1)Evolution and SED shape 2)Complementary tools to establish evolution 3)A possible evolutionary sequence Evolutionary stages of high-mass star formation from the observational viewpoint Riccardo Cesaroni INAF-Osservatorio Astrofisico di Arcetri

Evolution from SED more difficult than for low-mass: High-mass stars born embedded  SED changes less prominent than in low-mass stars Formation in clusters  affects theory and observations Large distances  confusion Two important parameters from SED:  L bol from integral of SED (basically for λ ~ λ peak ): major contribution from OB stars  M gas from (sub)mm emission (λ > λ peak ) and T estimate Evolution and SED shape

Star formation  gas converted into (OB) stars  L/M grows  SED shape changes SED shape ↔ L vs M ↔ evolution E.g. Molinari et al. (2008): IRAS colour selected sample of massive YSOs (Palla et al. 1991, Molinari et al. 1996, 1998, 2000) (sub)mm cont. maps of southern sub-sample with SEST (Beltran et al. 2006) plus MSX data Analysis of morphology and SED

Result: 3 source types 1.mm only, no IR  modified black body SED 2.Prominent mm and IR  complex SED 3.IR with marginal mm  complex SED Is this and evolutionary sequence? L vs M plot: Different distributions for different types Distributions fitted with evolutionary model

Type 1 IRAS H II regions 850µm 15µm

Type 2 IRAS H II regions

Type 3 IRAS H II region

Result: 3 source types 1.mm only, no IR  modified black body SED 2.Prominent mm and IR  complex SED 3.IR with marginal mm  complex SED Is this and evolutionary sequence? L vs M plot: Different distributions for different types Distributions fitted with evolutionary model

M gas (M O ) L bol (L O ) =type 1 =type 2 =type 3 * time ZAMS ? Molinari et al. (2008)

Herschel/Hi-GAL (Elia et al. 2010) Robitaille SED grey-body SED ! Complex SED Grey-body SED

What can we learn? Precise estimate of L crucial  Herschel far- IR coverage & angular resolution needed Measurement of both L and M  Herschel wavelength coverage needed Theoretical models for L vs M with clusters needed

Complementary tools H II regions: association and physical properties Masers: detection rates and milli-arcsec images Chemistry: molecular abundance variations Kinematics: presence/structure of bipolar outflows, infall, disks  However… no silver bullet!

H II regions Pin-point OB stars on ZAMS  age lower limit Not in pressure equilibrium: size ~ age ? HC  UC  Comp. Free-free flux  N Lyman  spectral type… but surprises are still possible!  radio excess? Caveats: Confusion with extended radio emission Absorption of ionizing photons by dust Density gradients/infall squeeze/quench H II region (Keto 2002) “Flickering’’ of H II region (Peters et al. 2010) Confusion with thermal jets  good radio spectrum needed

H II regions Pin-point OB stars on ZAMS  age lower limit Not in pressure equilibrium: size ~ age ? HC  UC  Comp. Free-free flux  N Lyman  spectral type… but surprises are still possible!  radio excess? Caveats: Confusion with extended radio emission Absorption of ionizing photons by dust Density gradients/infall squeeze/quench H II region (Keto 2002) “Flickering’’ of H II region (Peters et al. 2010) Confusion with thermal jets  good radio spectrum needed

radio excess? dist. x 2 Herschel Hi-GAL SDP Cesaroni et al. in prep. cluster single star

H II regions Pin-point OB stars on ZAMS  age lower limit Not in pressure equilibrium: size ~ age ? HC  UC  Comp. Free-free flux  N Lyman  spectral type… but surprises are still possible!  radio excess? Caveats: Confusion with extended radio emission Absorption of ionizing photons by dust Density gradients/infall squeeze/quench H II region (Keto 2002) “Flickering’’ of H II region (Peters et al. 2010) Confusion with thermal jets  good radio spectrum needed

Masers Tight association between high-mass star forming regions and maser emission (H 2 O, OH, CH 3 OH class I and II) Good correlation L H2O with L IRAS Good positional association of masers with (hot) molecular cores (Codella et al. 1997, 2010) No CH 3 OH class II maser detected in low-mass YSOs (Minier et al. 2003) High detection rates of H 2 O masers towards UC H II regions (e.g. 67%, Churchwell et al. 1990) Low detection rates (12%) of H 2 O masers in IR-dark clouds (Wang et al. 2006)

Class II methanol Proposed evolutionary sequence (Breen et al. 2010): H2OH2O Type 1 Type 2 Type 3

Type 1 Type 2 Type 3 ATCA H 2 O maser det. rates: Type 1: 13% Type 2: 59% Type 3: 21% Sanchez-Monge et al.

Caveats: Statistics mostly based on low angular resolution surveys Detection rates affected by maser variability (days, months)

Kinematics Three main phenomena in star formation: infall, expansion (=outflows), rotation (=disks) Plausible evolutionary sequence: infall  …  infall+outflow+disk  …  disk Observational evidence of evolution?

Infall Difficult to detect: V free-fall ~ FWHM only on small scales (< 0.1 pc) Evidence from red-shifted (self-)absorption  skewed line profiles  blue asymmetry Asymmetries more common in IR-dark than IR- luminous  infall/outflow decrease with time Inverse P-Cyg profiles in HMCs  infall still active at HMC stage i.e. on early ZAMS

infallexpansion Sample of massive (>100M O ) clumps with dist. < 4 kpc: line asymmetries more common in IR-dark clumps Lopez-Sepulcre et al. (2010)

Infall Difficult to detect: V free-fall ~ FWHM only on small scales (< 0.1 pc) Evidence from red-shifted (self-)absorption  skewed line profiles  blue asymmetry Asymmetries more common in IR-dark than IR- luminous  infall (& outflow) decrease with time Inverse P-Cyg profiles in HMCs  infall still active at HMC stage i.e. on early ZAMS

G T B (K) G G NGC7538 IRS1 CO(2-1) 13 CO(2-1) CN(2-1) Klaassen et al. (2011) Furuya et al. (2011) Beltran et al. (2011) Beltran et al. (2006) Girart et al. (2009) G Infall in HMCs G NH 3 (2,2) 13 CO(3-2)

 Infall dominant in Type 1 (IR-dark) sources  Infall still present in Type 2 (HMC) sources

Outflow Easy to detect: broad wings, large (pc) scales Possible evolutionary sequence (Beuther & Shepherd 2005): loss of collimation. Some evidence from IR imaging (Qiu et al. 2008) and maser proper motions (Torrelles et al. 2003) SiO outflows get fainter with L/M, i.e. with evolution (Lopez-Sepulcre et al. 2011)

Beuther & Shepherd (2005)

Outflow Easy to detect: broad wings, large (pc) scales Possible evolutionary sequence (Beuther & Shepherd 2005): loss of collimation. Some evidence from IR imaging (Qiu et al. 2008) and maser proper motions (Torrelles et al. 2003) SiO outflows get fainter with L/M, i.e. with evolution (Lopez-Sepulcre et al. 2011)

Lòpez-Sepulcre et al. (2011): jet/outflow (i.e. SiO) strength decreases with age (i.e. L bol /M gas ) Note: L SiO /L bol and L bol /M gas are distance independent evolution IRD IRL

Disks Known only around B-type (proto)stars (Cesaroni et al. 2007) Elusive in O-type (proto)stars  Short lived? Observational bias?  ALMA will tell (Cesaroni 2008) Type 3: Disks of a few AU found in young B stars (Bik & Thi 2004)  remnant disks? Type 2: Disk+outflow systems in HMCs  embedded ZAMS phase Type 1: Accretion disk+outflow in an IR-dark core (Fallscheer et al. 2009)  could be difficult to detect (faint, thick envelope)  ALMA needed

IRDC (IR-dark cloud) Fallscheer et al. (2009) CH 3 OH 5000 AU disk model small-scale velocity field large-scale bipolar outflow disk-model velocity field 0.2 pc

Possible evolutionary sequence Pre-stellarProtostellarHMC ZAMSZAMS IR emission darkdark λ<70µmluminous FIRluminous SED grey body (mm-FIR)grey bodymulti-comp. H II region NO(thermal jet) ?  HC  UCCompact  masers NOCH 3 OH, H 2 OallOH infall NO?>10 -3 M O /yr YES (but no accretion?) NO outflow NOstrongfadingNO disk NOYES?YESremnant? TemplateΣ > Σ min Mol160 alias IRAS23385 Mol119 alias IRAS20126 IRAS Type 1 Type 2Type 3

Type 1

24µm Molinari et al. (2008) Herschel

G Herschel Hi-GAL Stramatellos et al. (2011) 8µm Spitzer 70µm160µm 500µm350µm 250µm

Type 2

MSX Herschel images

IRAS Herschel MSX 10 4 L O

IRAS Cesaroni et al. Hofner et al. Sridharan et al. Moscadelli et al. Image: 2µm cont. --- OH maser H 2 O masers 1000 AU Keplerian rotation+infall: M * =10 M O Moscadelli et al. (2010) CH 3 OHH2OH2O 200 AU jet disk+jet disk

Type 3

IRAS H II region

Possible evolutionary sequence Pre-stellarProtostellarHMC ZAMSZAMS IR emission darkdark λ<70µmluminous FIRluminous SED grey body (mm-FIR)grey bodymulti-comp. H II region NO(thermal jet) ?  HC  UCCompact  masers NOCH 3 OH, H 2 OallOH infall NO?>10 -3 M O /yr YES (but no accretion?) NO outflow NOstrongfadingNO disk NOYES?YESremnant? TemplateΣ > Σ min Mol160 alias IRAS23385 Mol119 alias IRAS20126 IRAS Type 1 Type 2Type 3

RMS sample (Hoare et al. 2004, Urquhart et al. 2007) dist. x 3 radio det. radio non-det. Reliable luminosity estimate needed to confirm radio excess in B-type stars Herschel  radio excess?

non-radio radio RMS sample (Hoare et al. 2004, Urquhart et al. 2007)

H 2 O masers Li et al. in prep.

Evolution from SED more difficult than for low-mass: High-mass stars born embedded  SED changes less prominent than in low-mass stars Formation in clusters  affects theory and observations Large distances  confusion Two important parameters from SED:  L bol from integral of SED (basically for λ ~ λ peak ): quite reliable because OB stars dominate cluster  M gas from integral of (sub)mm emission (λ > λ peak ) and T from peak of SED

time λFλλFλ 2M O 20M O Robitaille et al. (2006)

Caveats: Statistics mostly based on low angular resolution surveys Detection rates affected by maser variability (days, months) However, VLBI imaging may be rewarding: Proper motions  3-D velocity field, timescale Parallax  distances up to a few kpc

G Hypercompact H II region Moscadelli et al. (2007) Beltran et al. (2007) 7mm free-free & H 2 O masers 500 AU O9.5 star radio SED

7mm free-free & H 2 O masers 30 km/s G Hypercompact H II region t exp = 50 yr !!! distance independent Beginning of expansion phase? Flickering? 500 AU

Caveats: Infall does not always imply accretion! E.g. G24.78: infall on 0.1 pc, but expansion of H II region on 1000 AU (Beltran et al. 2006, 2008) However, accretion may go on even through H II region! E.g. G10.62: Keto & Wood (2006) H66α recomb. line

Molinari et al. (2008) 3.6 cm (H II regions) 850 µm (protostellar core) 24µm IRAS

Possible evolutionary sequence Pre-stellarProtostellarHMC ZAMSZAMS IR emission darkdark λ<70µmluminous FIRluminous SED grey body (mm-FIR)grey bodymulti-comp. H II region NO(thermal jet) ?  HC  UCCompact  masers NOCH 3 OH, H 2 OallOH deuteration large smallNO infall NO?>10 -3 M O /yr YES (but no accretion?) NO outflow NOstrongfadingNO disk NOYES?YESremnant? TemplateΣ > Σ min Mol160 alias IRAS23385 Mol119 alias IRAS20126 IRAS

MIPS 70µm

10 4 L O = L IRAS

Toroids M > 100 M O R ~ AU L > 10 5 L O  O (proto)stars small t acc /t rot  non-equilibrium, circum- cluster structures Disks M < a few 10 M O R ~ 1000 AU L ~ 10 4 L O  B (proto)stars large t acc /t rot  equilibrium, circumstellar structures disks toroids Beltran et al. (2010)

tidal destruction rotational period photo-evaporation Cesaroni et al. (2007)

Assumptions: HPBW = R disk /4 FWHM line = V rot (R disk ) M disk  M star same in all disks T B > 20 K obs. freq. = 230 GHz 5 hours ON-source spec. res. = 0.2 km/s S/N = 20 edge-on i = 35° circumstellar disks Keplerian

Assumptions: HPBW = R disk /4 FWHM line = V rot (R disk ) M disk  M star same in all disks T B > 20 K obs. freq. = 230 GHz 5 hours ON-source spec. res. = 0.2 km/s S/N = 20 no stars edge-on i = 35°

non-radio radio Molinari et al. (2008) RMS sample (Hoare et al. 2004, Urquhart et al. 2007) M gas (M O ) L bol (L O ) time L bol (L O )

ZAMS M gas (M O ) L bol (L O ) =type 1 =type 2 =type 3 * time H 2 O masers: type 1: 13% type 2: 59% type 3: 21%

Chemistry Abundance variations expected with time due to: Freezing on dust grains (e.g. CO depletion)  Evaporation when stars form Reactions on grains and in gas phase (e.g. deuteration) Changes of gas/dust density and temperature

Evolution and SED shape Proposed criteria to identify evolutionary phases of high-mass star forming regions: IR colours: IRAS, MSX, 2MASS (e.g. Wood & Churchwell 1989, Palla et al. 1991, Sridharan et al. 2002, Lumsden et al. 2002) IR-darkness: MSX, Spitzer (e.g. Rathborne et al. 2006)  SED conveys all infos: Herschel!

Evolution from SED more difficult than for low-mass: High-mass stars born embedded  SED changes less prominent than in low-mass stars Formation in clusters  affects theory and observations Large distances  confusion Two important parameters from SED:  L bol from integral of SED (basically for λ ~ λ peak ): quite reliable because OB stars dominate cluster  M gas from integral of (sub)mm emission (λ > λ peak ) and T from peak of SED

24µm Molinari et al. (2008)

L O << L IRAS L IRAS = L O Herschel