Masers and Massive Star Formation Claire Chandler Overview: –Some fundamental questions in massive star formation –Clues from masers –Review of three regions: W3, Cep A, Orion –Preview of a movie of SiO masers associated with Source I in Orion –What have we learned?

Current picture of low-mass star formation

The problem with extending the picture of low-mass star formation to massive stars is the following: Radiation pressure acting on dust grains can become large enough to reverse the infall of matter: –F grav = GM * m/r 2 –F rad = L  /4  r 2 c –Same dependence on r  happens at all radii Luminosity prior to onset of nuclear burning comes from –Accretion, L acc = GM * M acc /R * –Gravitational contraction, L int Transition from where the evolution is dominated by the accretion timescale (  1/L acc ) to the Kelvin-Helmholtz timescale (  1/L int ) is ~10 M  

So, how do stars with M * >10M  form? Accretion: –Need to reduce  e.g., by making accreting material very optically thick (high M acc ) –Reduce the effective luminosity by making the radiation field anisotropic Form massive stars through collisions of intermediate-mass stars in clusters –May be explained by observed cluster dynamics –Possible problem with cross section for coalescence –Observational consequences of such collisions? 

Other differences between low- and high-mass star formation Physical properties of clouds undergoing low- and high- mass star formation are different: –Massive SF: clouds are warmer, larger, more massive, mainly located in spiral arms; high mass stars form in clusters and associations –Low-mass SF: form in a cooler population of clouds throughout the Galactic disk, as well as GMCs Energetic phenomena associated with massive SF: UCH II regions, hot molecular cores Different environments observed has led to the suggestion that different mechanisms (or modes) apply to low- and high-mass SF

Clues from high-resolution maser observations Maser proper motions plus radial velocities give 3-D velocity fields Masers trace a variety of physical conditions, depending on the molecule and pump mechanism: –OH (1665/7 MHz): n ~ 10 6  7 cm  3, T ~ 100 K –CH 3 OH: n ~ 10 6  7 cm  3, T ~ 100 K –SiO (v=0): n ~ 10 6 cm  3, T ~ few 100 K (very rare) –H 2 O: n ~  12 cm  3, T ~ few 100 K –SiO (v=1, v=2): n ~  12 cm  3, T ~ 1000  2500 K (rare) –Other OH transitions, HCN, NH 3, HCO 2 H, etc… Zeeman effect for OH, H 2 O gives B-field

Case studies: W3, Cep A, Orion W3 contains two main sites of massive star formation at ~2kpc: W3(Main), W3(OH): W3(Main) from Tieftrunk et al. (1997) W3(OH) from Reid et al. (1995)

Argon, Reid & Menten (2003) Moscadelli et al. (1999) Bloemhof et al. (1992) H 2 O maser proper motions  outflow from TW object CH 3 OH masers roughly coincident with OH masers OH maser proper motions  expansion at ~ few km/s B-field from OH Zeeman  ~10 mG (Baudry & Diamond 1998) W3(OH):

W3(Main): H 2 O maser proper motions trace several outflows from the IRS5 region Zeeman effect in H 2 O masers close to “c” give B ~ 15  40 mG (Sarma et al. 2001) Tieftrunk et al. (1997) Imai et al. (2000)

Cepheus A Most dense molecular core in the molecular cloud complex associated with Cep OB3 association, d~725pc, L bol ~2.5  10 4 L 

VLA observations originally interpreted as a disk around HW2 H 2 O masers in the vicinity of Cep A HW2 HW2 Torrelles et al. (2001) Garay et al. (1996)

H 2 O maser proper motions of R1  3 VLBA and MERLIN observations identify multiple sources for the masers: –R1, R2, R3 shocks outlining walls of outflow cavity Torrelles et al. (2001)

H 2 O maser proper motions of R4 –R4 possibly a disk around a ~3 M  star Gallimore et al. (2003) Torrelles et al. (2001)

H 2 O maser proper motions of R5 –R5 edge of an expanding bubble caused by spherical mass ejection from an embedded protostar Torrelles et al. (2001) Curiel et al. (2002)

Orion BN/KL Shocked H 2 emission traces an explosive outflow event centred close to radio Sources “I” and “n” L bol ~ 5-8  10 4 L  Menten & Reid (1995) Schultz et al. (1999)

OH, H 2 O and SiO masers in the vicinity of Source I Johnston et al. (1989) H 2 O masers: Gaume et al. (1998), Greenhill et al. (1998) H 2 O masers trace a ~20km/s flow, and SiO v=1 masers trace an “X” centred on Source I SiO Greenhill et al. (1998); Doeleman et al. (1999) OH H2OH2OH2OH2O Model

Monthly monitoring of SiO masers in Source I with the VLBA Greenhill, Chandler, Reid, Moran, Diamond The velocity field traced by the SiO masers close to the protostar can potentially determine whether the MHD disk wind models currently in vogue for outflows from low- mass protostars will also work for massive protostars. Monthly monitoring of the v=1 and v=2, J=1  0 SiO masers ( ~43GHz) with the VLBA began in June 2000 and is continuing through this summer Data sets are large: , 512 channels/transition; image ~25% of each cube  ~60GB/epoch just for the images Sneak preview of results from 4 epochs…

Single epoch SiO radial velocities and VLA 7mm continuum North Arm East Arm South Arm West Arm Bridge

Source I: the movie preview

Models for the Source I disk/outflow system

Summary: what does it all mean? VLBA maser proper motion studies provide the highest resolution possible of the dynamics of star formation Maser spot geometry and kinematics resemble those of low- mass systems, suggesting formation via accretion: such ordered motions are unlikely to result from coalescence W3 –Masses of outflow sources unknown; probably less than 10M  Cep A –Mass of HW2 probably ~10M  ; other sources less massive Source I –If edge-on disk model is correct, rotation  M * ~10  15M  ; this may be the first demonstration that accretion models can be scaled to high-mass systems Future: more proper motion studies needed for M * >10M  ; B- field measurements needed to constrain MHD wind models