Young Stars I,II

1. 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

2. Stars form from the collapse of protostellar gas clouds, r  104 AU optical infrared Alves, Lada & Lada 2001

3. (GM/R2) M coef.  (d/dR) (B2/8π) (4πR3/3) Essentially ALL protostellar cloud magnetic flux must be lost during star formation (protostars don’t have such B) (Mestel & Spitzer) The magnetic flux “problem” For gravity to overcome magnetic pressure: 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! Why? low ionization at high  as collapse proceeds, so flux-freezing is not a good approximation (Umebayashi & Nakano 1988)

4. Therefore, even if (o/K)2 ~ 0.1, R(final) ~ 0.01 R ~ 1015 cm ~ 100 AU. 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 Stars must form from disk accretion (magnetic flux loss in low-ionization disks)

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

6. Magnetorotational instability (MRI)? Works well when ionization high enough (?) Why do disks accrete? Hydrodynamic exchange? Doesn’t seem to work Gravitational instability? May work; requires massive disk

7. Side view: initial vertical field Magnetorotational Instability? (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

8. Stone, Balbus, Hawley, Gammie 1996

11. X-ray or CR ionization Thermal ionization (T > 1000K) Dead zone (and layered accretion) (Gammie 1996) BUT: low ionization  no magnetic viscosity  no accretion! T Tauri disk (model): 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?

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 - 10-8M/yr  protostar;  Disk accretion  10-4M/yr  FU Ori object

15. Why unsteady accretion? if dM/dt (infall) > dM/dt (accretion): onto disk onto star mass buildup  eventual rapid disk 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. 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

18. Jet seen in [O I] (accretion-driven) 280 AU Flared disk seen in scattered light: dust lane obscures central star Burrows et al. 1996 Magnetic fields CAN couple to protostellar disks: Jets/Winds • Thermal pressure too low to accelerate flows • Radiation pressure negligible • Collimation!

19. Alfven surface bead on a wire analogy collimation

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 FU Ori winds must be heated to explain H, etc; numerical simulations of MRI show waves propagating upward and shocking Miller & Stone 2002 “Atmospheric” absorption line profiles show evidence for sonic “turbulence” (Hartmann, Hinkle & Calvet 04)

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

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

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

27. note: x-ray emission not affected (much) by disk accretion (“T”) 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

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

29. 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

30. 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”

31. “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) Why magnetospheric accretion?

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

34. Models for magnetospheric emission

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

36. Classical TTS Accretion power in T Tauri Stars 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

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? accreting non-accreting Taurus: accreting stars (stars with disks) rotate more slowly (Bouvier et al., Edwards et al. 1993)

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” (Collier Cameron & Campbell) Solution 2: exact co-rotation, no winding problem: unrealistic (axisymmetric, etc.) detailed assumptions not very clear

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

49. 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 Disk-star magnetic coupling: does it work? To spin down star, either wind or disk must carry away the stellar J! 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

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