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”?