Chapter 8 – Host stars PHY6795O – Chapitres Choisis en Astrophysique Naines Brunes et...

Chapter 8 – Host stars

PHY6795O – Chapitres Choisis en Astrophysique

Naines Brunes et Exoplanètes

8. Host stars 2

Contents

8.0 Introduction8.1 Knowledge from astrometry8.2 Photometry and spectroscopy 8.3 Evolutionary models 8.4 Element abundances 8.5 Asterosesimology8.6 Activity and x-ray emission 8.7 Stellar multiplicity8.8 Summary

PHY6795O – Naines brunes et Exoplanètes

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8.0 Introduction Host stars properties are fundamental for deriving the physical

parameters of the planets they host RV technique relies on the knowledge of M★ for inferring MP.

Transit technique relies on the knowledge of R★ for inferring RP.

• The atmospheric signal from a transit light curve scales as 1/R★2. Thus,

exoplanet characterization strongly depends on the knowledge of the host stellar radius.

Lensing technique relies on knowledge of the lens/source distance for inferring MP.

The imaging technique relies on the age to convert (from models) the observed flux into a mass.

Properties of host stars are derived from a combination of astrometric, photometric and spectroscopic observations interpreted within the context of stellar evolutionay models.

Planets are known to exist around a wide variety of stellar types: main sequence stars, from late B to M, BDs, SDBs, giant stars and binary stars.

Host star properties carry information on planet formation mechanisms (e.g. elemental abundance)

PHY6795O – Naines brunes et Exoplanètes 8. Host stars

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8.1 Knowledge from astrometry (1)

Distance and proper motions Distance is a key observable for deriving the host star

properties. Majority of stars hosting planets are within 100 pc and have

relatively good parallaxes from Hipparcos.


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Nearby star census Important for statistical studies and find suitable targets

(e.g. of specific age). Difficult to establish complete census even out to modest

distances of only 10-20 pc. Even though very numerous, low-mass stars are notoriously

difficult to find since too faint. Catalogue of nearby stars

Gliese/(Gliese-Jahreiss): ‘’GJ’’, out to 25 pc (1700 stars) Nstars Database, 2600 objects out to 25 pc. RECONS (Henry 2009); focus on finding missing (low-mass) stars

within 10 pc. 376 stars within 10 pc as of 2012. LSPM (Lepine & Shara Proper Motion): 62 000 entries, 539 new

candidate within 25 pc, 63 within 15 pc.


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Nearby star census Important for statistical studies and find suitable targets

(e.g. of specific age). Difficult to establish complete census even out to modest

distances of only 10-20 pc. Even though very numerous, low-mass stars are notoriously

difficult to find since too faint. Catalogue of nearby stars

Gliese/(Gliese-Jahreiss): ‘’GJ’’, out to 25 pc (1700 stars) Nstars Database, 2600 objects out to 25 pc. RECONS (Henry 2009); focus on finding missing (low-mass) stars

within 10 pc. 376 stars within 10 pc as of 2012. LSPM (Lepine & Shara Proper Motion): 62 000 entries, 539 new

candidate within 25 pc, 63 within 15 pc.


8. Host stars 9

Contents

8.0 Introduction8.1 Knowledge from astrometry8.2 Photometry and spectroscopy 8.3 Evolutionary models 8.4 Element abundances 8.5 Asterosesimology8.6 Activity and x-ray emission 8.7 Stellar multiplicity


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8.2 Photometry and spectroscopy (1)

Bolometric magnitudes (Mbol) Total energy integrated over all wavelengths Distance modulus: Bolometric correction:

Fν is the measured flux and Sν is the filter function.

Bolometric luminosity

MV=4.83, Mbol=4.75 BC given for various filters (e.g. Pecaut & Mamajek 2012)


(8.11)

(8.12)

(8.13)

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8.2 Photometry and spectroscopy (2) Effective temperature

Determination of Teff requires knowledge of both R★ and L★.

R★ determined through angular diameter θ (e.g. CHARA) combined with a distance and/or theoretical evolutionary models.

L★ determined from Mbol which requires a distance.


(8.14)

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Parameters from spectroscopy Requires 2-3m telescopes and high-resolution échelle

spectroscopy (R~50 – 70 000) covering most of the optical spectrum. At high (>100) S/N per resolution element.

Relative line strengths and line profiles enable determination of the following four stellar parameters: effective temperature Teff, surface gravity log g, micro-turbulence velocity and metallicity [Fe/H].

Analysis make use of atmosphere, typically assuming plane parallel geometry (e.g. Kurucz 1993) and LTE conditions.

Typical uncertainties on the parameters: 40-70 K for Teff, 0.02-0.05 dex for [Fe/H]

This calibration is accurate within 47 K over the interval 4500 K < Teff < 6400, -0.85 < [Fe/H] < 0.40, and 0.51 < B-V < 1.20


(8.15)

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Rotational velocity Line profile broadening yields an estimate v sin i Stellar rotation period P estimated from photometric

variability induced by spots. If R★ is known, P yields v at the stellar equator, hence an

estimate i through v sin i . Derived inclination is that of the stellar rotation axis which is

not necessarily align with the orbital inclination. If assumed the same, then sin i estimate lift ambiguity on M sin i for planets detected through RV.

Rotational velocity vs age (gyrochronology) Stellar rotation decreases with time due to angular

momentum loss dictated by stellar winds whose intensity is a function of magnetic activity and rotation. Rotation rate converges at about 1 Gyr.


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8.3 Evolutionary models (1) Evolutionary models constitutes a key theoretical tool

for constraining mass, radii and ages of host stars. Accurate stellar distances are essential in locating the position

of a star in the HR diagram. Models depend on various parameters

• Chemical composition: X (H), Y (He), Z (metal)• Detailed treatment of physical processes

– Line opacity – Convection overshooting (cold sinking material penetrating beyond the

tachocline, the transition between the radiative convective zones).– Element diffusion– Stellar rotation

Evolutionary tracks Table giving mass, Teff, radi, log g vs time for a given initial

mass. Isochrone

HR diagram of an ensemble of stars at a given age


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8.3 Evolutionary models (2) Example of evolutionary tracks – Padova models

PHY6795O – Naines brunes et Exoplanètes 8. Host stars Bertelli et al. (2008)

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8.3 Evolutionary models (3) Example of isochrones – Padova models

PHY6795O – Naines brunes et Exoplanètes 8. Host starsBertelli et al. (2008)

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8.3 Evolutionary models (4) Example of an isochrone table


Bressan et al. (2012)

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8.3 Evolutionary models (5) Evolution of stars, brown dwarfs and giant planets


Burrows et al. (1997)

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8.3 Evolutionary models (6) (Incomplete) list of evolutionary models

Low and massive stars Siess, Dufour & Forestini (2000)

• http://www.astro.ulb.ac.be/~siess/pmwiki/pmwiki.php/Site/WWWTools

Padova (PARSEC models; Bressan et al. 2012)• http://stev.oapd.inaf.it/cgi-bin/cmd

Yontsei-Yale (Demarque et al. 2001)• http://www.astro.yale.edu/demarque/yyiso.html

Darmouth (Dotter et al. 2008; recent models includes magnetic field)• http://stellar.dartmouth.edu/models/index.html

Geneva models (massive stars, rotation; Ekstrom et al. 2012)• http://obswww.unige.ch/Recherche/evol/-Database-

Low-mass, brown dwarfs and planets Lyon models (COND, DUSTY; Baraffe et al. 2003)

• http://perso.ens-lyon.fr/isabelle.baraffe/ Burrows et al. (1997,1998), Saumon & Marley (2008)


http://www.astro.ulb.ac.be/~siess/pmwiki/pmwiki.php/Site/WWWTools

http://www.astro.ulb.ac.be/~siess/pmwiki/pmwiki.php/Site/WWWTools

http://stev.oapd.inaf.it/cgi-bin/cmd

http://stev.oapd.inaf.it/cgi-bin/cmd

http://www.astro.yale.edu/demarque/yyiso.html

http://www.astro.yale.edu/demarque/yyiso.html

http://stellar.dartmouth.edu/models/index.html

http://stellar.dartmouth.edu/models/index.html

http://obswww.unige.ch/Recherche/evol/-Database-

http://obswww.unige.ch/Recherche/evol/-Database-

http://perso.ens-lyon.fr/isabelle.baraffe/

http://perso.ens-lyon.fr/isabelle.baraffe/

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8.3 Evolutionary models (7) (Commonly used) atmosphere models

Low and (mostly) massive stars Castelli & Kurucz (2004)

• Wavelength range: 0.1 – 10 microns• Teff range: 3000 – 50 000 K

Low-mass stars, brown dwarfs and giants planets Pheonix/NextGen group (Hauschildt et al. 1999; Allard et al. 2012)

• 400 < Teff < 70 000 K; -0.5 < log (g) < 5.5; -4.0 <[M/H] < 0.5• Most recent models: BT-Settl, include treatment of clouds.

Morley et al. (2012)• L & T dwarfs with clouds


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8.3 Evolutionary models (8) Age properties of planet host stars

Age determined through: Isochrone fitting Lower limits inferred from Li abundance bs Teff based on

calibrations using open clusters of different ages (Sestito & Randich 2005)

Ca II activity – age relation (Wright et al. 2004) Uncertainties are significant

Age distribution (Saffe et al. 2005) very similar to that of solar neighborhood stars. From Ca II: 5.2±4 Gyr; from isochrones: 7.4±4 Gyr

Velocity dispersion of planet host stars is lower compared to disk stars suggesting that their average age is ~60% that of representative subset of the disk (Reid 2002). Age distribution of nearby FGK stars consistent with a uniform star

formation rate over the 10 Gyr history of the Galactic disk. Most planet host stars are in contrast younger than 5 Gyr.


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8.3 Evolutionary models (9) Li abundance vs age


Sestito & Randich (2005)

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8.3 Evolutionary models (10) Ca II activity (R’HK index) vs age

Chromospheric emission from Ca II is correlated to the stellar rotation period which itself is related to age.


Wright et al. (2004)

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8.4 Elemental abundances (1) Definition of metallicity

Thus [Fe/H] is the abundance of Fe relative to the Sun. The fraction of metals referred to all elements more massive than

He.

Metallicity and Teff not easily measured. Establishing statistical differences between stars w/o planets [Fe/H]

accuracy at the level of 0.1-0.2 dex.

Nearby solar-type stars have [Fe/H]~0 while planet host stars have [Fe/H]~0.15

Planet occurrence appears to be a strong function of [Fe/H].

Extensive literature for abundance analysis.


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8.4 Elemental abundances (2) Compilation of [Fe/H] investigations


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8.4 Elemental abundances (3) Planet occurrence vs metallicity


(8.16)

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8.4 Elemental abundances (4) Planet occurrence vs metallicity

Correlation between planet occurrence and [Fe/H] extends as well to close-in transiting systems.

High [Fe/H] explains anomously large radii measured for some systems (e.g. HD209458b) Radiogenic elements are an internal energy source (~0.5% of the

incoming luminosity) that acts to slow down the cooling of the planet.

To reconcile observed and theoretical radii, ~20 ME is required to explain the [Fe/H]~0 of solar type stars.

Up to ~100 ME to explain stars with [Fe/H] three times solar.

These results add to the big picture that heavy elements play a key role in the formation of close-in giants.

Planet occurrence correlation with [Fe/H] not observed in giants stars.


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8.4 Elemental abundances (5)


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8.4 Elemental abundances (6) M dwarfs

Spectra of M dwarfs are complex due to the presence of several molecules. Difficult to determine [Fe/H] through standard LTE spectroscopic modelling. [Fe/H] is best determined indirectly through photometric calibration (V-Ks)

of more massive bound companions to M dwarfs. Results are inconclusive so far.

Calibration of Bonfils et al (2005) suggest that M dwarf host of planets are metal poor (Bean et al. 2006)

Calibration of Johnson & Apps (2009) suggests that mean [Fe/H] of M dwarf sample with planets is similar to FGK dwarfs with planet. • Consistent pattern of [Fe/H] excess being correlated with planet occurrence.

Planet occurrence around M dwarfs systematically lower compared to FGK dwarfs at any given [Fe/H], likely the result of their lower mass. M dwarf (M★ < 0.6 M) planet occurrence: 3.4+2.2

-0.9 % for Mp sin i > 0.3 MJ, a < 2.45 AU.

Restricted to [Fe/H] > 0.2, occurrence rate rises to 10.7+5.9-4.2 %

Neptune-mass planets do not form preferably around metal-rich stars. Results that rely on small number statistics.


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Spectra of M dwarfs are complex due to the presence of several molecules. Difficult to determine [Fe/H] through standard LTE spectroscopic modelling. [Fe/H] is best determined indirectly through photometric calibration (V-Ks)

of more massive bound companions to M dwarfs. Results are inconclusive so far.

Calibration of Bonfils et al (2005) suggest that M dwarf host of planets are metal poor (Bean et al. 2006)

Calibration of Johnson & Apps (2009) suggests that mean [Fe/H] of M dwarf sample with planets is similar to FGK dwarfs with planet. • Consistent pattern of [Fe/H] excess being correlated with planet occurrence.

Planet occurrence around M dwarfs systematically lower compared to FGK dwarfs at any given [Fe/H], likely the result of their lower mass. M dwarf (M★ < 0.6 M) planet occurrence: 3.4+2.2

-0.9 % for Mp sin i > 0.3 MJ, a < 2.45 AU.

Restricted to [Fe/H] > 0.2, occurrence rate rises to 10.7+5.9-4.2 %

Neptune-mass planets do not form preferably around metal-rich stars. Results that rely on small number statistics.


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8.4 Elemental abundances (9) Origin of the metallicity difference – two hypotheses

Primordial scenario High [Fe/H] is a bulk property of the star and represents the original

composition out of which the proto-planetary system formed Higher [Fe/H] means higher proportion of dust to gas in the protoplanetary

disk. Facilitates condensation and accelerates accretion before the gas disk is lost. Giant planets subsequently formed by runaway accretion of gas onto a rocky

cores of M~10 ME. Cut-off in [Fe/H] distribution for host stars > [Fe/H] > 0.5 represents the upper

limit to metallicities in the solar neighborhood. Probability of forming a giant planet appears to proportional to the square the

number of Fe atoms. Particule collision rate is also proportional to the square of the number of particules. This argues for a physical link between dust particule collisions in the promordial disk and the formation of giant planets.

Detailed core accration models reproduce the distribution of giant planets with [Fe/H]. Also successfully predict that short period giant planets should be rare around M dwarfs but that Neptune mass ice-giants might be more common.


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Self-enrichment of metallicity-enhanced material High [Fe/H] is a phenomenon restricted to the surface region of the star,

arising from capture of metal-rich material, polluting the outer convective envelope.

Might be the result of inward migration of a planet onto the star, slef-enrichment of gas-depleted, metal-rich material from the disk to the satr as a result of migration or breakup and infall of planets and planetesimals.

A planet added to a fully convective star would lead to a negligeable [Fe/H] enhancement but perturbation is significant for solar-type sytars with relatively small convective zones comprising only a few % of the total mass.

Plausible scenario• Such ‘’pollution’’ is directly observed in white Dwarfs.• Differences between photospheric solar and meteoritic abundances could be

explained if the Sun’s convective envelope has been enriched by 0.07 dex compared to its interior.

Prediction: there should be a correlation between planet host metallicity and the convection depth which is not observed, especially in subgiants.


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8.4 Elemental abundances (11) High-metallicity arising from pollution


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8.4 Elemental abundances (12)


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Self-enrichment of metallicity-enhanced material High [Fe/H] is a phenomenon restricted to the surface region of the star,

arising from capture of metal-rich material, polluting the outer convective envelope.

Might be the result of inward migration of a planet onto the star, slef-enrichment of gas-depleted, metal-rich material from the disk to the satr as a result of migration or breakup and infall of planets and planetesimals.

A planet added to a fully convective star would lead to a negligeable [Fe/H] enhancement but perturbation is significant for solar-type sytars with relatively small convective zones comprising only a few % of the total mass.

Plausible scenario• Such ‘’pollution’’ is directly observed in white Dwarfs.• Differences between photospheric solar and meteoritic abundances could be

explained if the Sun’s convective envelope has been enriched by 0.07 dex compared to its interior.

Prediction: there should be a correlation between planet host metallicity and the convection depth which is not observed, especially in subgiants.

General concensus is that the dominant effect is primordial.


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8.4 Elemental abundances (14) Another hypotheses - Different Galactic origin

Giant planet occurrence-metallicity correlation is a dynamical manifestation related to the migration of stars in the Galactic disk Giant planet formation correlates with Galactocentric distance.

Most metal-rich stars (Fe/H > +0.25 dex) found in the solar neighborhood (including those with planets) are considered to have migrated from the inner disk through ‘radial mixing’.

Given Galactic metallicity gradient (0.07-0.1 dex/kpc), stars with mean [Fe/H] of +0.35 would originate at ~3-5 kpc from the Sun in the direction of the Galactic Center.

Consistent with the origin of the Sun whose birthplace was inferred at a Galactocentric radius of 6.6±0.9 kpc based on its [Fe/H] larger by ~0.17 dex than the average of nearby stars os solar age. Hypothesis consistent with its space motion.


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8.4 Elemental abundances (15) Another hypotheses - Different Galactic origin


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8.4 Elemental abundances (16) Probability of forming a terrestrial planet


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8.4 Elemental abundances (17) Implication for terrestrial planet formation

It is estimated that a [Fe/H at least half solar is required to build a habitable terrestrial planet.

Concentrations of radiogenic isotopes 40K, 238U and 232Th will affect the efficiency of plate tectonics which may be an important recycling process providing feedback to stabilize temperatures.

Relative abundance of Si and Mg affect the mass of the core relative to mantle.

Since all elements and isotopes vary with time and location within the Galaxy as a result of star formation activity, planetary systems with the same metallicity as the Sun will not necessarily form Habitable Earth-like planets.

In general, low-metallicity regions of the Galaxy (halo, thick disk and other thin disk) are not optimal for forming Earth-like planets.


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8.4 Elemental abundances (18) Li abundance

Li provide an important (bu complex) diagnostic of both primordial nucleosynthesis and of (pre-) main sequence evolution.

7Li along with 2H, 3He and 4Hewere produced by the Big Bang. 6Li is produced primarily through spallation (cosmic-ray reactions) with the interstellar gas.

Both 6Li and 7Li are destroyed at T>2-3x106 K. This element can survive only in the outer 2-3% of the stellar mass.

Li abundance reflects complex processes of mixing of matter, diffusion, and angular momentum evolution throughout the star’s evolutionary history.

The Li test uses the presence or absence of Li to distinguish candidate brown dwarfs from low-mass stars.

High Li abundance combined with high chromospheric activity often indicates stellar youth.

Planet hosts are observed to be depleted in Li suggesting that the presence of planets/disks increase the amount of mixing within its host.


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8.4 Elemental abundances (19) Li abundance in planet hosts


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8.5 Asteroseismology (1) Asteroseismology, the study of oscillations, provides

independant constraints on fundamental parameters such as mass, density, radius, age, rotation period and chemical composition.

Done by comparing patterns of observed oscillations with predictions from theoretical predictions. Stellar osciallations are predicted over a wide range of mass and

evolutionary states. Use of asteroseismic HR diagram.

Information of high value since stars with the same externally observable parameters (L★, Teff and [Fe/H]) cab have very different interiors as a result of their evolutionary histories. Asteroseismology is the best way of estimating the age of an individual star.

Most observations are done by photometry, best achieved in space (e.g. COROT, MOST, Kepler,…), but also through radial velocity.


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8.5 Asteroseismology (2)


Applications

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8.5 Asteroseismology (3)


Applications

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8.6 Activity and X-ray emission (1) Magnetic and chromospheric activity

This activity is usually inherent to the star but can also be enhanced/triggered through (magnetic) interaction with the star.

Effects expected in outermost layers of the star (chromosphere, transition region, corona)

The broad, deep atmospheric absorption lines of Ca II H and K provide an important diagnostic of chromospheric emission.

Activity is usually modulated with the stellar rotation rate.


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8.6 Activity and X-ray emission (2) X-ray emission

Strongly related to magnetic field and stellar rotation.

X-ray emission stronger in young stars, especially low-mass stars Important age indicator for young M dwarfs.

Close-in planets could cause x-ray emission variation by 30% but effect is difficult to measure.


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