What have we learned from Eclipsing Binaries in the Orion Nebula? Keivan Guadalupe Stassun Vanderbilt University

Dynamical Masses of PMS Stars circa 2013 N=33 M < 2 M sun Single stars Circumstellar disk “rotation curve” Binary stars Astrometric Eclipsing

Accurate masses and radii: Eclipsing binaries Stassun et al. (2004) V1174 Ori M 1 = 1.01 ± M sun M 2 = 0.73 ± M sun R 1 = 1.34 ± R sun R 2 = 1.07 ± R sun T 1 = 4470 ± 120 K T 2 = 3615 ± 10 K log g 1 = 4.19 ± 0.01 log g 2 = 4.25 ± 0.01 Masses, Radii, and T eff ratio measured very accurately. Luminosities are distance independent.

Comparison of dynamical masses to theoretical models in HR diagram Above 1 Msun: Good agreement: Mean difference 10%, scatter ~10%. Below 1 Msun: Poorer agreement: Differences as large as ~100%, large scatter. Best overall agreement: Siess et al (2000), and Palla & Stahler (1999): Overall consistency to ~5%, though with scatter of ~25%. Updated from Hillenbrand & White (2004)

Case Study: Par 1802 A pair of dissimilar ‘identical twins’ Alicia Aarnio (Michigan) Yilen Gómez (Vanderbilt) Bob Mathieu (Wisconsin) Phillip Cargile (Vanderbilt)

Par 1802: Non-coevality? T 1 /T 2 = ± R 1 = 1.82 ± 0.05 R sun R 2 = 1.69 ± 0.05 R sun L 1 /L 2 = 1.62 ± 0.03 M 1 = 0.39 ± 0.02 M sun M 2 = 0.38 ± 0.02 M sun q = 0.98 ± 0.01 Stassun et al. (Nature, 2008) 1.In protobinary phase, original secondary preferentially accreted mass. 2.Secondary ended accretion phase later. 3.Apparent age difference because different t=0 due to different accretion histories. e.g. Simon & Stobie (2009), Clarke (2007), Bate (2004)

Par 1802: Tidal heating by third body? T 1 /T 2 = ± R 1 = 1.82 ± 0.05 R sun R 2 = 1.69 ± 0.05 R sun L 1 /L 2 = 1.62 ± 0.03 M 1 = 0.39 ± 0.02 M sun M 2 = 0.38 ± 0.02 M sun q = 0.98 ± 0.01 Stassun et al. (Nature, 2008) Gomez et al. (ApJ, 2012)

Triples among known PMS eclipsing binaries 7 triples : 40 doubles = 18% (MS all; Duquennoy & Mayor 1991) 2 triples : 19 doubles = 11% (PMS all; Ghez et al. 1997) 8 triples : 13 doubles = 61% (PMS SBs; Guenther et al. 2007) 8 triples : 6 doubles = 133% (PMS EBs; Stassun & Hebb in prep.)

Case Study: 2M The First Brown-Dwarf Eclipsing Binary Bob Mathieu (Wisconsin) Jeff Valenti (STScI) Yilen Gómez (Vanderbilt) Matthew Richardson (Vanderbilt)

M 1 = 60 ± 3 M Jup M 2 = 39 ± 2 M Jup R 1 = 0.67 ± 0.03 R sun R 2 = 0.51 ± 0.03 R sun Temperature reversal Stassun et al. (Nature, 2006) Stassun et al. (ApJ, 2007) Mohanty et al. (ApJ, 2009) T2/T1 = 1.051

Primary brown dwarf ~10  more magnetically active than secondary. Reiners et al. (ApJ, 2007) Gómez et al. (ApJ, 2009) P 1 = ± days P 2 = ± 0.05 days Primary BD has strong surface activity B ~ 4 kG

Active, primary brown dwarf shifted to cooler T eff, apparently lower mass Mohanty et al. (ApJ, 2009) Stassun et al. (ApJ, 2012) 1 Myr T 2 /T 1 = 1.051± M : Brown Dwarf Eclipsing Binary in Orion Stassun et al. (Nature 2006, ApJ 2007)

Inflated R and Suppressed T eff in Active Low-mass Eclipsing Binaries Lopez-Morales (ApJ, 2007) Stassun et al. (ApJ, 2012) Sample: 11 stars in 7 EB systems Direct, accurate masses, radii, and Teff X-ray luminosity measurements L X to L H  relation of Scholz et al. (2007)

Inflated R and Suppressed T eff in Active Field M-dwarfs Morales et al. (AJ, 2008) Stassun et al. (ApJ, 2012) Sample: PMSU catalog 700 M-dwarfs with distances, spectral types, K mags, H  measures Radii via Stefan-Boltzmann law Masses via BCAH98 tracks

An Empirical Relation Linking R Inflation and T eff Suppression to H  Emission Stassun et al. (ApJ, 2012) Applicability: M < 0.8 M sun -4.6 < log L H  /L bol < -3.3

Primary brown dwarf in 2M effectively corrected in HR diagram 1 Myr Stassun et al. (ApJ, 2012)

New generation PMS models: Incorporating effects of magnetic activity New models incorporating magnetic field effects on stellar structure and evolution can successfully explain R inflation (Feiden & Chaboyer, ApJ 2012) New models show promise when compared to PMS EBs: 2-3x better agreement with dynamical masses (Stassun, Feiden, & Torres, in prep.) Torres et al. (ApJ, 2013)

Summary  Large uncertainties remain in theoretical stellar evolution models. Some models perform better than others in HR diagram, but… PMS is complex: activity effects, accretion histories, third bodies  Magnetic activity alters T eff and R of low-mass stars and brown dwarfs (but L bol appears to be ok). Beware mass and age estimates for active low-mass field objects: Teff or color based mass estimates can be too low by factor of ~2. Empirical relations useful for correcting T eff, R, and M estimates via H .  Apparent non-coevality of ~50% possible in low-mass binaries at ~1 Myr. Accretion history of protobinaries may be important for setting final masses. Tidal heating history and role of third bodies may be important for observed luminosities.  New PMS models incorporating magnetic effects show promise.

Fabienne Bastien (Vanderbilt) Josh Pepper (Lehigh) Gibor Basri (Berkeley)

Measuring stellar structure from “flicker” and “crackle” Bastien et al. (2013, Nature)

Measuring stellar structure from “flicker” and “crackle” Bastien et al. (2013, Nature)

TiOK I

Results for the primary alone: TiO yields Teff = 2500±25 K. At true log g = 3.5, K I implies Teff = 2650±30 K. Mohanty, Stassun & Doppmann (ApJ, 2010) log g = 3.50 T2/T1 = log g = 3.50 T2/T1 = 1.051

Results for the primary alone: Spot model: Teff = 2700 K photosphere with cool spots of 2300 K, covering 70% of surface. Mohanty, Stassun & Doppmann (ApJ, 2010)

Results for the secondary alone Mohanty & Stassun (ApJ, 2012) TiOK I

log g = 3.50 T2/T1 = log g = 3.50 T2/T1 = Results for the secondary alone: Mohanty & Stassun (ApJ, 2012) K I TiO

Implications Theory: Suppression of T eff in primary by magnetic inhibition of convection Observational results: similar spectral discrepancies in both brown dwarfs either spots are not the cause or spots are equally present on both brown dwarfs BOTTOM LINE : Spots cannot be invoked for T eff suppression of primary Chabrier et al. (A&A, 2007)

Global suppression of convection: MacDonald & Mullan (2009) equipartition between magnetic and thermal energies, require interior fields MG BUT: (i) expect equipartition between magnetic and kinetic energies (turbulent, convective) ⇒ fields of only 10s of kG (e.g. Browning 2008, 2010) (ii) field tied to convective fluid ⇒ MG fields buoyantly unstable MacDonald & Mullan (ApJ, 2009) Implications Chabrier et al. (2007) parametrizing in terms of MLT ⇒ α << 0.5 BUT: need 3D simulations of cooling flows along flux tubes in a magnetized convecting medium: Simulations starting to get there (e.g., Browning et al. 2008, 2010; stay tuned…)