Young Stars I,II
Lee Hartmann, Smithsonian Astrophysical Observatory
• magnetic flux and primordial stellar fields
• infall and disk accretion
• magnetic fields and turbulence in disks
• winds/jets
• magnetospheric accretion and stellar spindown
Stars form from the collapse
of protostellar gas clouds, r 
104 AU
optical
infrared
Alves, Lada & Lada 2001
The magnetic flux “problem”
(Mestel & Spitzer)
For gravity to overcome magnetic pressure:
(GM/ R2) M  coef.  (d/dR) (B2/8π) (4π R3/3)
GM2 > () B2 R4 =() c
Flux-freezing:   const
(plasma drift t ~ 106 yr,
free-fall t ~ 105 yr)
no reason to expect << c (equipartition)
R ~ 1017 cm; R* ~ 1011 cm; conserve BR2;
Bo ~ 10-5 G,  B* ~ 107 G!
Essentially ALL protostellar cloud magnetic flux must be lost
during star formation (protostars don’t have such B)
Why? low ionization at high  as collapse proceeds, so flux-freezing is not a
good approximation (Umebayashi & Nakano 1988)
The angular momentum “problem”
R ~ 1017 cm; R* ~ 1011 cm; conserve angular
momentum during (nearly) free-fall collapse
 R2  constant
R(final)/R  (o/K)2
Therefore, even if (o/K)2 ~ 0.1,
R(final) ~ 0.01 R ~ 1015 cm ~ 100 AU.
Stars must form from disk accretion
(magnetic flux loss in low-ionization disks)
molecular cloud core undergoes free-fall
collapse to protostar with disk and jet
Why do disks accrete?
Hydrodynamic exchange?
Doesn’t seem to work
Gravitational instability?
May work;
requires massive disk
Magnetorotational instability (MRI)?
Works well when ionization high enough (?)
Magnetorotational Instability?
Side view: initial vertical field
(Balbus & Hawley)
Consistent with
“” disk formalism
(B& Papaloizou)
Disks with very low initial B
 dynamo activity  MRI!
But: dusty protostellar disks
have VERY LOW ionization;
B doesn’t couple to gas
Stone, Balbus, Hawley, Gammie 1996
BUT: low ionization  no magnetic viscosity  no accretion!
T Tauri disk (model):
Thermal ionization (T > 1000K)
X-ray or CR ionization
Dead zone (and layered accretion) (Gammie 1996)
Does any primordial magnetic flux survive infall to disk?
Even if it does, can it survive ohmic diffusion in disk?
What does the turbulence in MRI do?
Can there be any highly organized fossil field in A(p) stars?
Fleming & Stone 2003:
Simulation of shearing box with
dead zone:
MRI operates only in upper layers,
but Reynolds stress extends into
midplane
 “Dead zone” somewhat active,
can accrete?!
Disk accretion can be highly time-variable,
with short bursts of very rapid accretion.
FU Ori; outburst of disk accretion
Disk accretion  10-7 - 10-8M/yr  protostar;

Disk accretion  10-4M/yr  FU Ori object
Why unsteady accretion?
Infall to disk; high velocity
disk accretion; low radial velocity  no reason to balance!
if dM/dt (infall) > dM/dt (accretion):
onto disk
onto star
mass buildup  eventual rapid disk accretion
Outburst sequence (Armitage et al. 2002; Gammie & Hartmann 200?)
matter builds up in dead zone
mass added at
outer edge
(infall)
Grav. Instability  accretion heating  thermal ioniz.  rapid accretion
rapid accretion triggers thermal instability in innermost disk
What happens to the star??
During FU Ori outburst, L(acc) ~ 100 L*;
 Likely advection of large amounts of thermal energy,
(Popham et al 1996)  star expands (but relaxes
quickly if only 0.01 M is added in each outburst?)
Rapid episodic accretion may be typical of the earliest
phases of protostellar formation
Magnetic fields CAN couple to protostellar disks:
Jets/Winds
• Thermal pressure too low to accelerate flows
• Radiation pressure negligible
• Collimation!
Jet seen in [O I]
(accretion-driven)
280 AU
Burrows et al. 1996
Flared disk seen in scattered light:
dust lane obscures central star
bead on a wire analogy
collimation
Alfven surface
Accretion leads to ejection
dM/dt (wind)
= 0.1 dM/dt (acc)
Calvet 1997
Accretion power drives strong mass loss (NOT stellar
winds! Stars without disks do not show detectable mass loss)
FU Ori disk winds
Hartmann & Calvet (1995); accelerating
disk wind results in shifts increasing
with increasing strength (upper levels)
Petrov & Herbig 1992
disk rotation
Winds and turbulence
FU Ori winds are
extremely time-variable;
consistent with complex
disk magnetic field
geometry
FU Ori winds must be
heated to explain H, etc;
numerical simulations of
MRI show waves
propagating upward and
shocking
“Atmospheric” absorption
line profiles show
evidence for sonic
“turbulence” (Hartmann,
Hinkle & Calvet 04)
Blandford &
Payne 1982
Miller &
Stone 2002
HAe/Be
IMTTS: predecessors
of the HAeBe
T Tauri stars:
CTTS= accreting
WTTS=not acc.
T Tauri: (FGKM) pre-main sequence stars with disks
Hartmann 1998
T Tauri star spots (cool);
BIG! (large stellar B)
V410 Tau
Stelzer et al. 2003
(stellar luminosity perturbed? Rosner & Hartmann… - observational problems
Proxies for magnetic fields (activity): enhanced in
pre-main sequence stars - “saturated” behavior (i.e. not
strongly rotation-dependent)
Chromospheric fluxes
X-ray fluxes
(accretion)
Walter et al. 1988
note: x-ray emission not affected
(much) by disk accretion (“T”)
Orion Nebula cluster stars (ages ~ 1 Myr)
Flaccomio et al. 2003
“Saturation” : B or heating efficiency?
T Tauri magnetic fields
BP Tau:
Longitudinal (circular polarization) photospheric B < 200 G;
Mean Zeeman broadening ~ 2.8kG  cancellation!
Circular polarization of He I emission (magnetospheric): 2.5 kG
Johns-Krull
et al. 1999,
2001
Summary of magnetic properties of
pre-main sequence stars
• Spot areas > 30% of stellar surface (non-axisymmetric part)
• Measured field strengths ~ 2kG (average over visible surface!)
• Circular polarization low  cancellation (complex structure)
• Magnetic activity strongly enhanced from solar, “saturated”
Why magnetospheric accretion?
• “Hole” in inner disk (Bertout, Basri, Bouvier 1988)
• Periodic modulation of light from “hot spots” (BBB)
• High-velocity infall (Calvet, Edwards, Hartigan, Hartmann)
• Stellar spindown through “disk locking” (Königl 1991) (?)
• Stellar magnetic fields ~ several kG, strong enough to disrupt
disks (e.g., Johns-Krull, Valenti, & Koresko 1999)
Magnetospheric accretion: line profiles
Königl 1991
(Muzerolle et al. 1998):
line width  (2GM*/R*)1/2
Models for magnetospheric emission
Circularly polarized He I emission
LCP
Johns-Krull et al.
1999
RCP
Accretion power in T Tauri Stars
Classical TTS
Bertout et al. 88;
Kenyon &
Hartmann 87;
Hartigan et al.
90,91;
Valenti et al. 93
Weak TTS
Blue excess (veiling) continuum can be > L*;
 not stellar magnetic activity, but accretion powered;
inner disks (IR emission)  veiling  accretion
Magnetospheric
accretion and outflow
Numerical simulations
show complex accretion
pattern, not always polar,
even when pure aligned
dipole (Miller & Stone
1997)
Tilted dipole  asymmetric streams of accretion:
But: we don’t see implied strong variations of line profiles. Geometry must
be more complicated.
Romanova et al. 2003, 2004
Complex magnetosphere?
Continuum emission: (Calvet & Gullbring 1998)
• very small (~ 1% ) covering factors
• high dM/dt  larger covering factor on star
Line emission (Muzerolle et al);
• high dM/dt  larger magnetosphere area
 Flux tube accretion
The angular momentum problem
If stars accrete most of their mass from disks, why aren’t
they rotating rapidly?
dJ*/dt loss in wind? But then don’t get spin-up to main
sequence (Pleiades)
Solution: transfer J to disk with B (“disk-locking”) (??)
Why do young stars rotate so slowly if they are
formed from disk accretion?
And why faster for lower-mass stars??
Clarke & Bouvier 2000
Disk-star magnetic coupling: does it work?
accreting
non-accreting
Taurus: accreting stars (stars with disks) rotate more
slowly (Bouvier et al., Edwards et al. 1993)
Why do young stars rotate so slowly if they are
formed from disk accretion?
Bimodal? (Herbst et
al. 2002)??
(should plot in log P)
Note: wide range
The angular momentum problem
Accretion implies J(disk)  J(star); how to get rid of it?
Solution 1: different
field lines
problem: field lines
wind up unless
perfect “slippage”
Solution 2: exact corotation, no winding
problem: unrealistic
(axisymmetric, etc.)
detailed assumptions
not very clear
(Collier Cameron & Campbell)
The angular momentum problem
Shu et al. “funnel” flow + x-wind
Lovelace, Romanova, & Bisnovatyi-Kogan 1995
Disk-star magnetic coupling
Generally, field lines wind up
 accretion and spindown
alternate?
 intermingled accreting flux
tubes with spindown field lines?
 limits spindown too much?
(Matt & Pudritz 2004)
Reconnection? Flares? Not clear that accreting TTS
have more activity than non-accreting (weak) TTS
(n.b. Need to heat accreting loops somehow)
Disk dynamo? Opposed field to star?
Accretion, spindown oscillatory
von Rekowski & Brandenburg 2004;
also Goodson, Winglee, Matt
Disk-star magnetic coupling: does it work?
To spin down star, either
wind or disk must carry
away the stellar J!
Disk: to accrete at dM/dt, inner disk must carry away this
angular momentum; assume co-rotation (Keplerian)
s = I* */(dJ/dt) = k2 M* *R*2
dM/dt dRd2
 k2 (M*/dM/dt) (*/K) (R*/Rco)1/2
 0.2  108 yr  (*/K) / 2
so either slow rotation or need very high dM/dt to spin
down in 106 yr
Disk-star magnetic coupling: does it work?
or?? coronal mass ejection-type loss, except using disk material??
Need ~ no angular momentum
loss to explain fast rotators in
Pleiades
(spinup due to
contraction toward
MS; stellar winds
can’t be effective)
Bouvier et al. 1997
But! need spindown to ~
107 years! Disks? But
disks don’t seem to last
quite that long!
Questions about young stars:
• How does the dynamo work in young, completely
convective stars?
• How are magnetic fields distributed over surfaces of
young stars? What happens to surface convection, etc.
when PB ~ Pg (photospheric) everywhere??
• Why is activity “saturated”?
• How is stellar angular momentum regulated?