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Unformatted text preview: High-mass star formation IR Dark Cloud Ext. Map G28.37 (Spitzer/GLIMPSE) (Butler & Tan) Michael Butler, Audra Hernandez, Bo Ma, Yichen Zhang Sven Van Loo, Peter Barnes, Elizabeth Lada Charlie Telesco Friday, November 19, 2010 Orion Nebula Cluster (VLT; JHK) (McCaughrean) Jonathan Tan (University of Florida) Paola Caselli (Leeds), Francesco Fontani (IRAM), Izaskun Jimenez-Serra (CfA), Mark Krumholz (UCSC), Christopher McKee (UCB), Francesco Palla (Arcetri), Jan Staff (LSU), Leonardo Testi (ESO), Barbara Whitney (SSI) 1 Why work on massive star formation? The First (Pop III) Stars were likely massive, some potentially supermassive stars, reionizing the universe and producing the first metals. Galaxies form and evolve by forming star clusters, where the influence of massive stars is paramount. Massive stars are what tend to be seen in distant galaxies. Planets form from the crumbs left over from star formation. Planet & star formation in star clusters can be influenced by massive star feedback. Supermassive black hole formation may be via massive star clusters or Pop III stars. Supermassive black hole accretion is likely to be regulated by star formation. Friday, November 19, 2010 2 Why not to work on massive star formation... A complicated, nonlinear process Numerical models Wide range of scales (~12 dex in space, time) and multidimensional. Uncertain/unconstrained initial conditions/boundary conditions. Observations Complete theory of star formation Analytic theory Physics: Gravity vs pressure (thermal, magnetic, turbulence, radiation, cosmic rays) and shear. Heating and cooling, generation and decay of turbulence, generation (dynamo) and diffusion of B-fields, etc. Chemical evolution of dust and gas. Some notation: Core -> star or close binary Clump -> star cluster Friday, November 19, 2010 3 Outline • Physical properties of massive star-forming regions • Theoretical scenarios - core accretion, competitive accretion, mergers, etc. • The “Turbulent Core Accretion” Model • Initial conditions: IRDCs; how are they generated? Does the clump reach pressure equilibrium? Timescale of star cluster formation? Collapse of the core: fragmentation? • Massive protostars: star, disk, outflow formation and evolution. Radiative transfer modeling. • [Feedback: outflows, ionization, rad. pressure. On core & clump.] Friday, November 19, 2010 4 Overview of Physical Scales Friday, November 19, 2010 5 Overview of Physical Scales AV=7.5 A8μm=0.30 NH=1.6x1022cm-2 Σ=180 M pc-2 AV=1.4 NH=3.0x1021cm-2 Σ=34 M pc-2 Friday, November 19, 2010 6 Overview of Physical Scales 02) (20 n e s, le uel M van y, E e hirl r, S obs Jac AV=200 A8μm=8.1 NH=4.2x1023cm-2 Σ=4800 M pc-2 AV=7.5 A8μm=0.30 NH=1.6x1022cm-2 Σ=180 M pc-2 AV=1.4 NH=3.0x1021cm-2 Σ=34 M pc-2 Friday, November 19, 2010 7 Overview of Physical Scales AV=200 A8μm=8.1 NH=4.2x1023cm-2 Σ=4800 M pc-2 AV=7.5 A8μm=0.30 NH=1.6x1022cm-2 Σ=180 M pc-2 AV=1.4 NH=3.0x1021cm-2 Σ=34 M pc-2 Friday, November 19, 2010 8 M (M 82 cC SS ra Cs dy & Gr a ha m 20 07 ) Overview of Physical Scales AV=200 A8μm=8.1 NH=4.2x1023cm-2 Σ=4800 M pc-2 AV=7.5 A8μm=0.30 NH=1.6x1022cm-2 Σ=180 M pc-2 AV=1.4 NH=3.0x1021cm-2 Σ=34 M pc-2 Friday, November 19, 2010 9 .) t al re These are the environments where massive stars form: can we scale-up low-mass SF theory? ne rs J. Tur la egu icky, irr arf obuln w in d on, K s Cs SS John (K. nH~2x105cm-3 tff~1x105yr Overview of Physical Scales AV=200 A8μm=8.1 NH=4.2x1023cm-2 Σ=4800 M pc-2 AV=7.5 A8μm=0.30 NH=1.6x1022cm-2 Σ=180 M pc-2 AV=1.4 NH=3.0x1021cm-2 Σ=34 M pc-2 Friday, November 19, 2010 10 Schematic Differences Between Massive Star Formation Theories pre-massive-stellar core Rare evolution from magnetically subcritical state? Beuther, Churchwell, McKee, Tan (2007); Tan (2008) massive-star-forming core [protostar+gravitationally-bound gas] LIMP-MP massive-protostar (MP) core fragmentation disk fragmentation time massive star m*f>8M Is there any isolated massive star formation? Turbulent core model (McKee & Tan 2002, 2003) Competitive Bondi-Hoyle accretion model (Bonnell ea. 2001; Bonnell & Bate 2006; Dobbs+, R. Smith+, P. Clark+) Prestellar core mass function? Friday, November 19, 2010 t=0 protostar formation m*=8M Radiation pressure likely to prevent accretion of dusty, unbound gas (e.g. Edgar & Clarke 2004) 11 radiation can exert enough outward pressure to halt infall, inhibiting further stellar growth (1). The presence of a flattened accretion disk surrounding the protostar (2) can alleviate this in- 008 involve accretion through a flattened disk and molecular outflows. The magnetic field is thought to play an important role in the formation of Sunlike stars by shaping cloud collapse, removing ex- North Aohoku Place, Hilo, HI 9672 *To whom correspondence shou girart@ieec.cat †Present address: Osservatorio A Enrico Fermi 5, 50125 Firenze, I Strong magnetic fields in star-forming regions MAGNETICALLY ALIGNED VELOCITY ANISOTROPY ! 425 A B Fig. 1. (A) Contour map of the 879-mm dust emission superposed on the color ! image of the polarized flux intensity in units of Jy per beam. Black thick bars indicate the position s of the ght) (Left) Image of 12CO J ¼ 1 0 emission of a subfield within the Taurus molecular cloud integrated over the velocity interval 5.5–7.5 kmangleand (rimagnetic field. These maps were obtained by using a natural weighting line visibility O velocity centroid ( Narayanan et al. 2008), with overlay of optical polarization vectors from the compilation by Heiles (2000). The molecularto theemis- data, which yielded to a full width at half maximum synthesized method is ocities exhibit streaks that are aligned along the local magnetic field direction. The solid line box outlines the area on which the axis-constrained PCAbeam of 1.34″ × 0.83″ with a position angle of 67° (shown in e dotted-line box shows the area within which the polarization angles are averaged to estimate the mean magnetic field direction. the bottom left corner). Contour levels are 0.8, 1.5, 2.5, 4, 6, 16, 26, 36…96% of the peak intensity, 9.13 Jy per beam. (B) Contour map of the 879 mm dust emission superposed on the color image of the flux weighted Correlation of field orientations from ~100pc to <1pc scales (Hua-bai Li et al. 2009) and the mean position angle of the emission streaks of CO Taurus (Heyer et al. 2008) onstrained PCA eigenfunctions to show a clear signalocity anisotropy induced by MHD turbulence. 4. THE TAURUS MOLECULAR CLOUD urus molecular cloud provides a valuable platform to ininterstellar gas dynamics and the star formation process, its proximity (140 pc) and the wealth of complementary yanan et al. (2008) present new wide-field imaging obs of 12CO and 13CO J ¼ 1 0 emission from the ceneg2 of the Taurus cloud complex, obtained with the 4 m telescope. The images identify a low column denrate of gas that contains subtle streaks of elevated 12CO aligned along the local magnetic field direction as defrom stellar polarization measurements (Heiles 2000). f 12CO J ¼ 1 0 integrated intensity and centroid veh measured polarization vectors from this subfield are Figure 3. These show a connection between the density ity fields. While the origin of these streaks is unknown, rous alignment with the polarization vectors strongly that the interstellar magnetic field plays a prominent e gas dynamics of this low-density material. ess the degree of velocity anisotropy within this subthe Taurus molecular cloud, we have applied the axised PCA method to the 12CO data from this imaging he precise field is described by the solid box in Figure 3. Friday, November 19, 2010 "#$%&'! ()! ! *+$,'-#.! /#'012! #,! -3'! 4&#5,! 650'.%0+&! .05%1! &'$#5,"! #$%! &'()*+,-./! 01'*%!2$,32!4$%!5678!9:%-*%&'-%+!%4!';"!<=>?@!<AA! !1!1'B!0.!;,*'+04$10(!2(';%"! ! C%! À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velocity map of the CH3OH 147-156 A. Black thick bars of the magnetic field. These maps were obtained by usi of 0 to the visibility data, which yielded to a full wi synthesized beam of 1.04″ × 0.59″ with a position angl bottom left corner). Contour levels are the same as in th a peak intensity of 6.55 Jy per beam. (C) Spectrum of th position of the dust emission peak. The continuum has the line emission (this is valid for all the molecular line Girart et al. (2009) 12 emission. The x- and y-axis structure functions derived at MAX ¼ 1408 46 are shown in Figure 5. These distributions show the same pattern of offsets between the parallel and perpendicular structure functions measured in the strong-field simulation snapshots ( B2, B3) shown in Figure 2. For the Taurus field, the power-law index of the structure function derived from 12CO along the x-axis (i.e., the direction aligned with the polarization) is steeper (0:81 Æ 0:05) than the index of the y-axis structure function (0:34 Æ 0:06). The steeper power law along the x-axis is indicative of a velocity field more dominated by large scales. Similar to the model structure functions in the strong magnetic field cases, the normalization of the y-axis structure function, v0; y , is 0.08 km sÀ1 and larger than the value of the x-axis structure function (v0; x ¼ 0:02 km sÀ1). Thus, the smooth variation of density along the presumed magnetic field is mirrored by a smooth variation in the velocity, and the stronger variation in density in the perpendicular direction (streakiness) is mirrored by a stronger variation in the velocity. Indeed, preliminary analysis shows that in the direction perpendicular to the projected magnetic field, displacements between the peaks in integrated intensity and velocity centroids are similar with typical values 0.2–0.4 pc. The results shown in Figures 3, 4, and 5 are suggestive of velocity anisotropy induced by strong MHD turbulence, as described by GS95 and verified by computational simulations (Cho ! Supercritical 12 JUNE 2009 VOL 324 SCIENCE www.sciencemag.org Σ = 1 g cm-2 Strength of B-field vs. Σ (Crutcher 2005; Falgarone et al. 2008) 13 Subcritical 12 From Cores to Stars: RATHBORNE ET AL. Individual Stars Appear to Form from Cores Beuther & Schilke (2004) Nutter & Ward-Thompson (2007) Rathborne et al. (2009) erived CMFs shown as binned histograms. The four panels correspond to different core samples when taking into consideration the C18 O n. Note that the shape of the mass function changes considerably when we use the C18 O (1–0) emission to guide the core extraction fro core * core between panels (a) and (d)). The dashed line corresponds to the scaled field star IMF of Muench et al. (2002). For each panel we dete n, the offsets between the CMF and the IMF. To accurately determine these scaling factors, we minimize the χ 2 between the distributions mpletenesse.g. Motte et al. 1998; Testi & Sargent 1998; Motte et al..2001vertical dottedWilson 2005; Alves et al. 2007; Li et al. (Kainulai See also: limit. The derived parameters are summarized in Table 1 The ; Mike Reid & line marks the mass completeness limit 2007; the histograms are the errors for each bin (calculated2009; André etroot2010 number per bin). Enoch et al. 2008; Pineda et al. 2009; Ragan et al. as the square al. of the ε = m /m Friday, November 19, 2010 the significant change in -> ~0.06 shape of the CMF occurs 0.22+-0.08 13 a mass of 2.7 ± 1.3 M . Using the derived scalin Overview of Physical Scales Turbulent Core Model of Individual Massive Star Formation (McKee & Tan 2003) nH~2x105cm-3 tff~1x105yr AV=200 A8μm=8.1 NH=4.2x1023cm-2 Σ=4800 M pc-2 AV=7.5 A8μm=0.30 NH=1.6x1022cm-2 Σ=180 M pc-2 AV=1.4 NH=3.0x1021cm-2 Σ=34 M pc-2 Friday, November 19, 2010 14 How many Massive Starless Cores in the Galaxy? Number in the Galaxy: (see also Zinnecker & Yorke 2007) If lifetime of this phase is ~ t*f ~ 105yr, then for a Galactic SFR of 3Myr-1 and an IMF yielding 1 massive star per 130 M (Salpeter 0.1-120M) and 2/3 of massive stars are forming in binaries, we expect 1500 Massive Starless Cores in the Galaxy. Friday, November 19, 2010 15 Mid-IR Extinction Mapping of Infrared Dark Clouds (Butler & Tan 2009; see also Peretto & Fuller 2009; Ragan et al. 2009; Battersby et al. 2010) G28.37+00.07 Spitzer - IRAC 8µm (GLIMPSE) Median filter for background around IRDC; interpolate for region behind the IRDC Correct for foreground emission - tricky-> choose nearby clouds 16’ Extinction map to derive Σ Distance from molecular line velocities (GRS) -> M(Σ) g cm-2 MJy sr-1 Friday, November 19, 2010 16 Application to Filamentary IRDCs G035.39−00.33 Comparison to mm dust emission (Rathborne et al. 2006) and 13CO and C18O line emission (Hernandez & Tan, submitted), give agreements at ~factor of 2 level 3’ I8μm (MJy sr-1) Friday, November 19, 2010 Σ (g cm-2) 17 Formation of IRDCs Some evidence that filamentary IRDCs are not yet virialized: Comparing to models of Fiege & Pudritz (2000) Filamentary virial analysis of 2 IRDCs (Hernandez & Tan, submitted) Extended SiO emission along one IRDC (Jimenez-Serra et al. 2010) But, the regions closer to virial equilibrium do appear to be those forming stars Friday, November 19, 2010 18 MIPS 24μm IRAC 8μm Extinction Map Massive Starless Cores Butler & Tan (2009), Butler & Tan, in prep. Σ = 0.26 g cm-2 mcore = 205 M 10” Σ = 0.12 g cm-2 mcore = 94 M Σ = 0.12 g cm-2 mcore = 50 M 10” Cores show central concentration; can fit power law radial density profiles, index ~-1.5. They contain many thermal Jeans masses. B-fields may be suppressing fragmentation within the core. 10” nH~105cm-3, B~1mG -> MB~100 M Friday, November 19, 2010 19 Overview of Physical Scales nH~2x105cm-3 tff~1x105yr Butler & Tan 2009 AV=200 A8μm=8.1 NH=4.2x1023cm-2 Σ=4800 M pc-2 AV=7.5 A8μm=0.30 NH=1.6x1022cm-2 Σ=180 M pc-2 AV=1.4 NH=3.0x1021cm-2 Σ=34 M pc-2 Friday, November 19, 2010 20 We expect massive star forming environments exist for >1tff and so can achieve approx. pressure equilibrium (proto star clusters take > 1tff(central) to form) Tan, Krumholz, McKee (2006) IRDC cores have tff~105yr, which is short Some (most?) star clusters appear to have age spreads >106yr, e.g. Orion Nebula Cluster median age of 2.5-3Myr (Da Rio et al. 2010) A plausible mechanism has been identified to maintain turbulence over many tff: protostellar outflow feedback (Norman & Silk 1980; Nakamura & Li 2007) While the issue of star cluster formation timescales is still debated (e.g. Elmegreen 2000, 2007; Hartmann & Burkert 2007), it seems likely that tform>tff(central). Friday, November 19, 2010 21 Collapse of the Core - Core Fragmentation? We expect most of these structures will fragment to form star clusters. Most mass -> low-mass stars. Fragmentation will be reduced by radiative feedback from the central star (Krumholz, Klein & McKee 2007; c.f. Dobbs, Bonnell, Clark 2005). Fragmentation should be reduced by radiative feedback from surrounding accreting low-mass stars (Krumholz & McKee 2008). Magnetic field support should increase the “magneto-Jeans” mass and reduce fragmentation: (Machida et al. 2005; Price & Bate 2007, Hennebelle & Teyssier 2008, Duffin & Pudritz 2009). However, see Li, Wang, Abel, Nakamura (2010). Friday, November 19, 2010 22 Overview of Physical Scales Fragmentation stopped by radiative heating (Krumholz & McKee 2008) But B-fields likely to also suppress fragmentation nH~2x105cm-3 tff~1x105yr Butler & Tan 2009 AV=200 A8μm=8.1 NH=4.2x1023cm-2 Σ=4800 M pc-2 AV=7.5 A8μm=0.30 NH=1.6x1022cm-2 Σ=180 M pc-2 AV=1.4 NH=3.0x1021cm-2 Σ=34 M pc-2 Friday, November 19, 2010 23 The later stages of individual massive star formation core bounded by pressure of clump Core Final mass accretion rate Friday, November 19, 2010 24 Outflow-confined HII Region The later stages of individual massive star formation Final mass accretion rate Protostellar evolution r* m* Friday, November 19, 2010 Disk structure Outflows Support by combination of large & small scale B-fields, and turbulent motions. Core boundaries fluctuate. 25 Collapse from Core to Disk Beltrán et al. (2004) Observational evidence for rotating toroids on scales ~1000AU, perpendicular to bipolar outflows, e.g. G24.78+0.08 A1 A2 A1 Also claims from maser observations (e.g. Wright et al., Greenhill et al., Goddi et al. in Orion KL; Pestalozzi et al. 2004 in NGC7538 IRS1N) Theory: Analytic study of disk accretion and fragmentation (Kratter, Matzner, Krumholz 2007) Radiation-hydro simulation of turbulent core collapse: modest disk fragmentation. (Krumholz, Klein, McKee 2007a). Simulated ALMA observations (Krumholz, Klein, McKee 2007b). Friday, November 19, 2010 CH3CN (12,0,0,13)->(11,0,0,12) [220GHz, 69K] 26 Protostellar Evolution [see also Hosokawa & Omukai 2009] Tan & McKee 2002 Radius Luminosity protostar+boundary layer+disk Ionizing L. Outflow momentum flux (scaled from Najita & Shu 94) Total outflow momentum Friday, November 19, 2010 27 Core Star Formation Efficiency from Outflow Feedback Tan & McKee, in prep. Assuming angular distribution of momentum flux of standard X-wind or disk-wind models (Matzner & McKee 1999), and work out the maximum angle from the rotation/outflow axis at which core gas is expelled, assuming a steady wind that either stays coupled beyond the core radius or decouples. This sets εcore (Matzner & McKee 2000) Friday, November 19, 2010 28 Overview of Physical Scales Bontemps et al. (2010) nH~2x105cm-3 tff~1x105yr Butler & Tan 2009 AV=200 A8μm=8.1 NH=4.2x1023cm-2 Σ=4800 M pc-2 AV=7.5 A8μm=0.30 NH=1.6x1022cm-2 Σ=180 M pc-2 AV=1.4 NH=3.0x1021cm-2 Σ=34 M pc-2 Friday, November 19, 2010 29 Radiative Transfer Modeling Zhang & Tan, in prep. boundary of the core (11757 AU) expansion wave front (10189 AU) photosphere, z ∼ 3H UV (200nm) optical (550nm) sonic point (2537 AU) star disk J band (1.25 µm) H band (1.7 µm) K band (2.2 µm) 8 µm outflow cavity wall star disk Rsub = 5.51 AU Rd = 449 AU Friday, November 19, 2010 30 Radiative Transfer Models Zhang &Tan, in prep. see also: Robitaille et al. 2006; Molinari et al. 2008. Rotation and outflow axis inclined at 60˚ to line of sight. Σ = 1 g cm-2 Mcore = 60 M⦿ m* = 8 M⦿ mdisk = m*/3 Lbol = 6x103 L⦿ Friday, November 19, 2010 31 Radiative Transfer Models Zhang &Tan, in prep. see also: Robitaille et al. 2006; Molinari et al. 2008. Rotation and outflow axis inclined at 60˚ to line of sight. d=1kpc convolved with telescope beam Friday, November 19, 2010 Σ = 1 g cm-2 Mcore = 60 M⦿ m* = 8 M⦿ mdisk = m*/3 Lbol = 6x103 L⦿ 32 Observations: Mid IR Emission - Outflow Cavity DE BUIZER 11.7μm Vol. 642 18μm G35.2N (De Buizer 2006) LMIR ~ 1.6x103L⦿ Rotation and outflow axis inclined at 60˚ to (see Fig. 2b). Interestingly, this emission is extremely weak at K, bright at L , and not detected at 11.7 mm, but it is present line of sight. at 18.3 mm (source 3 in Fig. 1). The MIR images were also registered with respect to the high-resolution 8.5 GHz radio continuum m* = Gibb M⦿ images of 8 et al. (2003) (Fig. 2a, gray contours) and with the low-resolution 15 GHz radio continuum image of Heaton & Little (1988) Lastrometric error 3 L⦿ bol = 6x10 (Fig. 2b, white contours). The 1 j relative —The region of G35.20 0.74 in false color as seen at (a) 11.7 mm and (b) 18.3 mm with T-ReCS. Plus signs in (b) show the locations of individual rces at 11.7 mm and are numbered by increasing right ascension. The plus sign in (a) shows the 18.3 mm location of source 3, which is not seen at s The origin is the location of the radio continuum source G35.2N, R.A. p 18h 58m13.033, decl. p 01 40 36 . 14 (J2000) (A. G. Gibb 2006, private cation). 41 mJy arcsec 2 at 11.7 mm and 283 mJy arcsec 2 at . Sources 1, 2, and 10 are seen at 11.7 mm but not at , and source 3 is seen at 18.3 mm but not 11.7 mm. 4 is marginally detected at 18.3 mm. The remaining are detected at both wavelengths and are mostly knots ssion associated with the MIR monopolar jet of 0.74. The origin of Figure 1 is the expected location outflow source itself. This source is a B2.6 star (as from the 8.5 GHz flux density of Gibb et al. [2003] ng the method described in De Buizer et al. [2005]) be seen as an ultracompact H ii region in the radio been dubbed G35.2N. Relations to Radio Continuum and NIR Emission IR images were registered with respect to the NIR K images of Fuller et al. (2001). Very accurate relative try (!0 . 15) was achieved because of the presence of Friday, November 19, 2010 between the MIR and radio continuum images is estimated to be 0 . 34 in right ascension and 0 . 18 in declination. The MIR and NIR images and contours shown in Figure 2 have been s shifted 0.023 ( 0 . 35) in right ascension to place G35.2N on the infrared outflow axis (this is approximately the estimated 1 j astrometric uncertainty). In Figure 2b, it can be seen that the overall extent of the northern radio lobe is comparable to that of the MIR emission. There is also considerable MIR emission coming from the central radio continuum–emitting region near the outflow source; 33 MID-INFRARED JET OF G35.20 0.74 L59 Outflow-Confined HII Regions (Thermal Radio Jets) 18μm 10μm 8.6GHz IRAS 16562-3959 Guzmán et al. (2010) 10.4 µm emission associated with the central source and more diffuse emission associated Fig. 5.— Grey scale: TIMMI2 10.4 µm emission. Contours: 8.6 GHz radio emission. The cross marks the position of the OH maser associated with the central source. There is strong with the Inner-East lobe. A number of ionized HCHIIs seen in other nearby sources G35.2N (De Buizer 2006) (e.g. van der Tak & ent wavelengths. (a) The 11.7 mm image in false color overlaid with K-band emission from Fuller et al. (2001, Menten 2005 radio continuum emission of Gibb et al. (2003, gray contours). (b) The 18.3 mm image in false color overlaid image of Heaton & Little (1988, white contours) and L image fromOrion(2001, gray contours). (Tan & McKee 2003) Fuller et al. source I: c) Zoom false color, the L contours in white and the high-resolution radio continuum contours in black. The OH masers 15GHz terisks, water masers of Forster & Caswell (1989) as crosses, and methanol masers of A. G. Gibb (2006, private ower right show the 1 j relative astrometric uncertainty between the radio continuum and NIR. mediately north of the to ∼16,000 AU by a B2.6 star. If the dust is made of graphite, infrared emission here one could heat out to the distance of source 6 with grains having uum emission. There- 2010 a typical size of 0.005 mm, still near the lower size limit. Friday, November 19, 34 Highly-collimated outflows from massive protostars Evidence for similar accretion processes as low-mass SF e.g. Beuther et al. (2002) H. Beuther et al.: IRAS 05358+3543: Multiple outflows at the earliest stages of massive star formation 937 Fig. 8. Presented are the PdBI observations as contour overlays on the grey-scale H2 data (McCaughrean et al., in prep.). Friday, Novembera) Outflow (A): CO 1–0, red wing emission (v = [–14, –7] km s−1 , levels 5(10)95% from the peak intensity, PdB data only), and 19, 2010 35 Conclusions Is massive star formation a scaled-up version of low-mass star formation? No. 1, 2006 MID-INFRARED JET OF G35.20 0.74 L59 We see massive pre-stellar and starforming cores; rotating toroids; ordered B-fields; collimated outflows; outflowconfined HII regions (thermal radio jets). “Turbulent Core Model”: normalize core surface pressure to surrounding clump pressure, i.e. self-gravitating weight. The cores are probably are marginally magnetically super critical, limiting their fragmentation. Fig. 2.—The G35.20 0.74 jet as seen at different wavelengths. (a) The 11.7 mm image in false color overlaid with K-band emission from Fuller et al. (2001, white contours) and the 8.5 GHz high-resolution radio continuum emission of Gibb et al. (2003, gray contours). (b) The 18.3 mm image in false color overlaid with the low-resolution 15 GHz radio continuum image of Heaton & Little (1988, white contours) and L image from Fuller et al. (2001, gray contours). (c) Zoom in on the central region of the 11.7 mm image in false color, the L contours in white and the high-resolution radio continuum contours in black. The OH masers of Hutawarakorn & Cohen (1999) are shown as asterisks, water masers of Forster & Caswell (1989) as crosses, and methanol masers of A. G. Gibb (2006, private communication) as large plus signs. The bars at lower right show the 1 j relative astrometric uncertainty between the radio continuum and NIR. infrared emission coincident with and immediately north of the position of G35.2N demonstrates that the infrared emission here is dominated by longer wavelength continuum emission. Therefore, the nature of the infrared emission is concluded to be predominantly continuum dust emission from the outflow cavity walls. This cavity was created by the molecular outflow, which punched a hole in the dense molecular material surrounding the young stellar source at the center of G35.20 0.74. The central source is mostly likely directly heating the walls of this cavity. The northern lobe of the outflow was found to be slightly blueshifted toward Earth (i.e., in CO by Gibb et al. 2003; in C i by Little et al. 1998). Given this fortuitous geometry, we can see directly into the outflow cavity as a consequence of the clearing away of material along our line of sight by the outflow itself. The sources farther north of G35.20 0.74, namely, sources 5–9, are expected to be knots of dust either in the outflow itself or clumps of preexisting material that are being impinged upon by the outflow. Source 6 lies 19,200 AU from G35.2N and is still at an estimated dust color temperature of 112 K. This is based on the 11.7 and 18.3 mm flux densities of this source and neglects the possible effects of silicate absorption (see De Buizer et al. [2005] for method and limitations). What is heating the dust this far out? Smaller dust grains can be heated out to farther distances than large dust grains. The typical size range of interstellar grains is believed to be 0.003–10 mm, and typical grain compositions include smooth astronomical silicates, graphite, and silicon carbide (Laor & Draine 1993; Draine & Lee 1984). In the following I use the equation for dust temperature given by Sellgren et al. (1983) and the ultraviolet and infrared emissivities of Draine & Lee (1984). Assuming the dust is made up of smooth astronomical silicates, dust with a lower size limit of 0.003 mm can be heated to 112 K only out to ∼16,000 AU by a B2.6 star. If the dust is made of graphite, one could heat out to the distance of source 6 with grains having a typical size of 0.005 mm, still near the lower size limit. However, if silicon carbide is the assumed composition of the dust, then one can get heating out much farther than source 6, namely, ∼52,000 AU at the 0.003 mm lower size limit. There is a possibility of some contribution from shock heating, although Fuller et al. (2001) claim no detection of shock-excited H2 in the region. Beaming of the MIR emission along the outflow axis, rather than the isotropic emission assumed in the above calculations, could also help in heating grains farther out. Interestingly, the MIR luminosity derived from the dust color temperature gives an estimated value of 1.6 # 103 L,. Assuming the MIR luminosity is all the luminosity of the source (an obvious underestimate) and calculating a spectral type from that bolometric luminosity using the method from De Buizer et al. (2005) gives a value of ∼B3, consistent with the radioderived spectral type. In summary, all of the dust, even as far out as source 6, can indeed be heated directly by G35.2N, depending on dust composition and size (as well as beaming), though we cannot rule out contributions from other possible heating mechanisms. As discussed in § 3.1, MIR source 3, coincident with NIR emission from the presumed infrared southern counterjet, does not have a smoothly increasing spectral slope typical of dust continuum emission but instead is only present at L and 18.3 mm. This implies that the emission in this southern source is dominated by line emission of some kind. The usual suspects are (1) H2 emission from shocks, although Fuller et al. (2001) claim no detection of H2 in the region; (2) PAH emission from the photodissociation region of the outflow interface with the molecular cloud, although the L and 18.3 mm filters do not encompass any PAH features; and (3) [Fe ii] emission from If there is a different mechanism, e.g. competitive accretion, stellar mergers, then one would expect some break in the IMF at the mass scale it takes over. How can competitive accretion, i.e. accretion of dusty gas initially unbound to the protostellar core, overcome radiation pressure feedback for m*>10Msun? Mergers require unrealistic stellar densities. ...and, yes, I too have a theory for Orion! Friday, November 19, 2010 36 Herschel: Find us the Massive Pre-Stellar Cores... ...and, if you are brave, do multi-layered kinematics to understand their formation. Find us the Massive Protostars... ...we probably already know (HCHIIs, UCHIIs) -> Lbol ALMA: Kinematics of the core envelopes and disks... Magnetic field strengths and orientations... e-VLA and ALMA: more outflow-confined HII regions, please... Friday, November 19, 2010 37 ...
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This note was uploaded on 01/22/2012 for the course AST 1022l taught by Professor Colon during the Fall '07 term at University of Florida.

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