Young Stars I,II
• magnetic flux and primordial stellar fields
• infall and disk accretion
• magnetic fields and turbulence in disks
• winds/jets
• magnetospheric accretion and stellar spindown
Lee Hartmann, Smithsonian Astrophysical Observatory
Alves, Lada & Lada 2001
Stars form from the collapse of protostellar gas clouds, r 104 AU
optical
infrared
Essentially ALL protostellar cloud magnetic flux must be lost during star formation (protostars don’t have such B)
no reason to expect << c (equipartition)
R ~ 1017 cm; R* ~ 1011 cm; conserve BR2;
Bo ~ 10-5 G, B* ~ 107 G!
Why? low ionization at high as collapse proceeds, so flux-freezing is not a good approximation (Umebayashi & Nakano 1988)
The magnetic flux “problem” (Mestel & Spitzer)
(GM/ R2) M coef. (d/dR) (B2/8π) (4π R3/3)
For gravity to overcome magnetic pressure:
GM2 > () B2 R4 =() c
Flux-freezing: const(plasma drift t ~ 106 yr, free-fall t ~ 105 yr)
Therefore, even if (o/K)2 ~ 0.1,
R(final) ~ 0.01 R ~ 1015 cm ~ 100 AU.
R ~ 1017 cm; R* ~ 1011 cm; conserve angular
momentum during (nearly) free-fall collapse
R2 constant
R(final)/R (o/K)2
Stars must form from disk accretion(magnetic flux loss in low-ionization disks)
The angular momentum “problem”
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?
Disks with very low initial B
dynamo activity MRI!
Side view: initial vertical field(Balbus & Hawley)
Consistent with “” disk formalism (B& Papaloizou)
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!
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?
T Tauri disk (model):
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
if dM/dt (infall) > dM/dt (accretion):
onto disk onto star
mass buildup eventual rapid disk accretion
Why unsteady accretion?
Infall to disk; high velocitydisk accretion; low radial velocity no reason to balance!
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
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
What happens to the star??
Magnetic fields CAN couple to protostellar disks:Jets/Winds
280 AU
Burrows et al. 1996Flared disk seen in scattered light:dust lane obscures central star
Jet seen in [O I] (accretion-driven)
• Thermal pressure too low to accelerate flows• Radiation pressure negligible• Collimation!
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
disk rotationHartmann & Calvet (1995); accelerating disk wind results in shifts increasing with increasing strength (upper levels)
Petrov & Herbig 1992
Winds and turbulence
FU Ori winds are extremely time-variable; consistent with complex disk magnetic field geometry
Blandford & Payne 1982
Miller & Stone 2002
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)
IMTTS: predecessorsof the HAeBe
T Tauri stars:CTTS= accretingWTTS=not acc.
HAe/Be
T Tauri: (FGKM) pre-main sequence stars with disks
Hartmann 1998
T Tauri star spots (cool);BIG! (large stellar B)
Stelzer et al. 2003
V410 Tau
(stellar luminosity perturbed? Rosner & Hartmann… - observational problems
Proxies for magnetic fields (activity): enhanced inpre-main sequence stars - “saturated” behavior (i.e. not strongly rotation-dependent)
Chromospheric fluxes X-ray fluxes
Walter et al. 1988
(accretion)
note: x-ray emission not affected (much) by disk accretion (“T”)
Flaccomio et al. 2003
Orion Nebula cluster stars (ages ~ 1 Myr)
“Saturation” : B or heating efficiency?
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
T Tauri magnetic fields
• 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”
Summary of magnetic properties of pre-main sequence stars
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
(Muzerolle et al. 1998): line width (2GM*/R*)1/2
Königl 1991
Models for magnetospheric emission
Circularly polarized He I emission
Johns-Krull et al. 1999
LCP RCP
Accretion power in T Tauri Stars
Blue excess (veiling) continuum can be > L*; not stellar magnetic activity, but accretion powered;
inner disks (IR emission) veiling accretion
Bertout et al. 88;Kenyon & Hartmann 87;Hartigan et al. 90,91;Valenti et al. 93
Classical TTS
Weak TTS
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?
Taurus: accreting stars (stars with disks) rotate more slowly (Bouvier et al., Edwards et al. 1993)
accreting
non-accreting
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 linesproblem: field lines wind up unless perfect “slippage”
Solution 2: exact co-rotation, 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)
von Rekowski & Brandenburg 2004;also Goodson, Winglee, Matt
Disk dynamo? Opposed field to star?Accretion, spindown oscillatory
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)
But! need spindown to ~ 107 years! Disks? But disks don’t seem to last quite that long!
Bouvier et al. 1997
Questions about young stars:
• How does the dynamo work in young, completely convective stars?
• How is stellar angular momentum regulated?
• 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”?