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					Clustered Massive Star Formation in Molecular Clouds
Jonathan C. Tan

ETH Zurich &
University of Florida
  Some open questions
What are the initial conditions for
proto star clusters and how do they arise?


What are the initial conditions for individual
massive star formation within star clusters?

How do massive protostars accumulate their
mass?
          Pressures in Star-Forming Regions
  Pressure due to self-gravity:                                             R
                                                                       M
  surface density


  pressure
  McKee & Tan (2002; 2003)




Local examples in the Galaxy:  ~ 0.1 - 1.0 g cm-2




  Infrared dark clouds (IRDCs)                   Star-forming clumps       Embedded clusters
Surface density vs. cluster mass
     Surface density vs. cluster mass

Morphologies Orion: tcluster
Age spread of of clumps: >~3Myr
Round and centrally concentrated
-> star formation occurs over
(Shirley et al. 2003)
many crossing times
                                          Galactic clumps: Mueller et al. (2002)
                                          IRDCs:       Carey et al. (2000);
                         =7000 M pc-3                Kirkland & Tan, in prep.
                                                 AV=200
                                                 NH=4.2x1023cm-2
                                                 =4800 M pc-2


                                                AV=7.5
                                                NH=1.6x1022cm-2
                                                =180 M pc-2
        CO clouds                               AV=1.4
                                                NH=3.0x1021cm-2
                                                =34 M pc-2
    Surface density vs. cluster mass

Westerlund 1
(Clark et al. 2005)


                             Galactic clumps: Mueller et al. (2002)
                             IRDCs:       Carey et al. (2000);
                                          Kirkland & Tan, in prep.
                                    AV=200
                             Galactic center clusters:
                             e.g. Figer et al. (1999);
                             Kim et al. (2000)

                             Super star clusters:
                             e.g. Gilbert & Graham (2001)
                             De Marchi et al. (1997)
                             Turner et al. (2000)
     Pressures in Star-Forming Regions
Ram pressure due to converging flows:

                                          v

pressure                                      



Local Galactic GMC-GMC collisions
(e.g. Scoville et al. 1986; Tan 2000)
         nH~500cm-3                     v~10km/s
Galaxy-galaxy collisions
      nH~1cm-3                          v~200km/s
Converging gas flows more efficient than shocks
from stellar feedback (Elmegreen 2003)
Surface density vs. cluster mass



                         Galactic clumps: Mueller et al. (2002)
                         IRDCs:       Carey et al. (2000);
                                      Kirkland & Tan, in prep.
                                AV=200
                         Galactic center clusters:
                         e.g. Figer et al. (1999);
                         Kim et al. (2000)

                         Super star clusters:
                         e.g. Gilbert & Graham (2001)
                         De Marchi et al. (1997)
                         Turner et al. (2000)

                         Ram pressures:
                         Nuclear GMC-GMC : 7.3x109 K cm-3
                         GMC-GMC : 1.3x107 K cm-3
                         galaxy-galaxy : 6.9x106 K cm-3
 What are the initial conditions for individual massive star formation?
        How do massive protostars accumulate their mass?




                                                          ?
Does a quasi-equilibrium core form first, and then collapse to a star?
Or does a protostellar seed grow via competitive Bondi-Hoyle accretion
and/or stellar collisions? (Bonnell, Dobbs, Dale, et al.; Klessen et al.)
What are the initial conditions for individual massive star formation?


                                                McKee & Tan (2002; 2003)




   Consider 60M pre-stellar core
   under high Pgrav~ G2 ~ 109 K cm-3
   at center of clump (=proto cluster)
   Surrounded by lower mass cores, perhaps
   with mass function similar to stellar IMF
   (Alves; Beuther & Schilke; Motte et al.; Testi & Sargent)
What are the initial conditions for individual massive star formation?
        Core surrounded by pressure of clump
                                                           m* ~ 0.5 Mcore




   Core


                                         Final mass accretion rate




                                          nH,s <<tcluster core/60M)1/23/2cm-3
                                            t*f = 106(M
  What are the initial conditions for individual massive star formation?
Turbulent cores, fragmenting from a turbulent medium,
reasonably close to virial, hydrostatic equilibrium




                                           Final mass accretion rate




                                            Support by combination of
                                            large & small scale B-fields,
                                            and turbulent motions.
                                            Core boundaries fluctuate.
             Evidence in favor of Turbulent Fragmentation
Observational:

~Quiescent cores are seen,
both with and without stars
(Walsh, Myers, & Burton 2004).


Mass function of cores
appears similar to stellar IMF
(Beuther & Schilke 2004; Motte et al. 2001).




 Collapsing cores?
 e.g. talk by Yuefang Wu

 Disks?
 e.g. Beltran et al. 2004;
 Sandell et al. 2003; Chini et al. 2004

 Collimated Outflows
 e.g. Beuther et al. 2002
        Evidence in favor of Turbulent Fragmentation


Theoretical:

By including B-fields, relatively long-lived cores are seen to form in
simulations, even with driven turbulence (Vasquez-Semadeni et al. 2004).

SPH sink particle technique,
the basis of competitive accretion models,
does not adequately resolve Bondi-Hoyle                             QuickTime™ an d a
                                                                TIFF (LZW) decomp resso r
                                                             are need ed to see this picture.

accretion (Krumholz, McKee, Klein 2005).
Vázquez-Semadeni, Kim, Shadmehri & Ballesteros-Paredes (2004)




                             QuickTime™ an d a
                        YUV420 codec decompressor
                       are need ed to see this p icture .
         Comparison to observed regions

     W49
     d=11.4kpc




                                      3.6cm - De Pree et al. (1997)


Core moves during its collapse: 5 km/s for 105yr is 0.5pc
   Comparison to observed regions

W3(OH), W3(H2O)
d=2.2kpc
                                           QuickTime™ an d a
                                       TIFF (LZW) decomp ressor
                                    are need ed to see this p icture .
1.3mm - Wyrowski et al. (1999)



                                                  QuickTime™ and a
                                              TIFF (LZW) decompressor
                                           are neede d to see this picture.




     3.6cm - Wilner et al. (1999)
    Comparison to observed regions
Orion Nebula
Cluster
d=450pc




               Near IR: VLT-ANTU+ISAAC (ESO)
                       What ejected BN?
                           BN: 2500-104L = B3-B4 -> mBN = 8-12M
                           vBN-vI~40 km/s (Plambeck et al. 1995; Tan 2004;
                           Bally & Zinnecker 2005; Rodriguez et al. 2005)
                           vr,BN = +21km/s; vONC = +8km/s
                           Look for massive star along past trajectory:
                           Source I, Trapezium stars, 2A
                           Of revealed sources, only 1C has the right
                           direction of motion (van Altena et al. 1988).
                           Amplitude of motion of 1C is 4.9±0.5km/s
Hut & Bahcall (1983)       -> mBN= 6.4±3M
                           1C should likely be a hard, eccentric binary,
                           with relatively massive secondary…
                           … it is! 45M primary and >~6M secondary,
                           Separated by ~17AU (->vesc=70 km/s)
                           (Schertl, Balega, Preibisch, Weigelt 2003).

                             BN ejected from 1C about 4000 years ago.
                             Both systems are leaving the cluster.
      Testing the model in Orion




              Near IR (VLT: ANTU+ISAAC)

Mcore ~60M, m*~20M, L*~8x104L          Tan 2003
               Protostellar Evolution

Bolometric
Luminosity


Protostellar
Radius


Ionizing
Luminosity
                         Outflows
              Density distributions of hydromagnetic outflows
              (e.g. disk wind, X-wind) approach a common form
              far from the star or inner disk: collimated wind
              (Shu et al. 1995; Ostriker 1997; Matzner & McKee 1999).
HH30 - HST

                                                 A massive hot protostar
                                                 may ionize the inner part
                                                 of the outflow.
             outflow      HII
                         Region


                                           disk
                                         Outflow-Confined HII Regions
                                            Treat sectors independently        Tan & McKee (2003)
                                                         ri
                                            S = 4   ∫rc
                                                              (2) nw2 r2 dr
                                                                                astro-ph/0309139


                                                                               Fiducial Model:
Radio Spectrum: thermal bremsstraulung




                                                                               X-wind
                                                                               m* = 20M
                                                                               r* = 16R
                                                                                .
                                                                               m* = 3x10-4M/yr
                                                                                     . .
                                                                               fw = mw/m* = 0.1
                                                                               fv = vw/vK* = 2.1

                                                     Radio source “I” appears
                                                     elongated (0.145”=65AU 43GHz continuum
                                                                              Menten & Reid
                                                     by <0.085”) at 22GHz
                                                     (K. Menten, priv. comm)   SiO (v=1,2; J=1-0)
                                                                               Greenhill et al. (2003)
              Relation to Star Clusters (1)
Outflows: summed from many protostars;
          attempt to measure SFR of cluster




Tan & McKee (2002)
proceedings of Boulder meeting
            Relation to Star Clusters (2)
Feedback: ionization, stellar winds, radiation pressure
          acting on a clumpy medium




Tan & McKee (2004)
proceedings of Cancun cluster
meeting
                               Conclusions
What are the initial conditions for proto star clusters and how do they arise?
  Infrared Dark Clouds: cold, dense, turbulent, self-
  gravitating, pressurized. Perhaps forming in compressive
  flows, i.e. cloud collisions.

What are the initial conditions for individual massive
star formation within star clusters?
 Turbulent cores, fragmenting from a turbulent medium,
 reasonably close to virial, hydrostatic equilibrium.

How do massive protostars accumulate their mass?
  Quantitative calculations of collapse of core to star
  via disk accretion: protostellar evolution, outflow,
  ionization -> outflow-confined HII region.
  Applied to source I in Orion KL.




                                     Near IR: VLT-ANTU+ISAAC (ESO)

				
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posted:7/28/2011
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