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3 on the cate the of stars. d movie ages. Comparison of low- and high-mass star formation — theory No. 1, 2010 C Krumholz D Jonathan Tan (University of Florida) Michael Butler (Zurich/MPIA) Audra Hernandez (Wisconsin) Karl Jaehnig Wanggi Lim Shuo Kong Bo Ma EBen Wu Yichen Zhang (Yale/U.Chile) OUTFLOW FEEDBACK REGULATED MASSIVE STAR FORMATION the massive star formation is simply part of the rapid “core” formation process. The use of sink particles in our simulation enables us to go beyond the calculations of Banerjee et al. (which stopped before significant mass accretion onto the stellar object) and follow the mass accretion for a long time, and to reveal a second phase of even more rapid accretion, driven by global collapse. The global collapse at the end of the pure hydro simulation (t = 0.63 Myr) was already shown in Figure 2. It is further illustrated in Figure 5, which shows the distribution of mass, velocity field, and contours of gravitation potential in a subregion centered on the most massive object, with the positions of stars superposed. It is clear that the global collapse has supplied the central region near the bottom of the gravitational potential well with plenty of dense material, which fuels both the enormous rate of total stellar mass accretion (close to the characteristic free-fall rate, see Figure 1) and the high accretion rate of the most massive object, which is but one of many stars in the region. In this crowded region, competitive rather than core accretion is likely at work. The reason is that the stars in are not enveloped by permanent host cores, because et this al.region (2009) Bonnell etonto al.stars (2001) the dense gas in the region is constantly drained and constantly replenished by the global collapse. The most massive object is located near the minimum of the global potential well and accretes the global collapse-fed material at a rate higher than any other object. Nevertheless, its accretion rate is only ∼10% of the total rate. In other words, the vast majority of the globally collapsing flow is diverted to stars other than the most massive object. Its accretion rate is large (∼2.5 × 10−4 M⊙ yr−1 ) during this late phase only because of an even larger global collapse rate. A common theme of the early and late phases of rapid accretion is that the dense gas that feeds the most massive object at a high rate is gathered by an agent external to the object. It is the converging flow set up by the initial turbulence in the early phase and the globally collapsing flow in the late phase. In this aspect, the formation mechanism is closer to core collapse than to competitive accretion. One may plausibly identify the dense filament in the early phase and the dense region at the bottom of the global gravitational potential well in the late phase as a McKee–Tan core (McKee & Tan 2003). However, the “cores” so identified are transient objects that are not in equilibrium. They are evolving constantly, with mass growing (from converging or collapsing flow) and depleting (into one or more collapsed objects) at the same time. The replenishment of dense gas in the “core” is the key to the long duration of high mass accretion rate for the most massive object, which is many times the freefall time of the small dense “core” and comparable to the global collapse time of the clump as a whole. In other words, there is no hint of any drop in the mass accretion rate after a finite reservoir of core material gets depleted over a (short) core free-fall time. Indeed, the nearly constant or increasing accretion rate for the most massive object highlights the important issue of when and how the mass accretion is terminated; there is an urgent need to tackle this issue if the final stellar mass is to be determined. We will return to a more detailed discussion of the nature of Sourav Chatterjee (NW) Nicola Da Rio Jan Staff Kei Tanaka Sven Van Loo (Leeds) 35 Wang et al. (2010) Figure 6. Slice of the density through the most massive object (denoted by an asterisk), showing a large magnetic bubble driven from near the central object at t = 3.75tff,c = 0.79 Myr in the MHD model. Overplotted are velocity arrows (white) and contours of the gravitational potential, which start at −10cs2 and decrease inward in steps of 20cs2 . plateau phases. The first starts somewhat later than that in Maite Beltrán (Arcetri) the HD case, indicating that the initiation of star formation in general and massive star formation in particular is somewhat Paola Caselli (MPE) delayed, as expected, because of magnetic cushion of turbulent compression. Once started, (SOFIA) the accretion rate onto the most James De Buizer massive object is, in fact, slightly higher in the MHD case than in the HD case, indicating that the magnetic field does not Francesco Fontani (Arcetri) significantly impede the initial phase of rapid mass accretion, which is enabled by the dense (UNC) structures formed by turbulent Matthew Goodson compression. Indeed, dense structure formation is aided by a strong magnetic field, which forces the turbulent flows to collide Jonathan Henshaw (LJMU) along the field lines (see, e.g., also Tilley & Pudritz 2007; Price Izaskun Jiménez-Serra (ESO) the dense & Bate 2008). When viewed in three-dimensional, structure resembles a warped, fragmented disk, with transient Jouni Kainulainen spirals that come and go. The (MPIA) spirals indicate significant levels of rotation on relatively small scales, which may hinder the Mark Krumholz stellar mass accretion due(UCSB) to centrifugal barriers. The centrifugal barrier may be weakened or even removed by Christopher magnetic braking. McKee Evidence for the(UCB) braking is shown in Figure 6, which shows a magnetic bubble driven from the central region, Romain Teyssier (Zurich) where active mass accretion is occurring. Magnetic braking driven bubbles provide a form of energy feedback that should accompany any star formed in a magnetized cloud (e.g., Draine 1980; Tomisaka 1998; Banerjee & Pudritz 2006; Mellon & Li 2008; Hennebelle & Fromang 2008) but is absent from purely Comparison of low- and high-mass star formation — theory & observation IRDC G028.37+00.07 NASA/Spitzer/IRAC+MIPS (Butler et al. 2014) IRDC G028.37+00.07 C1 NRAO/ALMA/N2D+(3-2) Tan et al. (2013) Orion Nebula Cluster (VLT; JHK) (McCaughrean) Jonathan Tan (University of Florida) Michael Butler (Zurich) Audra Hernandez (Wisconsin) Karl Jaehnig Wanggi Lim Shuo Kong Bo Ma Ben Wu Yichen Zhang (Yale/U.Chile) Sourav Chatterjee (NW) Nicola Da Rio Jan Staff Kei Tanaka Sven Van Loo (Leeds) Maite Beltrán (Arcetri) Paola Caselli (MPE) James De Buizer (SOFIA) Francesco Fontani (Arcetri) Matthew Goodson (UNC) Jonathan Henshaw (LJMU) Izaskun Jiménez-Serra (ESO) Jouni Kainulainen (MPIA) Mark Krumholz (UCSB) Christopher McKee (UCB) Romain Teyssier (Zurich) The Importance of Massive Stars Abel+ Vogelsberger+ Whitmore+ Gillessen+ McCaughrean+ O’Dell+ Massive Star Feedback and Regulation of SFRs 10 kpc 1 kpc Tasker & Tan (2009) ~1 pc Clump -> star cluster nH (cm-3): 102(green) 105 (red) ~100 pc GMC Van Loo, Butler, Tan (2013) Van Loo, Tan, Falle (2015) Butler, Tan, Van Loo (2015) The Physics of Massive Star Formation A complicated, nonlinear process: - Wide range of scales (~12 dex in space, time) and multidimensional. - Uncertain/unconstrained initial conditions/boundary conditions. Notation for the gas structures: Core -> star or close binary Clump -> star cluster Numerical models Observations Complete theory of star formation Analytic theory - 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. - Chemical evolution of dust and gas. - Stellar structure and evolution - Feedback • • • • • Massive Star Formation: Open Questions Causation: external triggering or spontaneous gravitational instability? Initial conditions: how close to equilibrium? Accretion mechanism: [turbulent/magnetic/thermalpressure]-regulated fragmentation to form cores vs competitive clump-fed accretion vs mergers Timescale: fast or slow (# of dynamical times)? Salpeter (1955) End result -2.35 – Initial mass function (IMF) – Binary fraction and properties N* dN*/dm* = A m* m*max? m* How do these properties vary with environment? Can we scale-up models of low-mass star formation to protocluster environments with Σ ~ 1 g cm-2? 198 “Standard” model of isolated low-mass star formation Shu, Adams, Lizano (1987) Many observed examples of this for low-mass star formation. Starless Core CFHT/Cuillandre Protostellar Core NASA/JPL/Tobin Protoplanetary Disk NASA/McCaughrean/OʼDell Can we find the equivalent stages for massive star formation? Or do the crowded cluster environments of massive star formation preclude this model? Massive Star Formation Theories Core Accretion: wide range of dm*/dt ~10-5 - 10-2 M! yr-1 (e.g. Myers & Fuller 1992; Caselli & Myers 1995; McLaughlin & Pudritz 1997; Osorio+ 1999; Nakano+ 2000; Behrend & Maeder 2001) Turbulent Core Model: (McKee & Tan 2002, 2003) Stars form from “cores” that fragment from the “clump”. If in equilibrium, then self-gravity is balanced by internal pressure: B-field, turbulence, radiation pressure (thermal P is small) Cores form from this turbulent/magnetized medium: at any instant there is a small mass fraction in cores. These cores collapse (quickly) to feed a central disk to form individual stars or binaries. Competitive (Clump-fed) Accretion: (Bonnell et al. 2001; Wang et al. 2010) Massive stars gain most mass by Bondi-Hoyle accretion of ambient clump gas. Massive stars form on the timescale of the star cluster. Violent interactions? Mergers? (Bonnell et al. 1998; Bally & Zinnecker 2005) Schematic Differences Between Massive Star Formation Theories massive prestellar core Rare evolution from magnetically subcritical state? massive-star-forming core [protostar+gravitationally-bound gas] massive-protostar (MP) massive Kunz & Mouschovias (2009) star m*f>8M! Outflowconfined HII Region Turbulent core model (MT02, 03) time Isolated massive star formation? (Bressert+2012; Oey+13) Competitive Bondi-Hoyle accretion model (Bonnell ea. 2001; Bonnell & Bate 2006; Dobbs+, R. Smith+, P. Clark+) Prestellar core mass function? (e.g. Motte et al. 1998; Testi & Sargent 1998; Alves et al. 2007) t=0 protostar formation m*=8M! Radiation pressure likely to prevent accretion of dusty, unbound gas (Edgar & Clarke 2004) What are the appropriate initial conditions for theoretical models? Do massive starless cores exist? How do we find them? Mid-IR Extinction Mapping of Infrared Dark Clouds (Butler & Tan 2009, 2012; see also Peretto & Fuller 2009; Ragan et al. 2009; Battersby et al. 2010) G28.37+00.07 Spitzer IRAC 8µm 16’ (GLIMPSE) (Churchwell et al. 2009) MJy sr-1 Mid-IR Extinction Mapping of Infrared Dark Clouds (Butler & Tan 2009, 2012; see also Peretto & Fuller 2009; Ragan et al. 2009; Battersby et al. 2010) G28.37+00.07 Spitzer IRAC 8µm 16’ (GLIMPSE) Median filter for background around IRDC; interpolate for region behind the IRDC g cm-2 Correct for foreground using “saturated” cores ~Arcsecond scale maps of regions up to Σ ~0.5 g cm-2; independent of dust temp. Distance from molecular line velocities -> M(Σ) MJy sr-1 V deep 8 µm Spitzer-IRAC imaging of massive Infrared DarkCloud Cloud (I se deep 8 µm Spitzer-IRAC imaging of massive Infrared Dark (IRD y extinction mapping. Merging with an NIR extinction map of the region nfrared (MIR) extinction map that probes mass surface densities -infrared (MIR) extinction mass densities upup to veals structures down to AV ∼ map 1 mag.that Weprobes utilize the mapsurface to: (1) measure stgst yet by extinction extinction mapping. Merging with Sample ofkinematics ~50 massive “starless” thehighest highest values yet probed probed by mapping. Merging with an a iusthe of ∼8 pc. 13values CO indicate that the cloud iscore/clumps gravitationally (Butler & Tan 2012; Butler et al. 2014) dynamic range reveals structures down 1 magW V∼∼ of the most massive youngthat starreveals clustersstructures known indown the Galaxy. esaone ahigh high dynamic range map map that totoAA 1 mag. V (2) 4 densities (M=60M!) Mass surface 1313 COpolytropic 10 M radius of ∼8 pc. kinematics indic emass cores within findinga can of be∼8 fit by singular ud mass∼7 ∼7××the 104IRDC, M⊙⊙ within within athey radius pc. CO kinematics indicate −2 form one of the most massive young star It They has potential 3. have ≃ 0.1–0.4 g cmto —relatively along young with star c d. Itthus thus hasΣthe the potential one oflow thevalues mostthat, massive t that magnetic fields, rather thanmassive accretion-powered radiative heating, are terize the of 16 cores within within theIRDC, IRDC,finding finding the acterize thestructures structures of concentration 16 massive cores the they Cores show central of these (3) the Σ (equivalently column density orgAcm ) −2−2 ρ ρ Determine 1.5±0.3 andkkρρ = = 1.3 ΣΣ≃≃ 0.1–0.4 —rela eswith withρcores. ρ∝∝r r−k−k and ± 0.3. 0.3. They Theyhave have 0.1–0.4 gV cm —r r a region that is nearly complete for AV > 3 mag. The PDF is well fit by measured coldtemperatures, temperatures, suggest that than accreti easured cold suggest thatmagnetic magneticfields, fields,rather rather than accr Contain many mag, high compared toJeans other masses. known clouds. It does not exhibit a separate rtant for controlling fragmentation fragmentation of (3) the Σ (eq ant for controlling of these thesecores. cores. (3)Determine Determine Fragmentation not suppressed claimed to indicate the importance of self-gravity. However, we suggest the Σ ( bility distribution function (PDF) for complete AV by radiative function heating (c.f., lity distribution (PDF) for aa region region that nearly complete lar, self-gravitating hierarchy of structures presentthat overisis anearly wide range of forfor Krumholz & McKee gle log-normal with mean2008). AVV ≃ ≃ 99 mag, clou e log-normal with mean A mag, high highcompared comparedtotoother otherknown known c end power lawtail, tail,which which has has been been claimed of os d –power law claimedtotoindicate indicatethe theimportance importance ds stars: formation hePDF PDFdoes doesresult resultfrom from aa self-similar, self-similar, self-gravitating structu self-gravitatinghierarchy hierarchyofof stru sninthe thecloud. cloud. Fragmentation suppressed by B-fields? Critical– Mass McKee 1992) words: Magnetic dust, extinction ISM: (Bertoldi clouds – &stars: formation rds: dust, extinction – ISM: clouds – stars: formation e-only material: color figures ′′ telescope; ∼22 for the 250 µm observations of Peretto et al. only material: color figures -3, B~200μG -> M ~100 M nH~105cm B 2010 with Herschel-SPIRE). ! clusters from Interferometric observations are possible at submillimeter and longer wavelengths, but on their own provide poor constraints on dust temperature. Dynamics of Massive Starless Cores: Are they close to Virial Equilibrium? The Astrophysical Journal, 754:5 (22pp), 2012 July 20 Butler & Tan Four IRDC core/clumps selected to be dark at 8, 24, 70 μm C1 F1 G2 F2 So use high angular resolution observations of N2D+(3-2) to 1. Identify exact location of (massive) starless cores 2. Measure core velocity dispersion, σ. High Deuterium Fraction [N2D+]/[N2H+] (Fontani et al. 2011) n CO freeze-out e.g. Hernandez et. al (2011) H3+ + CO → HCO+ + H2 (T<20K) H3 + HD → H2D+ + H2 + H2D+ + N2 → H2 + N2D+ aged radial profile figures (notation as in Figure 5(b)). ournal.) IRAM 30m Astrochemical Figure 10. Core E2, E3, F1, F2, F3, and F4 Σindicator maps (notation asthat in Figurethese 5(a)) and azimuthally av (A color version of this cores figure is available in the onlin are starless (Caselli et al. 2002) 13 C1 F2 G2 0.91 0.250 0.070 C1-N Tan, Kong et al. (2013) 0.246 -0.564 0.244 -0.566 C1-S G2-N 0.248 b (degrees) “Massive” “Starless” “Cores” with ALMA F1 G2-S 0.74 0.068 0.246 0.242 -0.568 0.240 -0.570 0.238 -0.572 0.066 0.244 0.58 0.064 28.328 28.326 28.324 28.322 28.32 4.424 34.422 34.420 34.418 34.416 0.250 440 34.438 34.436 34.434 34.432 34.784 0.246 -0.564 0.244 -0.566 0.242 -0.568 0.240 -0.570 34.782 34.780 34.778 34.776 0.070 C1-S 0.248 G2-N 0.068 G2-S 0.246 0.33 0.066 0.244 0.064 -0.572 0.238 28.328 28.326 28.324 28.322 28.32 4.424 34.422 34.420 34.418 34.416 440 34.438 34.436 34.434 34.432 34.784 34.782 34.780 34.778 34.776 0.23 0.250 0.246 -0.564 0.244 -0.566 0.242 -0.568 0.240 -0.570 0.238 -0.572 0.070 DCO+(3-2) b (degrees) DCO+(3-2) C1-N C1-S 0.248 G2-N G2-S 0.068 0.246 0.15 0.066 0.244 0.082 0.064 28.328 28.326 28.324 28.322 28.32 4.424 34.422 34.420 34.418 34.416 0.250 440 34.438 34.436 34.434 34.432 34.784 0.246 -0.564 0.244 -0.566 0.242 -0.568 0.240 -0.570 34.782 34.780 34.778 34.776 1.3mm 1.3mm continuum b (degrees) 0.070 0.036 C1-N C1-S 0.248 G2-N 0.068 0.246 G2-S 0.009 0.066 0.244 0.064 -0.572 0.238 28.328 28.326 28.324 l (degrees) 28.322 28.32 4.424 34.422 34.420 34.418 l (degrees) 34.416 440 34.438 34.436 34.434 l (degrees) 34.432 34.784 34.782 34.780 l (degrees) 34.778 34.776 0 Mass Surface Density (g cm-2) N2D+(3-2) b (degrees) N2D+(3-2) 0.45 C1-N Comparison to Turbulent Core Model core C1, ΣMIREX, N2D+(3-2) contours Σcl Mc clump 3.6” 0.09pc Core masses inside 3σ N2D+ contour: Σcl = 0.36 g cm-2 Mc,MIREX = 55.2 ± 25 M! Mc,mm = 62.5 129 26.9 M! ALMA beam Spitzer beam Predictions from Virial Equilibrium • 1D velocity dispersion if virialized: ( = 1) Core C1-N C1-S F1 F2 G2-N G2-S Σcl (g cm-2) 0.48 0.40 0.22 0.32 0.21 0.19 Mc 4 (M!) 16 6.5 4.7 2.4 0.83 Tan et al. 63 σvir (km/s)For0.66±0.22 0.88±0.30 the sample of 6 cores the0.43±0.15 ratio of the 0.44±0.15 observed to 0.33±0.11 the predicted0.25±0.09 virial equilib- rium velocity dispersion is 0.810.97 , while the ratio of the observed to predicted size is 0.71 σobs (km/s) 0.41±0.03 0.41±0.02 0.25±0.02 0.42±0.04 0.34±0.02 0.30±0.02 1.541.97 1.26 . Thus the cores appear to be close to the predictions of the model. For the most < σobs = 0.81 0.13 than massive core, C1-S, the observed velocity dispersion is /σ about a factor of 2±smaller vir > the predictions of the fiducial model. For it to be in virial equilibrium would require stronger magnetic fields of ∼ 1.0 mG, implying mA ≃ 0.3. In fact, given the core den5 cm= −30.28 -> Bvir=0.9mG sity of nH ≃ 6 ×m 10A,vir , the predicted median B-field strength using Crutcher et al.’s 0.65 (2010) relation Bmed ≃ 0.12nH µG (for nH > 300 cm−3 ) is(Crutcher 0.7 mG. et al. 2010) nH,c=6.4x105cm-3 -> Bmed=0.7mG (see also Pillai+ 2015) 4. The Chemical “Deuteration” Ages of the Cores Tentative Conclusion: Cores appear to be near virial equilibrium, after accounting for clump envelope. Possibly slightly sub-virial; or have stronger The Turbulent Core Model does not make any prediction about the timescale for the B-fields. But we need a larger sample - massive starless cores are rare! cores to assemble, except that at least one dynamical (i.e. or ≃ 2 free-fall times, tff ) is needed to reach an equilibrium state. If strong magnetic fields are regulating core formation, then the timescale for core formation could be considerably longer than this. CO freeze-out e.g. Hernandez et. al (2011) H3+ + CO → HCO+ + H2 (T<20-30K for small OPRH2) H3+ + HD → H2D+ + H2 H2D+ + N2 → H2 + N2D+ Observed Dfrac of C1-S Collapse rate αff ≲ 0.3 # =1 αff Model deuteration of N2H+ with “complete,” gas-phase, spin-state astrochemical network to constrain age and/or rate of collapse. αff#= Kong, Caselli, Tan, Wakelam, Sipilä (2015); see also Pagani et al. (2009, 2013) .01 The Deuteration Clock Summary of Initial Conditions • Massive starless cores exist • Require strong B-fields for dynamic stability and to prevent fragmentation • Chemically old compared to free-fall time Constraints for Initial Conditions of Numerical Simulations Peters et al. (2011) M = 100M!, R=0.5pc, nH = 5400cm-3, B=10μG Seifried et al. (2012) M = 100M!, R=0.25pc, nH = 4.4x104cm-3, B~1mG Myers et al. (2013) M = 300M!, R=0.1pc, Disc formation in turbulent mass 6cm-3, B>~1mG nH = 2.4x10 8 disc 1 disc 2 disc 3 disc 4 7 6 µ 5 4 3 2 1 0 0 5 10 t / kyr 15 0 5 t/ Figure 3. Mass-to-flux ratio µ (left) and incli magnetic field to the angular momentum vecto in spheres with a radius of 500 AU around th found in run 2.6-4-A. purpose we calculate the volume-weighted field hBi in a sphere with a radius of r = the CoM of each disc. In combination with M we obtain the mass-to-flux ratio Figure 2. Column density in logarithmic scaling for the top-on view of disc 1 (top left) and disc 2 (top right) of run 2.6-4-A and of the disc in run 2.6-4 without turbulence (bottom). The figures are 800 AU in size. lines. This is a remarkable result since for previous simulations of low- and high-mass cores with mass-to-flux ratios µ . 10 only sub-Keplerian discs were found (e.g. Allen et al. 2003; Price & Bate 2007; Mellon & Li 2008; Hennebelle & Fromang 2008; Duffin & Pudritz 2009; Seifried et al. 2011). We emphasise that for the other runs we find qualitatively similar results, i.e. discs with sizes of up to ⇠ 100 AU and masses of the order of 0.1 M . The number of discs per µ= M ⇡r2 | hBi | 0.13 /p . G We plot the time variation of µ in the lef for the same four discs as in Fig. 1. As can around a mean of 2 - 3. Hence, the val agree with the overall value of 2.6 and a the value of ⇠ 2 found in run 2.6-4. More the range where simulations without turbu found sub-Keplerian discs only. We therefo turbulent reconnection is not responsible fo Keplerian discs in our runs. Another way of reducing the magnetic was investigated by Hennebelle & Ciardi ( What is the accretion mechanism of massive protostars? Outflowconfined HII Region Continuum Radiative Transfer Modeling Temperature Density Robitaille, Whitney et al. (2006+); Indebetouw et al. (2006); Molinari et al. (2008); Zhang & Tan (2011), Zhang et al. (2013, 2014) hydrostatic core expansion wave rotating infall active accretion disk disk wind dust destruct. front gas + dust opacities protostellar evolution Parameters: Σclump = 1 g cm-2 Mcore = 60 M⦿ β = 0.02 m* = 8 M⦿ Lbol = 6x103 L⦿ Continuum Radiative Transfer Modeling Zhang, Tan & Hosokawa (2014) The Astrophysical Journal, 100 788:166 AU (35pp), 2014 June 20 1000 AU 20,000Zhang, AU Tan, & Hosokawa Figure 5. Input density and converged temperature profiles for the fiducial model (Mc = 60 M⊙ , Σcl = 1 g cm−2 , βc = 0.02) at selected stages with m∗ = 1, 2, 4, 8, density temperature temperature density temperature 12, 16, and 24 M⊙ . At each stage (each row), these profiles aredensity shown on three different scales (from left to right, 100 AU, 1000 AU, and 20,000 AU). In each panel, the protostar is at the (0,0) point, the x-axis lies on the disk midplane, and the y-axis along the outflow axis. The velocity fields are shown by the arrows. Note that the black arrows for the outflow have a much larger scale than the blue arrows for the infalling envelope. (A color version of this figure is available in the online journal.) 11 The Astrophysical Journal, 788:166 (35pp), 2014 June 20 Zhang, Tan, & Hosokawa 1M⦿ 2M⦿ 4M⦿ NIR to FIR morphologies Rotation and outflow axis inclined at 60˚ to line of sight. 8M⦿ 12M⦿ 16M⦿ 24M⦿ Figure 21. Resolved images for the selected evolutionary stages (m∗ = 1, 2, 4, 8, 12, 16, and 24 M⊙ , from top to bottom) of the fiducial model in various bands (columns) at the inclination of 60◦ between the line of sight and the axis. Each image is normalized to its maximum surface brightness, which is labeled in the bottom left corner. The total fluxes are labeled in the top right corners. A distance of 1 kpc is assumed. Each image has a field of view of 40′′ × 40′′ . The dotted lines mark the projected opening angle of the outflow cavity on the sky plane. Prediction: increasing symmetry from MIR-FIR Massive Protostar G35.2N: d=2.3kpc; L~105L⦿ Gemini-T-ReCS SOFIA-FORCAST radio (cm) continuum (ionized gas) De Buizer (2006) Zhang, Tan, De Buizer et al. (2013) Spectral energy distribution Flux profiles along outflow cavity axis o t t fi d o o Σ =1g g s e . M = 240 M d d i g e l v i n i o a m = 34 M t r m e p r l d o f e : s d s i e o l r fi m a t o c r s i p e tΣrcore/clump MIR SED requires high y v e i t i s m ens s a m y t m s n i , a e le w g p o a m h m i Si n o & s t D L ~ (0.66 - 2.2)×10 L n E i S stra M ~ 240M n co Ʃ ~ 0.4 - 1 g/cm clump cm-2 core * ⦿ ⦿ 5 ⦿ bol core t a s e g t l s i h f fi l 1 t P i c o W T fl ⦿ cl 2 θw ~ 35 - 51˚ θview ~ 43 - 58˚ m* ~ 20 - 34 M⦿ Fig. 4.— Intensity profiles along the outflow axis. The squares are observational data sampled at intervals of the resolutions of n u e Z s w m Massive stars play a crucial role in the production of heavy elements and in the evolution of the interstellar medium, yet how they form is still a matter of debate. We report high-angular-resolution submillimeter observations toward the massive hot molecular core (HMC) in the high-mass star-forming region G31.41+0.31. We find that the evolution of the gravitational collapse of the HMC is controlled by the magnetic field. The HMC is simultaneously contracting and rotating, and the magnetic field lines threading the HMC are deformed along its major axis, acquiring an hourglass shape. The magnetic energy dominates over the centrifugal and turbulence energies, and there is evidence of magnetic braking in the contracting core. lecular core surrounding the protostars figuration that was shown to be con theoretical models for the formation o stars, where well-ordered, large-scale turbulent, magnetic fields control th and collapse of the molecular cores stars form (12). We investigated the hot molecular in G31.41+0.31 (G31.41), a massive region [~500 to 1500 M◉ (13, 14)] l parsecs (pc) away (15). G31.41 has Continuities between Low- & High-Mass Star Formation S tars more massive than 8 M◉ (where M◉ is the mass of the Sun) account for only 1% of the stellar population in our Galaxy. Nevertheless they dominate the appearance and evolution of its interstellar medium and are responsible for the production of heavy elements. The formation of massive stars is not completely understood. Stars form when dense molecular clouds collapse as a result of gravity. But as the mass of a young star reaches 8 M◉, its own 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- Core Mass Function: e.g., Nutter & Ward-Thompson (2007) Protostellar Cores & B-field morphologies: hibition by shielding the infalling material from stellar radiation and by creating a lower density section along the rotation axis of the disk and a molecular outflow, which helps by channeling the radiation out, allowing the formation of stars more massive than 40 M◉ (3–5). Massive stars may also form through mergers of smaller stars (6). The scenario whereby massive stars form through disk-assisted accretion resembles the way stars like the Sun form. Both processes 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- 1 Institut de Ciències de l’Espai [Consejo Sup tigaciones Científicas (CSIC)–Institut d’Estudi (IEEC)], Campus Universitat Autònoma de Ba Facultat de Ciències, Torre C5 - parell 2a, 08 Catalunya, Spain. 2Departament d'Astronomia (IEEC-UB). Institut de Ciencies del Cosmos y Un CSIC, Universitat de Barcelona, Martí i Fran Barcelona, Catalunya, Spain. 3Harvard-Smithso Astrophysics, 60 Garden Street, Cambridge, M 4 Academia Sinica, Institute of Astronomy and A North Aohoku Place, Hilo, HI 96720, USA. *To whom correspondence should be addr [email protected] †Present address: Osservatorio Astrofisico d Enrico Fermi 5, 50125 Firenze, Italy. A B Girart et al. (2009); Q. Zhang et al. (2014) Outflow properties: - collimated morphologies (Beuther et al. 2002; Duarte-Cabral et al. 2013) [however, BN/KL - Bally & Zinnecker 2005] - momentum flux vs. luminosity - CO ladder (talk by van Dishoeck) Astrochemical properties: - Deuteration (Fontani et al. 2011, 2015) - Hot Cores/Corinos 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 angle of the magnetic field. These maps were obtained by using a natural weighting to the visibility data, which yielded to a full width at half maximum synthesized beam of 1.34″ × 0.83″ with a position angle of 67° (shown in 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 velocity map of the CH3OH 147-156 A. Black thick bars indicate of the magnetic field. These maps were obtained by using a robu Girart et al. (2009) of 0 to the visibility data, which yielded to a full width at ha 1408 12 JUNE 2009 VOL 324 synthesized beam of 1.04″ × 0.59″ with a position angle of 82° ( bottom left corner). Contour levels are the same as in the previou a peak intensity of 6.55 Jy per beam. (C) Spectrum of the C34S 7position of the dust emission peak. The continuum has been sub the line emission (this is valid for all the molecular line data pres SCIENCE www.sciencemag.org Duarte-Cabral et al. (2013) Conclusions Core Accretion, Competitive (Clump-fed) Accretion, Protostellar Mergers are potential formation mechanisms of massive stars. Key questions: Does clump fragment into cores with range of masses, including massive starless cores? If so, what prevents fragmentation? Magnetization strongly affects the outcome of numerical simulations. “Turbulent Core Model”: - Core surface pressure set by surrounding clump pressure - Core supported by nonthermal pressure (B-fields/ turbulence), but otherwise this is a scaled-up version of lowmass star formation - mG B-fields are likely to be preventing fragmentation - Massive protostars will ionize their outflows No. 1, 2006 Observations & Modeling: - Massive, deuterated, starless cores exist and may be near virial equilibrium, with slow collapse compared to free-fall - Massive protostars can have a similar morphology to low-mass protostars. B-field to regulate collapse & fragmentation. - Can sometimes be influenced by cluster environment (e.g., BN/KL) MID-INFRARED JET OF G35.20!0.74 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 white contours) and the 8.5 GHz high-resolution radio continuum emission of Gibb et al. (2003, gray contours). (b) Th with the low-resolution 15 GHz radio continuum image of Heaton & Little (1988, white contours) and L′ image from Full in on the central region of the 11.7 mm image in false color, the L′ contours in white and the high-resolution radio conti of Hutawarakorn & Cohen (1999) are shown as asterisks, water masers of Forster & Caswell (1989) as crosses, and meth communication) as large plus signs. The bars at lower right show the !1 j relative astrometric uncertainty between the infrared emission coincident with and immediately north of the to ∼16,000 AU by a B2.6 sta Figure 3: Cycle 2 observations of Cep A at 7.7µm (left) & 37.1µm (right). Similar to G35.2, at shorter wavelengths position of cavity G35.2N demonstrates that wavelengths the infraredthe emission here couldDetailed heat out to the distan only the near-facing outflow is seen, while at longer far-facing cavity is one revealed. by longer emission.models There-& test formation a typical size of 0.005 mm, RT models of thisis&dominated the other sources will bewavelength carried out tocontinuum constrain protostellar theories. Mass surface density → Σ - M Diagram Giant Molecular Clouds Star-Forming Clumps Young Star Clusters IRDC Clumps & Cores Physical Properties of Star-Forming Regions AV=230 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 Tan et al. (2014, PPVI) Mass → Fig. 1.— The Environments of Massive Star Formation. Mass surface density, ⌃ ⌘ M/(⇡R2 ), is plotted versus mass, M . Dotted