnetic ﬁeld morphology13with ∼10,000, 1000, and 140 au resolution, respectively. For the ALMA data, line
segments are plotted where the polarized intensity P > 3sP; the rms noise in the polarized intensity map sP = 25 μJy beam−1. The dust emission is shown starting at

Recent progress (2): ALMA Pol. maps
Hull+ (2017)
The Astrophysical Journal Letters, 842:L9 (10pp), 2017 June 20

initial mean field

sim
ulat
ion
Hull et al.

SP SD/SPICA
ALMA

random alignment, consistent with the “weak-field” case
Figure 2. Multi-scale projections of simulated data. Shown are the column densities and the magnetic ﬁeld orientation of protostellar cores that formed in our AREPO
simulations. As indicated, the initial mean magnetic ﬁeld is oriented in the vertical direction. The initial magnetic ﬁeld strength increases from the left to the rightcolumns, corresponding to Alfvén Mach numbers %A,mean ‐ field = 35, 3.5, 1.2, 0.35 (see Table 1). The top row shows the full-scale AREPO simulations, centered on
the cores shown in the bottom row. The full simulation boxes are 5.2 pc in extent (cloud scale). The middle row shows zoom-ins between JCMT and CARMA scales(core scale). The bottom row shows zoom-ins of cores at ALMA scales (protostar
14 scale). By allowing the magnetic ﬁeld orientations and zoom boxes to be toggled on
and off, the interactive ﬁgure enables the reader to compare more easily the background grayscale column density maps with the foreground magnetic ﬁeld

Nearby Star-forming regions withmisalignment between B & AM may produce
South-pole Large SD
Two types of outflows ? (Matsumoto+ 2017)

Figure 14. Orientations of the vector directions with respect to the z-axis as a function of the radius for the mean magneti
lines), minor axis of the density distribution (green lines), and outﬂow (black lines) for models M05B01 and M1B025 at
between parallel and antiparallel vectors, oV are plotted for each vector V .

B-field structure in size-scale ≳dense coreschange of field directions insmaller size-scales (ALMA)…
- statistics on protostellar disks- outflows’ structure
field strengths- Chandrasekhar-Fermi method- Other methods (e.g., Koch+ 2012)- need cross-check with Zeeman?

vs. SPICA/SAFARI
wavelength dependence- dust characterization,- alignment mechanism(environmental eﬀects, etc.)

is misaligned with theﬂattened envelope mastellar disk.
4. The internal distributi
cloud cores is nonunifo
disk accretes lumps othe angular momentu
hence, the disk is exp
size. This means that a
may have an orbital
misaligned with thetral star.
5. A strong magnetic ﬁel
infalling envelope; thpressure, and the gas
rim can account for thcondensation observe
cloud core MC27/L
veriﬁed by observatio
between the column
the rim.

Figure 15. Schematic diagram of two types of outﬂows: (a) magnetocentrifugal
wind, and (b) spiral ﬂow. The surfaces represent isodensity surfaces, and thetubes denote the magnetic ﬁeld lines. The arrows indicate the direction of theoutﬂow.

Kawabe, and N. Ohashi fo
λ-dependence
computations were carried o
the Center for Computation
BLAST observations in Vela C molecular clouds
nomical Observatory of Ja
Clustre System (T2K-Toda
Center,at
The University of T
(red) do not show “polarization-minimum”
by JSPS KAKENHI Grant
26400233, 26287030, 2540
λ~350μm (Gandilo+ 2016; Fissel+ 2016)
23244027.

We would like to thank K

infalling envelope, and bipolar outﬂows are ejected. The
nonmagnetized models produce massive circumstellar
disks, one of which undergoes fragmentation at ∼104
years following the formation of a protostar.
2. The radius of the circumstellar disk depends on the initialstrength of the magnetic ﬁeld. Models with a stronger
magnetic ﬁeld produce a circumstellar disk with a smaller
radius. The mass of the disk shows a similar dependence
on Journal,
the magnetic
ﬁeld, (9pp),
where a2016
stronger
The Astrophysical
824:84Juneﬁeld
20 produces aless massive disk. The ratio of disk mass to stellar mass
remains roughly constant at about ∼1%–10%, dependingon the strength of the magnetic ﬁeld.
3. The magnetized models reproduce the outﬂow, which can
be classiﬁed into two types: a magnetocentrifugal wind
and a spiral ﬂow. In the latter, the outﬂow is not alignedwith the rotational axis of the disk. In both the cases, theoutﬂow and the ﬂattened envelope are aligned with themagnetic ﬁeld on that scale. In some models, the outﬂow

Software: SFUMATO (MApMeasurement

The disk radius for eachdensity and velocity distrib
the disk Vd was deﬁned
1215放射は連続波なので，周波数に対する厳しい制約はない。ただし，遠赤外線からミリ波にの周波数依存性はダスト性質の特徴づけに役立つ可能性がある一方で，これまでの観測でAn
Observation
Plan
グができれば非常にユニークかつ有意義である。

れていない [38] ことから，1THz 近辺のサブミリ波帯を含む複数周波数（例えば 400GHz と

ーゲットとしては，ひとまずT ≈frequencies
15K ，AV ≥ 20available
mag，水素原子柱密度に換算
Unique if multiple(e.g., 400 & [39] して

m−2 の分子雲低温高密度領域を想定する。これは，近赤外における吸収減光による偏光観850GHz)に感度あり）や，SPICA で構想されている遠赤外放射偏光観測（比較的高温な領域に感度

assuming T=15K，Av≳20mag., or N(H)
9.4E22 cm
識している。また当面の感度計算にあたっては，偏光度≥ 5% ≳の領域が捉えられるよう，-21%) を目標感度に据えるのが良いだろう。この条件で求めた必要感度を表 1.1 にまとめた

to be complimentary to SPICA

表 1.1: ダスト偏光観測に必要な感度 (total intensity×1%に対するもの)
D=10m
D=30m
周波数

ビームサイズ

必要感度 (1σ)

ビームサイズ

必要感度 (1σ)

(GHz)

(′′ )

(mJy/beam)(′′ )(mJy/beam)

400

18.6

1.11

6.2

0.123

850

8.7

1.95

2.9

0.216

ground-based polarization observations above 850 GHz may be
極望遠鏡で現在想定されている感度と比較すると，比較的リーズナブルな時間で広域マッ
possible only from south pole regions.

16とがわかる。南極サイトでは，他の地上サイトでは偏光観測が実質的に極めて困難な
850G

(2) Polarization observations ofProtoplanetary (& Debris) Disks

17

Polarization in a protoplanetary disk
A new window opened by ALMA
HD142527 at λ=874μm

HL Tau at λ=3.1 mm

100 AUspatialresolution
isthe
criticalrevealIn the left panel,
the colorscale
represents
polarizedto
intensityin units of mJy beamalignment+theself-scattering
contours showsmall-scale
the continuum
emission,andwhite
vectors show thepolarization vectors.
structureofpolarization
vectors
olarization vectors are set to be the same. The18levels of the contours are (3, 10, 30, 100, 300, 60

e 1.

1

theoretical background

Origin of dust polarization at mm-submm
1. Thermal emission of “aligned” grains (Tazaki+ 2017)• Two alignment mechanisms

A. J ‖ B : Larmor precession (B: magnetic field)
B. J ‖ k : Radiative precession (k: net radiation flux)• Radiative alignment (J ‖ k) seems dominant for a large
grains (a > 100μm) in a protoplanetary disk

2. Self-scattering of anisotropic radiation fields by
dust grains (Kataoka+ 2015, 2016a; Yang+ 2016)• High albedo, and, High pol. eﬃciency are required ←prominent only at λ~ (2π)amax ; strong λ-dependence !

19

Two external alignment mechanisms
with radiation
flux

with toroidal
B-field

20

Tazaki et al.
(2017)

Various timescales of related processes ina protoplanetary disk (Tazaki et al. 2017)
Timescale : the shorter is more important(mid-plane)(surface)(black)gaseous damping(red)radiation alignment

magnetic precession(green)10% paramagnetic
inclusion(blue)
10% superparamagnetic inclusion(red)radiation precession
21

theoretical background

Origin of dust polarization at mm-submm
1. Thermal emission of “aligned” grains (Tazaki+ 2017)• Two alignment mechanisms

A. J ‖ B : Larmor precession (B: magnetic field)
B. J ‖ k : Radiative precession (k: net radiation flux)• Radiative alignment (J ‖ k) seems dominant for a large
grains (a > 100μm) in a protoplanetary disk

2. Self-scattering of anisotropic radiation fields by
dust grains (Kataoka+ 2015, 2016a; Yang+ 2016)• High albedo, and, High pol. eﬃciency are required ←prominent only at λ~ (2π)amax ; strong λ-dependence !

22

Condition for polarization due to scattering
Kataoka, Muto, MM et al. (2015)
・For eﬃcient scattering

λ=870 µm (ALMA Band 7)(grain size) ≳ λ/2π

1.2(grain size) ≲ λ/2π

There is a grain size which
contributes most to thepolarized emission

P・For eﬃcient polarization

1.4

Pol. eﬃciency P
albedo ω

1.2

1

1

0.8

0.8

0.6

0.6

0.4

0.4

0.2

0.2

0

0

-0.2
0.001

-0.2
0.01
0.1
Maximum grain size [cm](grain size)

If (grain size) ~ λ/2π, the polarized emission due
to dust scattering is strongest
23

1

ω

1.4

with the magnetic field. However, except for
LMA observations of HD 142527 (Kataoka et al.
obse
and HL Tau (Kataoka et al. 2017), the publishedrvat
ations are too coarse to resolve more than a few
ions
ndent beams across the disk, making it difficult
nguish between scattering and other polarization
nisms.
high-resolution ALMA observations of HD 142527
aoka et al. (2016) resolved polarization for many
independent resolution elements across the disk.
larization was radial throughout most of the disk,s expected for grains aligned with a toroidal field,
ward the edges the morphology changed from raazimuthal, which is more consistent with scatterodels in Kataoka et al. (2016) found that scatterbroadly reproduce the features observed in parts
disk – especially where the polarization orientahange sharply – but not everywhere. A complete
tanding of this interesting case is still missing.
Tau is one of the brightest Class I/II in the sky at
illimeter wavelengths, and thus the polarization
ology can be determined at high resolution withable integration times. Kataoka et al. (2017) folup on the Stephens et al. (2014) observations withobservations of HL Tau. Surprisingly, they foundhe polarization morphology was azimuthal, whichts grainsPolarizationaligned with theirdirections
long axes perpendicuFigure 1.
ALMA polarimetric observations at 3 mm (top,the radiation field,
as predicted
by Tazaki et al.
Kataoka et al.
2017), 1.3 mm (i.e.,(middle),Jand870 µm (bottom),
λ=3.1mm:azimuthal← radiativealignment‖ k)where the red vectors show the >3 polarization morphology (i.e.,
Henceforth, we will call this grain alignment
vectors haveaxis
not beenrotated).
Vector lengths are linearly propornism “alignment
with the radiation
anisotropy.”
λ=0.87mm:
parallelto the minor←
self-scattering
tional to P . The color scale shows the polarized intensity, which isvery di↵erent polarization morphologies observed
masked to only show 3 detections. Stokes I contours in each panel-3.5 ⇥ I , wheremm (0.00 6 resolution,
Stephens
et
al.
2014)
and
3
mm
areashown[3,
10, 25, 50, 100,
200, 325,
500, ∝
750,
1000]
consistent
with
the
case
of
≈
100μm
with
n(a)
a
max for1 for 3 mm, 1.3 mm, and 870 µm,
esolution, Kataoka et al. 2017) suggest that the
I is 44, 154, and 460 µJy bm
respectively.
ology of the polarization emission is strongly de24

HL Tau: Strong λ-dependence

Kataoka et al. (2017, 2015); Stephans et al. (2017; 2014)•••

Protoplanetary Disks/Debris Disks witha South-Pole Single Disk at THz
protoplanetary disks seem sufficiently bright tomake polarization observations at THz
HL Tau: ~10Jy @ λ=450μm (Andrews & Williams 2005)
DM Tau: ~ 1.08 Jy @ λ=350μm (Andrews & Williams 2005)

will not be able to spatially resolve them, but …
polarization will be detected only when the polarization
directions in the disk are rather uniform
λ-dependence of polarization detection — scattering ?

nearby debris disks: Pol may be difficult, but..
β Pic, Fomalhaut, ε Eri (Vega) : “The Fabulous Four”, Pol.?
τ Cet : 5.8mJy at λ=850μm(JCMT) rout = 52au at d=3.65 pc,
can be imaged in Total intensity with better sensitivity
25

in the optical (GO-13726; PI Kalas) where the nearest background source is 0.00 68 west and 0.00 03 north of the ALMA
position. Given that the ALMA beam
radius⇠ 0.00 78 along Right Ascension, it is likely the ALMA sourceis theage =0.44is Gyr;age
= 0.8-1.4Gyr;same background object as observed in optical data.
4 M. Chavez-Dagostino et al.

Fohmalhaut(MacGregor+ 2017)

d=7.66 pc; A4V

ε Eri(Chavez-Dagostino+ 2016)

d=3.22 pc; K2V

LMT at 1.1mm

11” beam

BackgroundsourceS1-S7: MacGreg
corner shows the natural weight beam size of 1.00 56 ⇥ 1.00 15. The rms noise is 14 µJy/beam. (right) The ALMA continuumBackgroundimage
overlaid as contours (white) on the HST STIS image from Kalas et al. (2013). Contour levels are in steps of [5, 15, 25, 35, 45, 55]⇥
the rms noise.
sources

The Astrophysical Journal, 828:113 (8pp), 2016 September 10

Figure1. 1.1 mm LMT/AzTEC
of ϵ Eridani.
The outerin
ringtheis fullylower
detected ata significance of >5.7σ. The central peakFigure 1. (left) ALMA image of the 1.3 mm continuum emission from
Fomalhaut.
The continuum
dashedmapwhite
ellipseleft

is detected at 7.5σ and likely corresponds to the sum of three contributing agents: the stellar photosphere, the stellar upper atmosphereand an (or perhaps two) inner warm disk(s). As many as seven (S1-S7) background objects are detected in this map, of which, four have
S/N>5. The source labeled S1 is the brightest with a 1.1 mm flux of 4.6±0.6 mJy. The smoothed beam size of 10.9′′ FWHM is given by
the white circle in the top-left. As a reference, we include the contour levels for S/N = -3.5, -2, 2, 3.5, 5.0, 6.5, 8.0.

ure 3 we show the positions of the dust ring with respect tothe stationary background sources at diﬀerent epochs correspondent to previous (sub-)mm observations.
To further verify the presence (or lack) of substructures
along the ring, we modeled the ϵ Eridani LMT image using
optically thin debris disk models as presented below.

3.2. Modeling Approacha reasonable extrapolation of the model to the photometric data at other wavelengths. For debris disks this model
can account for important eﬀects such as brightening at
disk ansae and brightness asymmetry for non-axisymmetricdisks.

τ Cet

For a given model viewed at some orientation, a highGiven the clear observed eccentricity in the Fomalhaut debris disk, we construct models that accountresolution
for thedisk orbital
image is first generated. The central stellar
emission of 0.7 mJy (i.e. the expected photospheric flux,
parameters of particles in the disk. A particle orbiting within 4a ϵcircumstellar
disk has both a proper
see nextand
section)forcedis then added and the model is convolved
ERIDANI MODELLING
with the LMT beam. A χ goodness of fit metric is theneccentricity, ep and ef , respectively, as well as a proper and forced
argument of periastron, ! parameand !fcomputed. We within
beginby
an 80×80 area shown in Figure 4, but
We modeled the emission from ϵ Eridani using a p(MacGregor+
2016)
terized model of an optically thin debris disk. This model
only pixels (about 70%) where the emission from either the
populating the complex eccentricity plane defined by these four parameters
following
Wyatt
et
al.
(1999).
The
forcedhas been described in detail in Wyatt et al. (1999) andmodel or the image is significant are used (for details see
Kennedyetal.
(2012).
Basically,a
3D
distributionofdust
Wyatt etforcing
al. 2012). The
eccentricity and argument of periastron, ef and !f , are imposed on
the particles by the massive perturber
thevariable background level means
surface area is generated, which can be viewed from any dithat simple least squares minimization does not necessarrection
to create eccentricity
synthetic images. The
yield parameter,
satisfactory
results,
so in mostcases some by-hand
age =
7.24
Gyr;
eccentricity in the disk, and are free parameters in our model. The
properis surface
alsobrightness
left as a ily
freeof the images is calculated by adding up the emission in inintervention was needed to obtain a smooth and continuous
cells alongorbit;
the line ofthesight,!
whereeach cell emits withbackground.
That is,models with χ values
ep , and describes the additional scatter in the eccentricity of eachdividual
particle’s
a givenepevenis thoughp associatedpc;
G8V
like a blackbody at some temperature that is proportional tolower thand=3.65
those presented are
possible,they remove whatassumed to be randomly distributed from 0 to 2⇡. By assumingthea distancesemi-major
axis,
a,
for
each
particleandrandomcould be astrophysical background behind the disk, in parto the central star. For the resolved√disk compoticular at the SE side, and produce negative residuals in the
we used a temperature law of T = 416/ R K, wheremean anomalies, we iterate to find the true anomaly, f , using the nents
code(Vallado
NW.2007).
While the χThen,is always several times higher than the
Rnewtonmis the disk radiusin AU.from
Thoughast2body
we model the data
at a
2

′′

2

res

26
single wavelength, this temperature was chosen to provide

2

number of degrees of freedom (because our model of the sky

Protoplanetary Disks/Debris Disks witha South-Pole Single Disk at THz
protoplanetary disks seem sufficiently bright tomake polarization observations at THz
HL Tau: ~10Jy @ λ=450μm (Andrews & Williams 2005)
DM Tau: ~ 1.08 Jy @ λ=350μm (Andrews & Williams 2005)

will not be able to spatially resolve them, but …
polarization will be detected only when the polarization
directions in the disk are rather uniform
λ-dependence of polarization detection — scattering ?

nearby debris disks: Pol may be difficult, but..
β Pic, Fomalhaut, ε Eri (Vega) : “The Fabulous Four”, Pol.?
τ Cet : 5.8mJy at λ=850μm(JCMT) rout = 52au at d=3.65 pc,
can be imaged in Total intensity with better sensitivity
27

Summary
Large vs. small scale B-fields and their connection
with disk/outflow structure and their evolution
B-Field’s directions & strengths at various size-scales
wavelength dependence of polarization eﬃciency

Small-scale structure of polarization in
protoplanetary disks has been detected by ALMA
no B-field alignment ? … but, wavelength dependence fora large sample -> dust size

Nearby protoplanetary & debris disks may be
important targets for the Single Disk in South Poleregions
28