1. 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

2. Alves, Lada & Lada 2001 Stars form from the collapse of protostellar gas clouds, r  10 4 AU optical infrared

3. 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 ~ 10 17 cm; R * ~ 10 11 cm; conserve BR 2 ; B o ~ 10 -5 G,  B * ~ 10 7 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/ R 2 ) M  coef.  (d/dR) (B 2 /8π) (4π R 3 /3) For gravity to overcome magnetic pressure: GM 2 > () B 2 R 4 =()  c Flux-freezing :   const (plasma drift t ~ 10 6 yr, free-fall t ~ 10 5 yr)

4. Therefore, even if (  o /  K ) 2 ~ 0.1, R(final) ~ 0.01 R ~ 10 15 cm ~ 100 AU. R ~ 10 17 cm; R * ~ 10 11 cm; conserve angular momentum during (nearly) free-fall collapse  R 2   constant R(final)/R  (  o /  K ) 2 Stars must form from disk accretion (magnetic flux loss in low-ionization disks) The angular momentum “problem”

5. molecular cloud core undergoes free-fall collapse to protostar with disk and jet

6. 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 (?)

7. 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

8. Stone, Balbus, Hawley, Gammie 1996

11. 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):

12. 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?!

13. Disk accretion can be highly time-variable, with short bursts of very rapid accretion.

14. FU Ori; outburst of disk accretion Disk accretion  10 -7 - 10 -8 M  /yr  protostar;  Disk accretion  10 -4 M  /yr  FU Ori object

15. if dM/dt (infall) > dM/dt (accretion): onto disk onto star mass buildup  eventual rapid disk accretion Why unsteady accretion? Infall to disk; high velocity disk accretion; low radial velocity  no reason to balance!

16. 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

17. 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??

18. Magnetic fields CAN couple to protostellar disks: Jets/Winds 280 AU Burrows et al. 1996 Flared 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!

19. bead on a wire analogy collimation Alfven surface

21. 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)

22. FU Ori disk winds disk rotation Hartmann & Calvet (1995); accelerating disk wind results in shifts increasing with increasing strength (upper levels) Petrov & Herbig 1992

23. 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)

24. IMTTS: predecessors of the HAeBe T Tauri stars: CTTS= accreting WTTS=not acc. HAe/Be

25. T Tauri: (FGKM) pre-main sequence stars with disks Hartmann 1998

26. T Tauri star spots (cool); BIG! (large stellar B) Stelzer et al. 2003 V410 Tau (stellar luminosity perturbed? Rosner & Hartmann… - observational problems

27. Proxies for magnetic fields (activity): enhanced in pre-main sequence stars - “saturated” behavior (i.e. not strongly rotation-dependent) Chromospheric fluxesX-ray fluxes Walter et al. 1988 (accretion) note: x-ray emission not affected (much) by disk accretion (“T”)

28. Flaccomio et al. 2003 Orion Nebula cluster stars (ages ~ 1 Myr) “Saturation” : B or heating efficiency?

29. 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

30. 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

31. 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)

33. Magnetospheric accretion: line profiles (Muzerolle et al. 1998): line width  (2GM * /R * ) 1/2 Königl 1991

34. Models for magnetospheric emission

35. Circularly polarized He I emission Johns-Krull et al. 1999 LCP RCP

36. 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

37. Magnetospheric accretion and outflow Numerical simulations show complex accretion pattern, not always polar, even when pure aligned dipole (Miller & Stone 1997)

38. 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

39. 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

40. 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”) (??)

41. Why do young stars rotate so slowly if they are formed from disk accretion? And why faster for lower-mass stars?? Clarke & Bouvier 2000

42. 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

43. 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

44. 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 co- rotation, no winding problem: unrealistic (axisymmetric, etc.) detailed assumptions not very clear (Collier Cameron & Campbell)

45. The angular momentum problem Shu et al. “funnel” flow + x-wind

46. Lovelace, Romanova, & Bisnovatyi-Kogan 1995

47. 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)

48. von Rekowski & Brandenburg 2004; also Goodson, Winglee, Matt Disk dynamo? Opposed field to star? Accretion, spindown oscillatory

49. 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) = k 2 M *  * R * 2 dM/dt  d R d 2  k 2 (M * /dM/dt) (  * /  K ) (R * /R co ) 1/2  0.2  10 8 yr  (  * /  K ) / 2 so either slow rotation or need very high dM/dt to spin down in 10 6 yr

50. Disk-star magnetic coupling: does it work? or?? coronal mass ejection-type loss, except using disk material??

52. 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 ~ 10 7 years! Disks? But disks don’t seem to last quite that long! Bouvier et al. 1997

53. 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 P B ~ P g (photospheric) everywhere?? Why is activity “saturated”?