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¿QUÉ SABEMOS DE LA FORMACIÓN DE ESTRELLAS MASIVAS? Miriam Garcia

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¿QUÉ SABEMOS DE LA FORMACIÓN DE ESTRELLAS MASIVAS? Miriam Garcia ¿QUÉ SABEMOS DE LA FORMACIÓN DE ESTRELLAS MASIVAS? Miriam Garcia IX REUNIÓN CIENTÍFICA DE LA SEA 14 SEPTIEMBRE, MADRID ¡POCO! ¡¡Gracias por su atención!! IX REUNIÓN CIENTÍFICA DE LA SEA 14 SEPTIEMBRE, MADRID OutlineOutline • Introduction • Main stages of massive star formation • Proposed scenarios of the accretion phase • The observed stages • Feedback • Stellar masses: • Spectroscopic mass: (need to know Mdot) • Evolutionary mass • Dynamical mass (but few eclipsing binary systems known, plus binarity modifies massive star evolution in a way we do not fully undestand yet.) • Obscured IR clusters IX SEA 2010, 14/SEPT/2010 INTRODUCTIONINTRODUCTION IX SEA 2010, 14/SEPT/2010 GeneralGeneral processprocess • A gas condensation (core) collapses inside a larger unit of the molecular cloud (clump). • The protostar forms, • it increases its mass by accretion • with simultaneous mass loss through bipolar outflow and/or collimated jet • Initial conditions • Dense compact clump: 10 23 – 10 24 cm -2 • Cold cores: 10 5 cm -3, 10-20K, 0.5pc diameter Bate et al. UK Astroph. fluids facility IX SEA 2010, 14/SEPT/2010 MainMain stages stages 5 From a Giant Molecular cloud (10 M⊙): 1. Fragmentation /compresion: formation of cold cores (100 M⊙) (induced by external SN shock, UV rad, ρ-wave) 2. Collapse into protostellar embryos that can heat up the core (ρ increases, κ increases, the IR does not escape). Collimated jets and outflows appear. 3. Disk-accretion onto the protostar as it evolves towards MS Massive stars start H-burning before the ZAMS and develope radiation-driven winds Formation of hyper-compact HII region from the photoionized and partially photoevaporated disk. 4. Disruption of the cloud by the combined action of winds, outflows, and UV radiation IX SEA 2010, 14/SEPT/2010 UCHII -> CHII -> normal HII EmbeddedEmbedded pre-main pre-main sequence sequence stages stages To be observed at mid-IR and radio wavelengths (see summary by Menten et al. 2005; van der Tak & Menten 2005): • IR dark clouds • Hot molecular cores • Hypercompact and ultra-compact HII regions • Compact and classical HII regions (can be seen in optical range) So, if you are looking for YSO go to: • Dark clouds • Bright-rimmed pillars • Compact HII regions IX SEA 2010, 14/SEPT/2010 IR dark clouds • Significant mid-IR opacity • Cold (<15 K) Glimpse, 8 µm • Dense (>10 4 cm -3) • Large column densities (>10 23 –10 25 cm -2) • Believed to be the starting point of high mass star formation Perault et al. 1996; Egan et al. 1998; Carey et al. 1998, 2000; Hennebelle et al. 2001; Simon et al. 2006; Rathborne et al. 2006 IX SEA 2010, 14/SEPT/2010 AccretionAccretion (or (or how how toto outpower outpower radiation radiation pressure) pressure) • The problem: massive stars ignite H at 20 M⊙, before the accretion stage is over (Krumholz et al. 2009) . 40 M⊙ in models with rotation (Yorke & Sonnhalter 2002) . • Three main mechanisms: (see reviews by Zinnecker & Yorke 2007; Beuther et al. 2007) • Monolithic collapse (Yorke & Sonhalter 2002; Krumholtz 2007, 2009) • Competitive accretion (Bonnell et al. 1997, 2001) . Can reproduce the apparently universal IMF. • Stellar mergers • Triggered OB star formation (Elmegreen & Lada 1997): stars formed in shock matter (from rad. pressure of other OB stars, SN, …) are more massive. IX SEA 2010, 14/SEPT/2010 KrumholzKrumholz et et al.al. 20092009 • Radiation pressure does not halt accetion in 3D models: “- Gravitational + Rayleigh-Taylor instabilities channel gas onto the star system through non-axisymmetric disk that self- shield against radiation, and allow radiation to escape through optically-thin bubbles. -The instabilities in accretion disks, and the resulting formation of companion stars, can only be simulated when the disk is represented non-axisymmetrically. It is important to consider these effects because most massive stars are members of multiple system”. • Simulation. Initial cloud properties: • Mcore = 100 M⊙ • Radius = 0.1 pc • Density profile ρ∝r−1.5 • T=20K • Slow, solid-body rotation IX SEA 2010, 14/SEPT/2010 FeedbackFeedback IX SEA 2010, 14/SEPT/2010 THEORYTHEORY MEETSMEETS OBSERVATIONSOBSERVATIONS IX SEA 2010, 14/SEPT/2010 Current mass Initial mass Contrast to predictions & favor one or another scenario IX SEA 2010, 14/SEPT/2010 TheThe maximum maximum predicted predicted stellar stellar mass mass Figer (2005) • Obvious limit: when radiation pressure stops accretion (at 60 M⊙, 40 M⊙ or 20 M⊙, depending on the author!) • But now, radiation pressure can be overcome with jets • (In opacity limited accretion, limit set at 200 M⊙) IX SEA 2010, 14/SEPT/2010 TheThe IMF IMF • φ(m)dm ∝ m-2.35 dm or d(log N)/d(log m)= Γ • Figer (2005) from populations of very massive young clusters: Max(M ini )=150 M⊙ • Oey & Clarke (2005) statistical approach: 120-200 M⊙ • Koen (2006) 140-160 M ⊙ • Also: it takes a GMC 5 Figer (2005) (10 M⊙) to make a 50 M⊙ star (Larson, 1982). IX SEA 2010, 14/SEPT/2010 EMPIRICALEMPIRICAL STELLARSTELLAR MASSESMASSES IX SEA 2010, 14/SEPT/2010 EvolutionaryEvolutionary masses: masses: A2 Ia B1.5 Ia O9.5 III Garcia et al. (2010) (Spectral types: Bresolin et al. 2007, Garcia et al. in prep.) IX SEA 2010, 14/SEPT/2010 AUTOPOP:AUTOPOP: automaticautomatic finding finding & & characterizationcharacterization ofof OB OB associationsassociations Photometric Catalog + target selection criteria Isochrones Friends of Friends Automatic M, Age determin. (plus prop. in isochrones) Masses and ages associations members VO-friendly tables to Garcia (2009, 2010) al. et X-match with cats IX SEA 2010, 14/SEPT/2010 SomeSome results results on on IC IC 16131613 Log age= 6.5 , 6.8 , 7.1 , 7.4 , 7.7 , 8.0 Garcia et al. 2010 IX SEA 2010, 14/SEPT/2010 SpectroscopicSpectroscopic mass: mass: • From gravity and radius: M= gR 2/G • log g : Quantitative spectroscopic analysis of the 4000-5000 Å range, of Balmer lines (Hβ, Hγ and Hδ) with FASTWIND (Puls et al. 2005) . • R from V-mag + A v + distance, or calibrations • Hα to estimate the mass loss rate, and infer M ini (also need age). IX SEA 2010, 14/SEPT/2010 ToTo determine determine massmass loss loss rate… rate… • Can be done with Hα (better characterized) but UV lines more sensitive Mdot, and also may provide terminal velocity. • Need also to constrain clumping filling factor • Panchromatic spectroscopic studies with CMFGEN: IX SEA 2010, 14/SEPT/2010 HybridHybrid spectroscopic-evolutionaryspectroscopic-evolutionary massmass Quantitative spectroscopic analysis of the 4000-5000 Å range: • Effective temperature: • Silicon lines for B superg. ( SiIII4552,4567, 4574, SiII4128 and SiIV4116 ) • HeI and HeII for O stars • log g BEWARE OF R !! prep in 2010, al. et Simón-Díaz IX SEA 2010, 14/SEPT/2010 BUT!!BUT!! MassMass discrepancydiscrepancy • Reported by Herrero et al. (1992) • Partly reconciled when a more detailed description of the OB star atmospheres was used to derive the spectroscopic masses (Herrero, Puls, & Najarro, 2002; Repolust, Puls, & Herrero, 2004; see also Puls 2008) • However, some discrepancy remains for: • Galactic B-type supergiants (Markova & Puls, 2008) • OB stars in the LMC and SMC (Massey et al., 2009) • IC 1613 (Garcia et al. 2010) IX SEA 2010, 14/SEPT/2010 TheThe IACOB IACOB databasedatabase • IACOB database: high resolution spectroscopy of Galactic OB stars (Simón-Díaz et al. 2010). • Large sample of OB stars with spectroscopic and evolutionary masses. • Tackle the mass discrepancy • Build the upper end of the IMF IX SEA 2010, 14/SEPT/2010 PAUSAPAUSA FRANCESA:FRANCESA: UNCERTAINTIESUNCERTAINTIES ATAT WORKWORK IX SEA 2010, 14/SEPT/2010 • Unknown extinction (can be overcome with mutiwavelength photometry as opposed as to only use B-V). • Uncertainties in distance (that propagate to R*): • From photometric analysis • Spectroscopic paralax (which relies on calibrations of radius with spectral type, which in turn require previous distance determinations) • Physical paralaxes : • Hipparcos (not so good if outside of Solar neighborhood Maíz-Apellániz 2001, 2008) • GAIA will play a crucial role with its improved paralaxes • Unresolved binaries (real or projected), which increase the apparent V-magnitude (HD64315, S. Simón-Díaz, J. Lorenzo- Espinosa ). IX SEA 2010, 14/SEPT/2010 WHEREWHERE ISIS THETHE MOSTMOST MASSIVEMASSIVE STAR?STAR? IX SEA 2010, 14/SEPT/2010 So…So… where where do do wewe find find the the most most massive massive stars? stars? • Considering a Salpeter IMF and that the cluster has not lost any member to SN explosion (age < 3Myr), 4 A cluster of Mtot > 10 M⊙ is needed to produce a 150 M⊙ star (Figer 2005, see also Crowther et al. 2010) 30 Dor in the LMC: • R136: 150 M ⊙ (Parker & Garmany 1993, Massey & Hunter 1998). IX SEA 2010, 14/SEPT/2010 Are there 10 44 M clusters in the MW? Are there 10 M⊙⊙⊙⊙⊙⊙clusters in the MW? Q: Do we have a 30 Dor in the MW? A: Except for NGC3603, we did not know until the advent of IR surveys. In the MW, we know about 10 4 very massive clusters (Mtot > 10 M⊙) →But expect of the order of 80 (Hanson & Popescu, 2008; Ivanov et al., 2010) Figer (2007) IX SEA 2010, 14/SEPT/2010 Cygnus Center Arches Quintuplet Norma RSG2 RSG1 Scutum Crux RSGC3 Westerlund 1 Alicante 8 NGC 3603 Sagitarius Carina Tr14+16 Cyg OB2 Westerlund 2 Orion Perseus Accessible from ORM IX SEA 2010, 14/SEPT/2010 MASGOMAS: en busca de los ccúúúúmulosmulos perdidos MASGOMAS: MAssive Stars in Galactic Obscured Massive clusterS IX SEA 2010, 14/SEPT/2010 TheThe IR IR piercespierces through through Galactic Galactic dust dust Optical (DSS) G61.48+0.09 Marín-Franch et al. (2009) Near-IR (LIRIS@WHT) Near-IR (2MASS) IX SEA 2010, 14/SEPT/2010 MASGOMASMASGOMAS • Select [new] clusters from IR surveys (GLIMPSE, 2MASS, UKIDDS ) • If needed, obtain new IR photometry with LIRIS • Obtain NIR spectra with LIRIS Marín-Franch et al.
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