LOOKS MOVES BUT - Microsoft · CMD of NGC 6791 (left-hand panel) and NGC 6819 (right-hand panel). Photometric data are taken from Stetson, Bruntt & Grundahl (2003) and Hole et al

Andrea Miglio, Karsten Brogaard, Rasmus Handberg

SACSTELLAR ASTROPHYSICS CENTRE

SAC

STELLAR ASTROPHYSICS CENTRE

LogodesignStellar Astrophysics Centre

LOGOVARIANTER

Aarhus University, Denmark

and

LOOKS LIKE A DUCK, MOVES LIKE A DUCK, BUT DOES IT QUACK LIKE A DUCK?

ASTEROSEISMOLOGY OF RED-GIANT STARS IN CLUSTERS

School of Physics and Astronomy, University of Birmingham, UK

4000450050005500600065007000

100

101

102

103

Teff [K]

L/L su

n

Kepler KASC starsKepler Objects of Interest

Chaplin & Miglio, ARAA, 2013

0

1.0e+04

2.0e+04

3.0e+04

4.0e+04

KIC 6949816

0

2.0e+04

4.0e+04

6.0e+04KIC 9269772

0

5.0e+04

1.0e+05

1.5e+05

2.0e+05

KIC 3100193

10 20 30 40 50 60 70 80 90 100 110 0

5.0e+04

1.0e+05

1.5e+05

2.0e+05

KIC 7522297

Frequency [µHz]Po

wer

spe

ctra

l den

sity

[p

pm2 /µ

Hz]

νmax

5

10

15

20

KIC 8006161

20

40

60

KIC 12069424

100

200

300

KIC 6442183

Pow

er s

pect

ral d

ensi

ty

[ppm

2 /µH

z]

200

400

600

KIC 12508433

500 1000 1500 2000 2500 3000 3500 40000

200

400

600

800

KIC 6035199

Frequency [µHz]

νmax

SOLAR-LIKE OSCILLATIONS

Brown et al. 1991Kjeldsen&Bedding 1995

Asteroseismology of old open clusters 2079

0.8 0.9 1 1.1 1.2 1.3 1.4 1.5 1.6

13

14

15

16

17

18

19

20

V

B−V

NGC6791

Stetson et al. photometryRGB stars used in this studyRC stars used in this study

0.4 0.6 0.8 1 1.2 1.4

12

12.5

13

13.5

14

14.5

15

15.5

16

16.5

V

B−V

NGC6819

Hole et al. photometryRGB stars used in this studyRC stars used in this study

Figure 1. CMD of NGC 6791 (left-hand panel) and NGC 6819 (right-hand panel). Photometric data are taken from Stetson, Bruntt & Grundahl (2003) andHole et al. (2009), respectively. RGB stars used in this work are marked by open squares and RC stars by open circles. See Section 3.2 for a description of thetarget selection.

et al. 1991; Kjeldsen & Bedding 1995; Mosser et al. 2010; Belkacemet al. 2011), and therefore

νmax ≃ M/M⊙(R/R⊙)2

!Teff/Teff,⊙

νmax,⊙ , (2)

where νmax,⊙ = 3100 µHz and Teff,⊙ = 5777 K.These scaling relations are widely used to estimate masses and

radii of red giants (see e.g. Stello et al. 2008; Kallinger et al. 2010;Mosser et al. 2010), but they are based on simplifying assumptionsthat remain to be independently verified. Recent advances have beenmade on providing a theoretical basis for the relation between theacoustic cut-off frequency and νmax (Belkacem et al. 2011), andpreliminary investigations with stellar models (Stello et al. 2009)indicate that the scaling relations hold to within ∼3 per cent on themain sequence and RGB (see also White et al. 2011).

Depending on the observational constraints available, we mayderive mass estimates from equations (1) and (2) alone, or via theircombination with other available information from non-seismic ob-servations. When no information on distance/luminosity is avail-able, which is usually the case for field stars, equations (1) and (2)may be solved to derive M and R (see e.g. Kallinger et al. 2010;Mosser et al. 2010):

M

M⊙≃

"νmax

νmax,⊙

#3 ""ν

"ν⊙

#−4 "Teff

Teff,⊙

#3/2

, (3)

R

R⊙≃

"νmax

νmax,⊙

# ""ν

"ν⊙

#−2 "Teff

Teff,⊙

#1/2

. (4)

In the case of clusters, we can use the distance/luminosities esti-mated with independent methods (i.e. via isochrone fitting or eclips-ing binaries) as an additional constraint. Including this informationallows M to be estimated also from equation (1) or equation (2)

alone (see equations 5 and 6, respectively), or in combination lead-ing to a mass estimate with no explicit dependence on Teff (as inequation 7):

M

M⊙≃

""ν

"ν⊙

#2 "L

L⊙

#3/2 "Teff

Teff,⊙

#−6

, (5)

M

M⊙≃

"νmax

νmax,⊙

# "L

L⊙

# "Teff

Teff,⊙

#−7/2

, (6)

M

M⊙≃

"νmax

νmax,⊙

#12/5 ""ν

"ν⊙

#−14/5 "L

L⊙

#3/10

. (7)

In the following sections, we use equations (3)–(7) directly toestimate M (and R) without adding any extra dependence on stellarmodels. As illustrated in detail e.g. by Gai et al. (2011), additional(so-called ‘grid-based’) methods to estimate M and R can be de-signed. These procedures are also based on equations (1) and (2)but, by searching solutions for M and R in grids of evolutionarytracks, have the advantage of reducing uncertainties on the derivedmass and radius, at the price of some model dependence that weprefer to avoid in this study.

2.1 Error estimates

The formal uncertainties (σ i) on the masses (Mi) of the stars wereused to compute a weighted average mass for stars belonging to theRGB and for stars in the RC:

M =$N

1 Mi/σ2i$N

1 1/σ 2i

.

C⃝ 2011 The Authors, MNRAS 419, 2077–2088Monthly Notices of the Royal Astronomical Society C⃝ 2011 RAS

Belkacem et al. 2011

2000 2200 2400 2600 2800 3000 3200 3400 3600 3800 40000

0.02

0.04

0.06

0.08

0.1

0.12

ν [µHz]

m2 s

−2µ

Hz−

1

2000 2200 2400 2600 2800 3000 3200 3400 3600 3800 40000

0.02

0.04

0.06

0.08

0.1

0.12

ν [µHz]

m2 s

−2µ

Hz−

1

BiSON data

�⌫ =⇣2R R0

drc(r)

⌘�1

/�M/R3

�1/2

Davies et al. 2014

SOLAR-LIKE OSCILLATIONS

NON-RADIAL MODES IN K GIANTS

He core

H-burning shell

H-rich radiative core

convective envelope

gravity mode acoustic mode

An asteroseismic membership study of NGC 6791, NGC 6819, and NGC 6811 3

FIG. 2.— Color-magnitude diagrams of the clusters. Photometry is fromStetson et al. (2003) (NGC 6791), Hole et al. (2009) (NGC 6819), and theWebda database (http://www.univie.ac.at/webda/) (NGC 6811). Representa-tive isochrones from Pietrinferni et al. (2004) (NGC 6791 and NGC 6811) andMarigo et al. (2008) (NGC 6819) are matched to the red giant stars to guidethe eye. Near horizontal dashed lines mark the sampling limit for solar-likeoscillations with Kepler’s long-cadence mode. Large black dots show starsfor which Kepler data are currently available, while large colored dots (red,purple and blue) indicate the subset of these which show evidence of solar-like oscillations. All large dots in the two lower panels show likely membersfrom radial velocity surveys (Hole et al. 2009 and unpublished work by Mei-bom). Vertical dashed lines mark the approximate location of the classicalinstability strip. Stars marked as ’Clump star’ in Tables 1 and 2 are encircled.

The photometric time series data presented here were ob-tained in ’long cadence’ (∆t ∼ 30min, Jenkins et al. (2010a))between 2009 May 12 and 2010 March 20, known as ob-serving quarters 1–4 (Q1–Q4). Within this period the space-craft’s long-cadence mode provided approximately 14,000data points per star. The raw images were processed by thestandard Kepler Science Pipeline and included steps to re-move signatures in the data from sources such as pointingdrifts, focus changes, and thermal variations all performedduring the Pre-search Data Conditioning (PDC) procedure(Jenkins et al. 2010b). PDC also corrects for flux from neigh-boring stars within each photometric aperture based on astatic aperture model. However, this model is not adequatefor all stars, due to small changes in the telescope point-spread-function and pointing between subsequent quarterlyrolls when the spacecraft is rotated 90 degrees to align itssolar panels. As a result the light curves show jumps in theaverage flux level from one quarter to the next. To correctfor that we shifted the average flux levels for each quarterto match that of the raw (pre-PDC) data before stitching to-gether the time series from all four quarters. This ensured thatthe relative flux variations were consistent from one quarterto the next. We compared our corrected (post-PDC) data withthe raw data and also after we performed a number of “man-ual” corrections based entirely on the appearance of the lightcurves (hence not taking auxiliary house-keeping data such aspointing into account). These corrections included removalof outliers, jumps, and slow trends in a similar way as the ap-proach by Garcia et al. (2011). The comparison revealed thatfor a few stars PDC did not perform well, in which case wechose the raw or “manually” corrected raw data.

4. BLENDING AND LIGHT CURVE VERIFICATIONThe super stamps in Figure 1 clearly illustrate that blend-

ing is an issue we need to address before proceeding withthe analysis of these cluster data. Some stars show clearsignatures of blending arising from the relatively large pixelscale (∼ 4′′) of the Kepler photometer compared to the fairlycrowded cluster fields. Blending will give rise to additionallight in the photometric aperture, which will reduce the rela-tive stellar variability, and increase the photon counting noise.In severe cases, the detected stellar variability arises from ablending star and not the target.We have studied the effects from blending by looking at

correlations between light curves of all the target stars (black,red, purple, and blue dots in Figure 2). The light curve corre-lations show no significant increase unless the stars are withinapproximately five pixels of each other and the blending staris at least as bright as the target (Figure 3). We visuallyassessed the light curves of all stars separated by less thanfive pixels, and identified those that showed clear correlationover extended periods of time as blends. We list the blend-ing stars in Table 1 (column-3), Table 2 (column-4), and Ta-ble 3 (column-4). Also listed here, is the variability type ofblending stars identified from single light curves that clearlyshowed variability from two stars. If the variability includedthe expected seismic signal of the target, under the assumptionthat the target was a cluster member, we interpreted the addi-tional signal as caused by a blend. We see no blending forour NGC 6811 targets. We note that the stars for which wecurrently have light curves are far from all stars in the vicinityof the clusters (see Figure 1). It is therefore likely that ourcorrelation analysis has not revealed all blends. For the twomost crowded clusters (NGC 6791 and NGC 6819) there are













~8 Gyr

~2.5 Gyr

~1 Gyr

acoustic spectrum of a R~10 Rsun star

M~ 1.1 Msun

ν~ 35 μHz

M~ 1.6 Msun

ν~ 55 μHz

M~ 2.5 Msun

ν~ 80 μHz

overmassiveundermassive

SOUNDING K GIANTS IN CLUSTERS

OVERMASSIVE STARS

0.4 0.6 0.8 1 1.2 1.4

12

12.5

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13.5

14

14.5

15

15.5

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NGC6819

Hole et al. photometryRGB stars used in this studyRC stars used in this studyOvermassive stars

Masses ~ 2-3 MSUN

Handberg et al. 2015, in preparationBrogaard et al. 2015

in depth studies of their structuree.g. internal rotational profile,

core structure

compare with expectations from binary evolution

detection

A LI-RICH 1.6 MSUN RGB STAR

Anthony-Twarog et al. 2013

*

Rosvick & VandenBerg 1998

Carlberg et al. 2015

✗ ✗0.7 RC


0.8 0.9 1 1.1 1.2 1.3 1.4 1.5 1.6

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20

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B−V

NGC6791


0.4 0.6 0.8 1 1.2 1.4

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14.5

15

15.5

16

16.5

V

B−V

NGC6819





!Teff/Teff,⊙

νmax,⊙ , (2)




M

M⊙≃

"νmax

νmax,⊙

#3 ""ν

"ν⊙

#−4 "Teff

Teff,⊙

#3/2

, (3)

R

R⊙≃

"νmax

νmax,⊙

# ""ν

"ν⊙

#−2 "Teff

Teff,⊙

#1/2

. (4)



M

M⊙≃

""ν

"ν⊙

#2 "L

L⊙

#3/2 "Teff

Teff,⊙

#−6

, (5)

M

M⊙≃

"νmax

νmax,⊙

# "L

L⊙

# "Teff

Teff,⊙

#−7/2

, (6)

M

M⊙≃

"νmax

νmax,⊙

#12/5 ""ν

"ν⊙

#−14/5 "L

L⊙

#3/10

. (7)


2.1 Error estimates


M =$N

1 Mi/σ2i$N

1 1/σ 2i

.


Stello et al. 2010,2011

Corsaro et al. 2012Miglio et al. 2012

Basu et al. 2011

[Fe/H]~0.3-0.4 [Fe/H]~0

Hekker et al. 2011

EXPLOITING ENSEMBLES

14.4 14.6 14.8 15 15.2 15.4 15.6 15.8 16 16.20.6

0.8

1

1.2

1.4

1.6

1.8Eq. 6

V

M/M

sun

RGB

14.5 14.6 14.7 14.8 14.9V

RC

14.5 14.6 14.7 14.8 14.90.6

0.8

1

1.2

1.4

1.6

1.8

V

M/M

sun

RC correction applied

Mass(RGB) vs Mass(RC)

quantitative estimate of integrated RGB mass loss∆M = 0.09±0.03 (random) ±0.04 (systematic) MSUN

INTEGRATED RGB MASS LOSS

NGC6791

Miglio et al. 2012

2 3 4 5 6 7 8 9 10 11 120

50

100

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350

∆ν

∆Π

1

2 3 4 5 6 7 8 9 10 11 120

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100

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∆ν

∆Π

1

NGC6791 NGC6819bare-schwarzschild1Hp overshooting1Hp penetrative convection

TESTING HECB MODELS

Vrad, et al., in preparation Bossini, et al., in preparationArentoft, et al., in preparationBossini, et al., 2015

Handberg, et al, in preparation

16 18 20 22 24 26 28 30 32 342.8

2.9

3

3.1

3.2

3.3

3.4

3.5

3.6

Frequency (µHz)

∆ν

(µH

z)

KIC 5023732 − ∆νfit = 3.0823 µHz

l=0l=1l=2

25 30 35 40 45 503.6

3.7

3.8

3.9

4

4.1

4.2

4.3

Frequency (µHz)

∆ν

(µH

z)

KIC 5113041 − ∆νfit = 3.9943 µHz

l=0l=1l=2

40 45 50 55 60 65 70 75

4.9

5

5.1

5.2

5.3

5.4

5.5

5.6

5.7

5.8

Frequency (µHz)

∆ν

(µH

z)

KIC 5023889 − ∆νfit = 5.3663 µHz

l=0l=1l=2

Kepler giants in NGC6819

ACOUSTIC GLITCHES IN GIANTS

ANY HOPE FOR SEISMOLOGY IN GCS?

A&A proofs: manuscript no. M4

Fig. 2.

Fig. 5.

initial helium abundance Y = 0.25 (Marino et al. 2008; Villanovaet al. 2012; Malavolta et al. 2014).

Mention multiple populations.

We proceed by adopting the recently determined values ofdistance, extinction and reddening, and mass of RGB stars from

isochrone fitting as described in the literature (M = 0.87 M�,see Malavolta et al. 2014 ). For each star in the sample we de-termine Te↵ using the de-reddened (B � V) colours and the re-lations by Ramírez & Meléndez (2005) and the photosphericradius/luminosity by combining distance, apparent magnitude

Article number, page 2 of 5

M4: a first glimpse

Miglio, Chaplin, Brogaard, et al. in preparation


K2P2 pipeline, Lund et al. 2015

Miglio, Chaplin, Brogaard, et al. in preparation

M4: a first glimpse


0.4 0.6 0.8 1 1.2 1.4 1.6 1.8 2 2.2

11

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17

V−I

V

Cluster membersDetections

0 10 20 30 40 50 600

10

20

30

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ν max Expected [µHz]

ν max

Obs

erve

d [µ

Hz]

0 1 2 3 4 5 6 70

1

2

3

4

5

6

7

∆ν Expected [µHz]

∆ν

Obs

erve

d [µ

Hz]


0.8 0.9 1 1.1 1.2 1.3 1.4 1.5 1.6

13

14

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20

V

B−V

NGC6791


0.4 0.6 0.8 1 1.2 1.4

12

12.5

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13.5

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14.5

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15.5

16

16.5

V

B−V

NGC6819





!Teff/Teff,⊙

νmax,⊙ , (2)




M

M⊙≃

"νmax

νmax,⊙

#3 ""ν

"ν⊙

#−4 "Teff

Teff,⊙

#3/2

, (3)

R

R⊙≃

"νmax

νmax,⊙

# ""ν

"ν⊙

#−2 "Teff

Teff,⊙

#1/2

. (4)



M

M⊙≃

""ν

"ν⊙

#2 "L

L⊙

#3/2 "Teff

Teff,⊙

#−6

, (5)

M

M⊙≃

"νmax

νmax,⊙

# "L

L⊙

# "Teff

Teff,⊙

#−7/2

, (6)

M

M⊙≃

"νmax

νmax,⊙

#12/5 ""ν

"ν⊙

#−14/5 "L

L⊙

#3/10

. (7)


2.1 Error estimates


M =$N

1 Mi/σ2i$N

1 1/σ 2i

.


�⌫ 'q

MR3�⌫�

SUMMARY

asteroseismology of K giants in clusters:

overmassive starsexposing camouflage:

first look at M4 with K2: detection of solar-like oscillations

undermassive Li-rich RC star

mass loss, mass loss dispersion

testing stellar models

age-mass relations

Documents

LOOKS MOVES BUT - Microsoft · CMD of NGC 6791 (left-hand panel) and NGC 6819 (right-hand panel). Photometric data are taken from Stetson, Bruntt & Grundahl (2003) and Hole et al