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18h29m49.6s49.8s50.0s
Right Ascension (J2000)
+1�1501600
1800
2000
2200
2400
Dec
linat
ion
(J20
00)
400AU
1.5
3.0
4.5
6.0
7.5
9.010.5
Pol
ariz
edflux
den
sity
(mJy
bea
m�
1)
Serpens SMM1
Hull+2016, in prep. See also Hull+2016a
ALMA
18h29m48.5s48.6s48.7s48.8s48.9s49.0s
Right Ascension (J2000)
+1�1605200
5400
5600
5800
Dec
linat
ion
(J20
00)
400AU
0.10
0.15
0.20
0.25
0.30
0.35
0.40
0.45
0.50
Pol
ariz
edflux
den
sity
(mJy
bea
m�
1)
Ser-emb 8(N)
Hull+2016, in prep.
ALMA
18h29m47.5s48.0s48.5s49.0s49.5s
Right Ascension (J2000)
+1�1603500
4000
4500
5000
5500
1700000
0500
Dec
linat
ion
(J20
00)
Ser-emb 8, 8(N)
8
8(N)
(b)
B-field
Ser-emb 8, 8(N)
2000 AU
Hull+2014
Ser-emb 8, 8(N) CARMA 1mm
Will ALMA see toroidally wrapped B-fields (or mm-wave dust scattering?!) at ~25 AU disk scales in these sources?
Stay tuned for ALMA Cycle 3 data!
Proposal 2015.1.00768.S PI: Hull, high priority
0.06ʺ (~25 AU) resolution dust polarization (@ Band 7)
Class 0 protostellar dust polarization
HH 211
REFERENCES Hull et al. 2014, ApJ, 768, 159 (outflows vs. B-fields) Hull et al. 2013, ApJS, in press (data release)
ABSTRACT Magnetic fields are thought to play an important role in the formation of stars. However, that importance has been called into question by previous observations showing misalignment between protostellar outflows and magnetic fields (B-fields), as well as inconsistency in B-field morphology between 10,000 AU and 1,000 AU scales. To investigate these inconsistencies, we used the 1.3 millimeter full-Stokes polarimeter at CARMA to map dust polarization with ~2.5" resolution toward 30 star-forming cores and 8 high-mass star-forming regions as part of the TADPOL survey.
Our main findings are: (1) A subset of the sources have consistent magnetic
field orientations between large (~20") and small (~2.5") scales. Those same sources also tend to have higher fractional polarizations than the sources with inconsistent large-to-small-scale fields. We interpret this to mean that in at least some cases B-fields play a role in regulating the infall of material all the way down to the ~1000AU scales of protostellar envelopes.
(2) Outflows appear to be randomly aligned with B-fields; although, in sources with low polarization fractions there is a hint that outflows are preferentially perpendicular to small-scale B-fields, which suggests that in these sources the B-fields have been wrapped up by envelope rotation.
Bipolar outflows are randomly aligned with magnetic fields at the 1000 AU
scales of protostellar envelopes
CD
F
𝜽outflow – 𝜽B-field (deg)
0 1
0º 90º
Hull et al. 2014, ApJ, 768, 159
Photo credit: Chat Hull
with ALMA [email protected] — Harvard/NRAO — Jansky Fellow — Chat Hull
More info — chathull.com References
• Hull et al. 2016, ApJL, 823, 27 • Hull et al. 2015, JAI, 450005 • Hull et al. 2014, ApJS, 213, 13 • Hull et al. 2013, ApJ, 768, 159 • Machida et al. 2006, ApJ, 645, 1227 • Yang et al. 2016, MNRAS, 456, 2794
Hull+2016, in prep.
an angle of !B ’ 90!. On the other hand, the rotation axis hardlychanges its direction from the initial state and remains directedalong the z-axis. The disk normal is also oriented along thez-axis (i.e., the rotation axis). From Figure 8b, it can be seen thata disk forms by the effect of the rotation and the disk normaldirection coincides with the rotation axis.
Figure 8c shows the central region at the core formation epoch(nc ¼ 2:3 ; 1011 cm#3). It can be seen from this figure that anonaxisymmetric structure has formed and the central core haschanged its shape from a circular disk (bottom panels of Figs. 8aand 8b) to a bar (Fig. 8c, bottom). The magnetic field lines runlaterally, i.e., jBrj, jB"j3 jBzj, in the adiabatic phase (Figs. 8cand 8d ). Figure 8d shows an adiabatic core when the centraldensity has reached nc ¼ 6:9 ; 1014 cm#3. A spiral structure isseen in this figure, which indicates that a nonaxisymmetricpattern has formed, even if no explicit nonaxisymmetric densityperturbation is assumed at the initial stage. (Although the non-axisymmetric patterns also appear in some models of Papers I,II, and III, it should be noted that these patterns are due to anonaxisymmetric perturbation added to the density and magneticfield at the initial stage.) The magnetic field lines are considerablytwisted in Figure 8d . It should be noted that in this model, theinclined magnetic field induces nonaxisymmetric perturbations,on behalf of the initial explicit perturbation.
Figure 11 shows the magnetic field lines, the shape of thecore, and the velocity vectors on the z ¼ 0 plane in the adiabaticphase for model C00. This figure shows that a ring is formed, asfound in Paper III, without any growth of a nonaxisymmetricpattern. In Papers I, II, and III, we assumed a cylindrical cloud inhydrostatic equilibrium, in which the magnetic field and angularvelocity are functions of the radius r in cylindrical coordinates.On the other hand, the cloud is assumed to be spherical with auniform magnetic field and angular velocity at the initial stagein model C00. In spite of these differences, a similar ring struc-ture appears in both models C00 and CS of Paper II. Figure 9(bottom) plots the evolution of the axis ratio against the centraldensity for group C. The axis ratios for models C30, C45, andC60 begin to grow after a thin disk is formed (nc k 5 ; 106 cm#3)and reach "ar ’ 0:5 at the core formation epoch. The axis ratio
grows to "ar ’ 1 at nc ¼ 1012 cm#3 in models C30, C45, andC60, while no nonaxisymmetric pattern appears in model C00.This shows that the nonaxisymmetric pattern arises from theanisotropy of the Lorentz force around the rotation axis. A barstructure is formed by the nonaxisymmetric force exerted by theinclined magnetic field, as shown in Figures 8b–8d . This isconfirmed by the fact that the short axis of the density distribu-tion on the z ¼ 0 plane (the disk midplane) and the bar patternrotate together with the magnetic field lines. The axis ratio (thenonaxisymmetry) grows in proportion to #1=6 (107 cm#3 P nc P1010 cm#3 in Fig. 9, bottom), as found by Hanawa &Matsumoto(1999). Since the lateral component of the magnetic field(jBjsin !0) is large (Fig. 9, bottom), the axis ratio grows more inmodels with large !0.
The evolution of the angles !B, !!, !P, and "B for group C isplotted against the central density in Figure 12. The angle be-tween the magnetic field and z-axis becomes !B ’ 90! even inthe early phase of isothermal collapse for all the models C30,C45, and C60. The rotation axis and the disk normal maintaintheir angles !!; !P ’ 0!. Figures 4 and 12 show that in bothmagnetic- and rotation-dominant models the directions of themagnetic field, rotation axis, and disk normal are qualitativelythe same for models with the same $ and !, irrespective of !0 inthe range 30! $ !0 $ 60!.
Figure 13 shows the magnetic field lines, velocity vectors,and density distribution for the epoch t ¼ 1:52 ; 106 yr (nc ¼1:5 ; 109 cm#3) for model C30. Note that the box scale andlevel of grid are different for each panel. The spatial scale ofeach successive panel is different by a factor of 4, and thus thescale between Figures 13a and 13d is different by a factor of 64.The magnetic field has an angle !0B % 30! in Figure 13a, where!0B is defined as the angle between the volume average magnetic
Fig. 11.—Same as Fig. 1, but for C00 at the core formation epoch.
Fig. 10.—Same as Fig. 5 (left), but for model C45. The lower left inset is anenlarged view of the center.
CLOUD EVOLUTION WITH OBLIQUE FIELD 1237No. 2, 2006
Future
?
8H
.Yan
get
al.
Figu
re5.
Mod
elfo
rthe
pola
rizat
ion
ofH
LTa
udi
sk.P
lotte
dar
eth
epo
lar-
ized
inte
nsity
(col
orm
ap,i
nun
itsof
2⌃2 0
abs
sca⌫
2 kT0/
c2 )and
pola
rizat
ion
vect
ors(
line
segm
ents
,with
leng
thpr
opor
tiona
lto
the
pola
rizat
ion
frac
tion
ofth
eto
tali
nten
sity
)for
i=
45� .
Ach
arac
teris
ticsc
atte
ring
optic
alde
pth
�⌧ s,c⇡
0.07
isne
eded
tobr
ing
the
high
pola
rizat
ion
frac
tion
ofte
nsof
perc
entr
elat
ive
toth
esc
atte
red
inte
nsity
(sho
wn
inth
elo
wer
-left
pane
lof
Fig.
2)do
wn
toth
eob
serv
edle
velo
f⇠1%
rela
tive
toth
eto
tali
nten
sity
.
the
low
pola
rized
inte
nsity
near
the
edge
.Thi
stre
ndw
aspr
esen
tin
our
mod
elw
hen
only
the
scat
tere
dra
diat
ion
was
cons
ider
ed(s
eeFi
g.2
and
3),b
utw
asw
ashe
dou
tby
the
tota
lint
ensi
tyin
Fig.
5.Th
ese
mi-a
naly
ticth
eory
that
we
have
deve
lope
dso
faru
nder
the
assu
mpt
ion
ofge
omet
rical
lyth
indi
sk,o
ptic
ally
thin
emis
sion
,an
don
lyR
ayle
igh
scat
terin
gis
inde
pend
ento
fthe
deta
iled
prop
er-
ties
ofdu
stgr
ains
.Thi
sin
depe
nden
cem
akes
the
broa
dag
reem
ent
betw
een
ourm
odel
and
the
mai
npo
lariz
atio
nfe
atur
esob
serv
edin
HL
Tau
rath
erro
bust
.In
the
next
sect
ion,
we
will
tryto
putc
on-
stra
ints
onth
egr
ain
size
dist
ribut
ion
inth
eH
LTa
udi
sk,w
hich
ism
uch
mor
eun
certa
in.
4IM
PLIC
ATIO
NS
ON
GR
AIN
SIZE
4.1
Scat
teri
ngO
paci
tyan
dth
eN
eed
for
Larg
eG
rain
s
The
mai
nfr
eepa
ram
eter
ofth
efin
alm
odel
for
the
HL
Tau
disk
pola
rizat
ion
disc
usse
din
§3
isth
ech
arac
teris
ticsc
atte
ring
op-
tical
dept
h�⌧ s,c
,w
hich
cont
rols
the
pola
rizat
ion
frac
tion.
Inor
-de
rto
prod
uce
the
obse
rved
pola
rizat
ion
frac
tion
of⇠
1%,
ava
lue
of�⌧ s,c⇡
0.07
isre
quire
d.Th
isva
lue
yiel
dsa
scat
terin
gop
acity s
ca⇡
10�3
cm2
g�1
(cro
ssse
ctio
npe
run
itto
tal,
rath
erth
andu
st,m
ass)
at1.
3m
m(th
ew
avel
engt
hof
the
HL
Tau
disk
pola
rizat
ion
obse
rvat
ion)
,usi
ngth
ech
arac
teris
ticco
lum
nde
nsity
⌃0=
68g
cm�2
from
the
best
-fitd
isk
mas
sof
0.13
M�
ofK
won
etal
.(20
11).
This
scat
terin
gop
acity
can
putc
onst
rain
tson
the
grai
nsi
zedi
strib
utio
n,al
thou
ghth
eyde
pend
onth
edu
stco
mpo
sitio
n,w
hich
isun
certa
in.A
san
illus
tratio
n,w
eco
nsid
erth
em
odel
ofdu
stgr
ains
adop
ted
byK
atao
kaet
al.(
2015
),w
hich
are
sphe
res
with
am
ixtu
reof
silic
ate
(8%
),w
ater
ice
(62%
)and
orga
nics
(30%
).A
llfr
actio
nala
bund
ance
sar
ein
volu
me
and
are
take
nfr
omPo
llack
etal
.(19
94).
We
assu
me
aca
noni
calg
as-to
-dus
tmas
sra
tioof
100,
and
use
the
Mie
theo
ryto
calc
ulat
eth
eab
sorp
tion
and
scat
terin
gop
aciti
es(B
ohre
n&
Hu↵
man
1983
).Th
ein
ferr
edsc
atte
ring
opac
-ity
corr
espo
nds
toa
grai
nra
dius
a=
37µ
mfo
rgr
ains
ofa
sing
le
size
.For
theM
RN
-type
pow
er-la
wdi
strib
utio
nn(
a)/
a�3.
5(M
athi
set
al.1
977)
,we
obta
ina
max
imum
grai
nsi
zeof
a max=
108µ
m.
The
incr
ease
ofth
ism
axim
umov
erth
esi
ngle
size
case
com
esfr
omav
erag
ing
over
the
grai
nsi
ze.I
nbo
thca
ses,
the
dim
ensi
onle
sspa
-ra
met
erx=
2⇡a/�⌧
1,so
that
the
Ray
leig
hlim
itus
edfo
rtre
atin
gth
esc
atte
ring
inth
epr
evio
usse
ctio
nsis
self-
cons
iste
nt(s
eeFi
g.6
belo
w). In
sum
mar
y,to
repr
oduc
eth
e⇠
1%po
lariz
atio
nfr
actio
nob
-se
rved
inth
edi
skof
HL
Tau
thro
ugh
dust
scat
terin
g,th
egr
ains
mus
thav
egr
own
tote
nsof
mic
rons
(the
exac
tval
uede
pend
son
the
assu
med
grai
nsi
zedi
strib
utio
nan
dco
mpo
sitio
n).H
owev
er,t
his
pict
ure
isco
mpl
icat
edby
the
opac
itysp
ectra
lind
ex�
infe
rred
for
HL
Tau,
asw
edi
scus
sne
xt.
4.2
Opa
city
Spec
tral
Inde
x�
and
the
Nee
dfo
rLa
rger
Gra
ins?
Kw
onet
al.(
2011
)ob
tain
eda
best
-fitv
alue
0.73
for
the
spec
tral
inde
x�
ofth
edu
stop
acity a
bs/⌫�
for
the
HL
Tau
disk
base
dm
ostly
onC
AR
MA
obse
rvat
ions
at1.
3an
d2.
7m
m.I
tis
inag
ree-
men
twith
the
spat
ially
aver
aged
valu
eob
tain
edfr
omA
LMA
ob-
serv
atio
nsfr
om0.
87to
2.9
mm
(ALM
APa
rtner
ship
:B
roga
net
al.
2015
).Th
isva
lue
issi
gnifi
cant
lylo
wer
than
the
typi
cal
ISM
valu
eof�⇠
1.5�
2.Th
edi↵
eren
ceis
usua
llyta
ken
asev
iden
cefo
rgr
ain
grow
thto
mill
imet
ersi
zeor
larg
er(T
esti
etal
.20
14),
alth
ough
othe
rin
terp
reta
tions
are
poss
ible
.For
exam
ple,
Ric
ciet
al.(
2012
)sh
owed
that
ava
lue
of�⇠
1or
low
erca
nbe
obta
ined
with
outm
m/c
msi
zed
grai
nsif
part
ofth
edi
skis
optic
ally
thic
k.So
me
supp
ort
for
this
poss
ibili
tyis
prov
ided
byth
esp
atia
llyre
-so
lved
dist
ribut
ion
of�
deriv
edfr
omth
eA
LMA
data
,whi
chsh
ows
�⇠
0in
dica
tive
ofop
tical
lyth
ick
emis
sion
atth
ece
ntra
lcon
tin-
uum
peak
and
two
rings
(B1
and
B6,
ALM
APa
rtner
ship
:Bro
gan
etal
.201
5,se
eth
eir
Fig.
3).A
noth
erpo
ssib
ility
isth
atth
ein
dex
�is
sens
itive
tono
tonl
yth
esi
zebu
tals
oth
esh
ape
ofth
egr
ains
.In
deed
,Ver
hoe↵
etal
.(20
11)
was
able
tore
prod
uce
the
spec
tral
ener
gydi
strib
utio
n(S
ED)
ofth
edi
skof
HD
1425
27(w
ith�⇠
1in
the
mill
imet
erre
gim
e)w
ithirr
egul
argr
ains
ofsi
zes
upto
only
2.5µ
m;t
hegr
ain
shap
ew
astre
ated
with
the
dist
ribut
ion
ofho
llow
sphe
res
Min
etal
.(20
05).
The
grai
nsin
ferr
edin
ourm
odel
ofdu
stsc
atte
ring-
indu
ced
pola
rizat
ion
fort
heH
LTa
udi
skha
vea
sign
if-ic
antly
larg
erm
axim
umsi
ze(o
ford
erte
nsof
mic
rons
).Th
eym
ayst
illbe
able
tore
prod
uce
the
obse
rved
(ave
rage
d)op
acity
spec
tral
inde
xof�⇠
0.73
ifth
egr
ains
are
irreg
ular
and/
orpa
rtof
the
disk
isop
tical
lyth
ick.
Det
aile
dex
plor
atio
nof
this
poss
ibili
tyis
beyo
ndth
esc
ope
ofth
epr
esen
twor
k.If
larg
e,m
m/c
msi
zed,
grai
nsar
ere
spon
sibl
efo
rthe
rela
tivel
ylo
wva
lue
of�
obse
rved
inth
eH
LTa
udi
sk,i
tis
natu
ral
toas
kw
heth
erth
eyca
npr
oduc
ea
pola
rizat
ion
patte
rnth
atm
atch
esth
eob
serv
edon
eth
roug
hsc
atte
ring.
Itis
unlik
ely,
beca
use
the
key
topr
oduc
ing
the
obse
rved
patte
rnis
the
pola
rizat
ion
degr
eeof
the
scat
tere
dlig
htpe
akin
gne
ar90�
(as
inth
eR
ayle
igh
limit)
,and
this
requ
irem
enti
sno
tsat
isfie
dfo
rmm/c
msi
zed
grai
ns.F
orex
ampl
e,fo
rthe
grai
nm
odel
adop
ted
byK
atao
kaet
al.(
2015
),th
epo
lariz
a-tio
nde
gree
(defi
ned
asth
era
tioof
the
two
elem
ents
inth
esc
at-
terin
gm
atrix
,�Z 1
2/Z 1
1,w
hich
ises
sent
ially
the
pola
rizat
ion
frac
-tio
nbu
tcan
beei
ther
posi
tive
orne
gativ
e)is
near
lyze
roat
0.87
mm
fora
llsc
atte
ring
angl
esex
cept
arou
nd13
5�,w
here
itre
ache
sa
(neg
ativ
e)“p
eak”
valu
eof⇠�0.2
for
a max=
1m
man
d1
cm(s
eeth
erig
htpa
nelo
fth
eir
Fig.
2).T
hene
gativ
eva
lue
iskn
own
asth
epo
lariz
atio
nre
vers
al(e
.g.,
Mur
akaw
a20
10;K
irchs
chla
ger&
Wol
f201
4)w
hich
,tog
ethe
rwith
the
shift
ofth
epo
lariz
atio
n“p
eak”
MN
RA
S00
0,1–
12(2
015)
Yang+2016
?
x [AU]
y[AU]
shock
shock
shock
shock
shock
0 0.5 1 1.5 2 2.5 3x 104
0
0.5
1
1.5
2
2.5
3
x 104
log10(ρ)[g
cm−3]
−20.5
−20
−19.5
−19
−18.5
−18
−17.5
−17
−16.5
Hull, Mocz, Burkhart, Hernquist+2016, in prep.
2000 AU
10000 AU
Compression of B-fields along
walls of outflow cavities?
18h29m47.9s48.0s48.1s48.2s48.3s
Right Ascension (J2000)
+1�1604000
4200
4400
4600
Dec
linat
ion
(J20
00)
400AU 0.08
0.16
0.24
0.32
0.40
0.48
0.56
0.64
Pol
ariz
edflux
den
sity
(mJy
bea
m�
1)
Ser-emb 8
Hull+2016, in prep.
Accretion along magnetized filaments?
ALMA
outflows
magnetic fields
850 micron dust emission
sources
ALMA 850 micron dust emission
AREPO simulation
ALMA figure panels (top-left, top-right, bottom-right): 850 micron observations of dust polarization toward Serpens SMM1, Ser-emb 8(N), and Ser-emb 8, three Class 0 young stellar objects in the Serpens star-forming region (d = 415 pc). Yellow line segments are magnetic field orientations, rotated by 90° from the polarization orientations. Grayscale is polarized thermal dust emission. Black contours are total intensity (Stokes I) dust emission, starting at 3𝜎 and increasing logarithmically (SMM1 has a gray 3𝜎 contour to highlight a “hole” in the dust in the SE). Blue and red arrows indicate outflow orientations. The black ellipse in the lower-left is the synthesized beam (resolution element), which is ~0.3ʺ (or ~100 AU). Data are from ALMA Cycle 2 project 2013.1.00726.S. Figures were produced using APLpy.CARMA/ALMA/simulation comparison (bottom-left): comparison of CARMA data (right), ALMA data (left), and an AREPO simulation (bottom). CARMA: black and gray line segments are inferred magnetic field orientation; contours are 1mm dust emission; arrows indicate outflows. ALMA: colorscale is total intensity 850 micron dust emission at the same scale as the CARMA map. Simulation: a slice from a simulation of a star-forming core produced using AREPO (credit: Phil Mocz); colorscale is gas density; white line segments are magnetic fields; magenta vectors are velocity vectors; white line is a shock front created by colliding flows.
ALMA & CARMA Same sources!