\Delta\Pi_{\mathrm{1}} in field stars and in old-open clusters

Kepler red-clump stars in the field and in open clusters: constraints on core mixing

D. Bossini, A. Miglio, M. Salaris, M. Vrard, S. Cassisi, B. Mosser,J. Montalbán, L. Girardi, A. Noels, A. Bressan, A. Pietrinferni, J. Tayar
School of Physics and Astronomy, University of Birmingham, Edgbaston, Birmingham B15 2TT, United Kingdom
Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Aarhus, DK
Astrophysics Research Institute, Liverpool John Moores University, 146 Brownlow Hill, Liverpool L3 5RF, UK
Instituto de Astrofísica e Cincias do Espaço, Universidade do Porto, CAUP, Rua das Estrelas, 4150-762 Porto,Portugal
Osservatorio Astronomico di Collurania – INAF, via M. Maggini, 64100 Teramo, Italy
LESIA, Observatoire de Paris, PSL Research University, CNRS, Université Pierre et Marie Curie, Université Paris Diderot,
92195 Meudon, France
Dipartimento di Fisica e Astronomia, Università di Padova, Vicolo dell’Osservatorio 3, I-35122 Padova, Italy
Osservatorio Astronomico di Padova – INAF, Vicolo dell’Osservatorio 5, I-35122 Padova, Italy
Institut d’Astrophysique et Géophysique de l’Université de Liège, Allée du six Août, 17 B-4000 Liege, Belgium
SISSA, via Bonomea 265, I-34136 Trieste, Italy
Department of Astronomy, Ohio State University, 140 W 18th Ave, OH 43210, USA
E-mail: dbossini@bison.ph.bham.ac.uk
Accepted 2017 May 3. Received 2017 May 3; in original form 2017 April 7
Abstract

Convective mixing in Helium-core-burning (HeCB) stars is one of the outstanding issues in stellar modelling. The precise asteroseismic measurements of gravity-modes period spacing () has opened the door to detailed studies of the near-core structure of such stars, which had not been possible before. Here we provide stringent tests of various core-mixing scenarios against the largely unbiased population of red-clump stars belonging to the old open clusters monitored by Kepler, and by coupling the updated precise inference on  in thousands field stars with spectroscopic constraints. We find that models with moderate overshooting successfully reproduce the range observed of  in clusters. In particular we show that there is no evidence for the need to extend the size of the adiabatically stratified core, at least at the beginning of the HeCB phase. This conclusion is based primarily on ensemble studies of  as a function of mass and metallicity. While  shows no appreciable dependence on the mass, we have found a clear dependence of  on metallicity, which is also supported by predictions from models.

keywords:
stars: evolution – asteroseismology – stars: low-mass – stars: interiors
pubyear: 2017pagerange: Kepler red-clump stars in the field and in open clusters: constraints on core mixingKepler red-clump stars in the field and in open clusters: constraints on core mixing

1 Introduction

Modelling Helium-core-burning (HeCB) low-mass stars has proven to be complicated, given the lack of a detailed physical understanding of how energy and chemical elements are transported in regions adjacent to convectively unstable cores. In particular, this phase is characterised by convective cores that tend to grow with evolution (hence generating sharp chemical profiles), and by the insurgence of a convectively unstable region separated from the core (called Helium-semiconvection, Castellani et al. 1971). Overall, these uncertainties limit our ability to determine precise stellar properties (such as, e.g, mass and age), which are necessary in the context of studying stellar populations. Moreover, they generate uncertainties in model evolutionary tracks which affect a wide range of applications, including the theoretical calibration of red clump stars as distance indicators and the reliability of theoretical predictions about the following evolutionary stages such as the AGB and the WD ones (see e.g. Girardi, 2016).

It has been recently recognised that the gravity-mode period spacing () measured in solar-like oscillating stars provides stringent constraints on core mixing processes in the Helium-core burning phase (Montalbán et al., 2013; Bossini et al., 2015; Constantino et al., 2015). In our previous work (Bossini et al., 2015, hereafter B15), we investigated how key observational tracers of the near-core properties of HeCB stars (the luminosity of the AGB bump and, primarily, ) depend on the core-mixing scheme adopted. By comparison with data from Pinsonneault et al. (2014) and Mosser et al. (2012) we concluded, in agreement with independent studies (Constantino et al., 2015), that no standard model can satisfactorily account for the period spacing of HeCB stars. We then proposed a parametrised model (a moderate penetrative convection, i.e. overshooting with adiabatic stratification in the extra-mixed region, see Sec. 2) that is able to reproduce at the same time the observed distribution of   and the . However, we were prevented from drawing any further quantitative conclusions because of the inherent limitation of comparing model predictions against a composite stellar population of less than stars, and of potential biases affecting the measurement and detectability of the period spacing (as flagged, e.g. in Constantino et al. 2015).

Here, we specifically address these concerns by studying  of red-clump (RC) stars in old-open clusters, and by investigating the occurrence of any significant detection bias (Sec. 3). Moreover, we take a further step and compare our predictions to the more stringent tests provided by analysing the period-spacing structure (Mosser et al., 2014; Vrard et al., 2016) coupled to spectroscopic constraints, which are now available for thousands of solar-like oscillating giants (SDSS Collaboration et al., 2016), which allows investigating trends of   with mass and metallicity (Sec. 4).

2 Models

We use the stellar evolution code MESA (Paxton et al., 2013) to compute internal structures of stars during the helium-core-burning phase. The default set of physical inputs is described in Rodrigues et al. (2017). We test several types of parametrised mixing schemes during the HeCB phase, which are classified based on the thermal stratification adopted in the region mixed beyond the Schwarzschild border, following the terminology introduced by Zahn (1991). With the term overshooting (OV) we refer to models in which the gradient of temperature in such region is radiative () while penetrative convection (PC) indicates the cases where we assume an adiabatic gradient (). The size of the extra-mixed region is parametrised as , where  is the overshooting parameter and is the minimum between one pressure scale height, , and the radius of the convective core, .111This differs from the default parametrisation of overshooting in MESA, where is instead considered equal to the minimum between  and the mixing length (see Deheuvels et al. 2016) The mixing schemes tested in this work are:

  • MOV: , (Moderate OV),

  • MPC: , (Moderate PC),

  • HOV: , (High OV),

  • HPC: , (High PC).

In B15 we tested several of these schemes, and concluded that only MPC (computed using PARSEC, Bressan et al. 2015) was compatible with the observed  and the luminosity distribution in the early AGB. Regarding the large extra-mixing schemes () we found that only HOV had a good agreement with the observed . However, HOV fails to describe the luminosity distribution (too high ). Finally, the plausibility of a bare-Schwarzschild scheme, here not included, which had already been ruled out theoretically (Gabriel et al., 2014), was also rejected by comparison with observations, including star counts in globular clusters (Cassisi et al., 2003) and . Compared to B15, we have modified the mixing scheme in models that develop He-semiconvection zones (MOV, MPC). We prevent the overshooting region to be (suddenly) attached to the He-semiconvection zone (which in MESA is treated as a convective region), by redefining  to the minimum of . While this should be considered as an ad-hoc treatment with limited physical significance, it provides a stable numerical scheme and mimics an efficient mixing in the He-semiconvective region. For further details see Bossini (2016).

3 Old-open clusters

Differently from field stars, open clusters are simple stellar populations, i.e coeval stars with the same initial chemical composition, and a similar mass for their evolved stars. We can therefore perform a stringent test for the proposed mixing schemes in samples free of selection biases due to age, mass, and metallicity.

Our observational constraints on RC stars in the old-open clusters NGC6791 and NGC6819 are taken from Vrard et al. (2016) and, crucially, include measurements of the gravity-mode period spacing. Among all the stars observed by Kepler in NGC6791 and NGC6819, we exclude those not belonging to the RC ( s), stars that are likely to be product of non-single evolution (over- / under-massive stars, Handberg et al., submitted), and stars in which, according to Vrard et al. (2016), the SNR is too low for a robust inference on  (5 stars in NGC6819, and 3 in NGC6791). The final sample of RC stars in NGC6819 and NGC6791 consists of, respectively, 14 and 16 stars (see Table A.1 and A.2 for a complete list of targets). To compare data with theoretical predictions, we compute models representative of stars in the RC of the two clusters, adopting different extra-mixing schemes described in section 2. For NGC 6791 we calculate an evolutionary track with M, , and (Brogaard et al., 2011), while for NGC 6819 we consider models with M, , and (solar metallicity, Handberg et al., submitted). Figure 1 shows the period spacings of the final samples of the observed RC stars as a function of their average large frequency separation () and the comparison with our model predictions.  in the models is computed from individual radial-mode frequencies (see Rodrigues et al. 2017). As in B15, the high-penetrative-convection scheme predicts a range of  which is too high compared to the observations. The high-overshooting scheme, on the other hand, provides a range which is compatible with the observations, however, it predicts too high a luminosity of the AGB bump. We note that models computed with the OV scheme have an upper limit of  (during the HeCB phase) which does not monotonically increase with   but rather has a maximum at . For higher values of  the He-semiconvection zone, that develops in the late phases of HeCB, remains separate from the convective core, allowing a larger radiative region, thus effectively decreasing . This is the reason why in Fig. 1 MOV reaches higher  values than HOV.

The comparison between models and stars in NGC6791 and NGC6891 supports the conclusions reached in B15, i.e. that a moderate extra-mixed scheme reproduces well the maximum  in the HeCB phase. However, while in NGC6891 the MPC model cannot reach the small period spacings of two stars (21 and 43), that are likely to be early HeCB stars, the moderate overshooting scheme (MOV, green line in Fig. 1) provides a better representation of the data. This model starts the HeCB with a lower , since the overshooting region is radiative, and reaches  as high as the MPC in the late HeCB, since the overall mixed core has in both MOV and MPC schemes. On the other hand, NGC6791 does not present early HeCB stars with small  as predicted by the MOV scheme. A possible cause for this may be ascribed to the limited number of stars in the cluster, or to the three RC stars for which  cannot be determined (see Table 2). The tentative evidence for such a discrepancy is supported by the general trend of the lower limit of  with metallicity in field stars (see Sec. 4 and Fig. 2). However, our sample around the cluster metallicity does not contain a sufficient number of stars to draw strong conclusions.

Offsets in  between models and observations may be attributed to either small differences in the reference mass, or to systematic shifts in the effective temperature scale (due to e.g. uncertainties related to near-surface convection and to outer boundary condition), which modify the predicted photospheric radius, and hence  (). We notice that, for the models used in this study, e.g. increasing the solar-calibrated mixing-length parameter by is sufficient to recover a good agreement (see Fig. 1). We stress that changing the outer boundary conditions / mixing length parameter has no impact on the predictions related to , which is determined by the near-core properties.

Figure 1: Red-clump stars in NGC6791 and NGC6819 in a - diagram. Model predictions based on different near-core mixing schemes are shown by solid lines. The vertical lines indicate the range of  covered by each mixing scheme. The dashed line shows MOV models computed with slightly () increased mixing-length parameters (compared to the solar-calibrated value). The red triangles mark the presence of buoyancy glitches which, to the first order, induce a modulation in  with respect to the asymptotic value (Miglio et al. 2008, Cunha et al. 2015, and Fig. 10 of Mosser et al. 2015), potentially hampering an accurate inference of the asymptotic .

4  of Field stars

Figure 2: Period spacing of HeCB stars with APOGEE DR13 spectroscopic parameters crossed with Vrard et al. (2016) plotted against mass (upper panel) and metallicity (lower panel). Black lines correspond to the 95th and 5th percentiles of the data distribution along , while green and orange lines represent the model predictions (respectively MOV and MPC schemes) for  and . An indication of the typical error on the data is visible in the top-right corner of each panel. NGC6791 (grey dots) and NGC6819 (yellow dots) cluster stars are also shown.

In this section we explore the effects of mass and metallicity on the asymptotic period spacing of stars in the HeCB phase. The dataset we use contains field stars with spectroscopic constrains available from APOGEE DR13 (SDSS Collaboration et al., 2016) and  reported in Vrard et al. (2016). RC stars are selected looking for  greater than s. The range of metallicity considered is . We limit the mass range to M in order to avoid stellar masses that are approaching the secondary clump condition (e.g., see Girardi, 1999). Figure 2 shows the  of the final selection plotted against the mass (upper panel) and metallicity (lower panel). It can be noticed that, in the interval considered, the period spacing is limited in a band between a maximum () and a minimum () value. To measure robustly the observed values of  and  we bin the dataset in mass and metallicity and for each bin we determine the 95th and 5th percentiles of the  distribution (representing  and , respectively). In order to evaluate the uncertainties on the percentiles, taking in account also uncertainties on ,   and [Fe/H], we create 1000 realisations of the observed population. We use these to calculate means and standard deviations of  and   which we then compare to model predictions (see the black lines in Fig. 2).

As evinced from the upper panel of Fig. 2, the data show that the range of  is largely independent of mass, while its upper and lower boundaries decrease with increasing metallicity. To investigate whether models can account for such a behaviour, we compute a small grid of tracks that covers the range of mass and metallicity explored. In the lower panel of Fig. 2 we consider models at different metallicities with mass equal to 1.20 M (close to the average mass of the M of the observed distribution), while for the upper panel we fixed the metallicity to (mean observed value is ) and vary the mass.

We notice that models computed with the MOV scheme are in overall good agreement with the observational constraints. , which is determined by the initial size of the adiabatically stratified core, is well reproduced by the MOV scheme, suggesting that models with PC are disfavoured at least in the initial phases of HeCB. , which depends on the core properties of the much more delicate advanced phases of HeCB, is also in good qualitative agreement with the MOV scheme. Interestingly, models also show decreasing  and  as metallicity increases. The offset between the observed  and the low-mass models (upper panel of Fig. 2) originates primarily from a metallicity effect. Metallicity is not uniformly distributed across the range of masses considered, with small-mass stars being older and hence more metal poor, while the models shown in the upper panel are computed at fixed (solar).

Figure 3: Brunt-Väisälä frequency in the model grid presented in Rodrigues et al. (2017) at the start of the helium burning. profiles are shown on the upper panel for models with fixed metallicity (), while in the lower panel for fixed a mass ( M).

To interpret the behaviour of  and   it is worth recalling that the asymptotic period spacing of dipolar modes is related to the inverse of the integral of the Brunt-Väisälä frequency () over the radius (r) in the g-mode propagation cavity (Tassoul 1980):

(1)

The region that mostly influences the integral in Eq. 1 is the radiative region near the centre (due to the dependency on ). Since is typically null in the deep fully convective regions, larger convective cores will lead to larger values of  (Montalbán et al., 2013). Looking at the Brunt-Väisälä frequency at the very beginning of the HeCB phase (Fig. 3), we notice that for a fixed metallicity all the profiles overlap, while visible differences are found by changing the metallicity. The reason behind this has to be searched in the mass of the helium rich core () at the beginning of HeCB, which determines the physical conditions of the central regions.  is similar to the critical mass , which is needed for the plasma to reach temperatures high enough to burn helium and start the helium flash. For stars with masses in our range of interest, the critical mass  is the result of two competing mechanisms along the RGB: the central cooling due to the degeneracy and the hydrogen-burning shell that constantly deposits helium on the core. While the first is independent of metallicy, the second has an efficiency that increases with . The final effect tends to decrease  with increasing (Cassisi & Salaris, 2013). Indeed, the different  and its properties influences the H-burning shell efficiency also in the HeCB phase, leading to high- stars having a more efficient H-burning shell, which contributes more to the total luminosity than in the low- stars. Therefore, in more metal rich stars the contribution of the He-core-burning to the whole energy budget is lower than in metal-poor ones. This occurrence has the consequence that metal-rich stars develop smaller convective core - and hence smaller  - with respect low-Z stars. On the other hand, low-mass stars with the same end up with similar  and therefore similar helium core during the HeCB and similar . However, this is true only for - M (depending on ). Above this value,  starts to decrease since we approach the secondary clump (the degeneracy of the He core decreases and hence , Girardi 1999).

Figure 4: Left panel: RC- evolution as a function of central helium mass of a M solar metallicity star with three different mixing schemes and four different C(,)O reaction rates. Right panel: RC- evolution as a function of central helium mass fraction of a M star computed assuming three different combinations of and .

An additional test is made to quantify the effect of the initial helium on . In our grid is in fact coupled with via linear chemical enrichment relation (see Rodrigues et al., 2017). To decouple the effect of and we compute five tracks of mass and () but with initial helium (with as our default value). Figure 4, right panel, shows the evolution of  with central helium for the 5 tracks. We notice that the effect on  and  grows linearly with . However for variation of from the default value ( and ) the deviation on  () lays within the observed uncertainty on  (- s). The deviation we tested is compatible with typical spread on disk populations (Casagrande et al., 2007). Nevertheless, the effect becomes substantial for extreme enrichments, e.g. in bulge or globular clusters where populations with very different He abundances and same metallicity may coexist (Renzini et al., 2015).

5 Additional uncertainties on : C(,)O nuclear reaction rate

As shown above  is strongly dependent on assumptions related to core convection and metallicity, however, additional parameters and uncertainties may also impact on , in particular the C(,)O reaction rate, which, along with triple-, plays a fundamental role, especially at the end of the HeCB (e.g. see Straniero et al., 2003; Cassisi et al., 2003; Constantino et al., 2015).

We compute a series of HeCB evolutionary tracks ( M, solar abundance) in which we adopt four C(,)O reaction rates in conjunction with 3 mixing schemes: Bare-Schwarzschild model (BS), step function overshooting (HOV), and 1 penetrative convection (HPC). The C(,)O reaction rates considered are the tabulated values given by JINA (Cyburt et al., 2007), K02 (Kunz et al., 2002), CF88 (Caughlan & Fowler, 1988), and NACRE (Angulo et al., 1999) and already made available in MESA. While no difference can be noticed at the beginning of the phase, the impact of the different mixing schemes is evident at the maximum period spacing (end of the HeCB), where HOV tracks show a scatter of around - s between them, compared to only s for BS (Figure 4). We therefore expect an uncertainty between 6 and 2 s on the MPC and MOV models. This value is comparable with the observed average  uncertainty for clump stars ( s).

6 Discussion and Conclusion

The precise measurements of gravity-modes period spacing () in thousands of He-core-burning stars has opened the door to detailed studies of the near core structure of such stars, which had not been possible before (Montalbán et al., 2013; Bossini et al., 2015; Constantino et al., 2015)

Here we provide additional stringent tests of the mixing schemes by stress-testing the models presented in B15 against results on the simple population of RC stars belonging to the old open clusters monitored by Kepler, and making use of the updated precise inference on  presented by Vrard et al. (2016) coupled with spectroscopic constraints from APOGEE DR13 (SDSS Collaboration et al., 2016).

We find that in clusters  is measured in all RC stars with few exceptions (as discussed in Section 3, Table 1 and 2), and that models with moderate overshooting can reproduce the range of period spacing observed. In particular our models do not support the need to extend the size of the adiabatically stratified core, at least at the beginning of the HeCB phase. This conclusion is based primarily on ensemble studies of  as a function of mass and metallicity, where we could also show that models successfully reproduce the main trends (or their absence). While  shows no appreciable dependence on the mass, we have found a clear dependence of  on metallicity (Figure 2) also shown by the models, which strengthens even further the result on the clusters. We complement the study by considering how theoretically predicted  depends on the initial helium mass fraction and on the nuclear cross sections adopted in the models, and conclude that the adopted mixing scheme and metallicity are the dominant effects.

The parametrised model presented here should be considered as a first approximation which broadly reproduces the inferred average asymptotic period spacing. Significant improvements can be made by looking at signatures of sharp-structure variations (Mosser et al., 2015; Cunha et al., 2015; Bossini et al., 2015) which will enable to test in greater detail the chemical and temperature stratification near the edge of the convective core, providing additional indications that will eventually be compared with more realistic and physically justified models of convection in this key stellar evolutionary phases.

Acknowledgments

AM acknowledges the support of the UK Science and Technology Facilities Council (STFC). We acknowledge the support from the PRIN INAF 2014 – CRA 1.05.01.94.05. JM acknowledges support from the ERC Consolidator Grant funding scheme (project STARKEY, G.A. n. 615604). MV acknowledges funding by the Portuguese Science foundation through the grant with reference CIAAUP 03/2016BPD, in the context of the project FIS/04434, cofunded by FEDER through the program COMPLETE. SC acknowledges financial support from PRIN-INAF2014 (PI: S. Cassisi)

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N KIC Vrard et al. class selected notes
(Hz) (Hz) (s) (s) (2016) RC stars
1 4937011 RGB
2 4937056 46.10 4.76 yes unclear low l=1 structure
3 4937257 RGB
4 4937576 32.95 3.56 yes RGB
5 4937770 94.30 7.83 160.70 1.99 yes unclear Possible second clump.
6 4937775 89.92 7.33 226.50 4.61 yes unclear little l=1 structure. No photometric clump
7 5023732 27.08 3.12 yes RGB
8 5023845 109.94 8.96 yes RGB
9 5023889 RGB
10 5023931 50.57 4.92 yes RGB little l=1 structure
11 5023953 48.75 4.74 299.00 4.36 yes RC no
12 5024043 55.98 5.64 61.00 2.00 yes RGB
13 5024143 122.84 9.68 68.40 6.59 yes RGB
14 5024240 153.85 12.00 yes RGB
15 5024268 RGB
16 5024272 RGB
17 5024297 46.24 4.60 yes RGB
18 5024312 96.68 8.13 yes RGB
19 5024327 44.18 4.72 269.50 3.21 yes RC yes
20 5024329 RGB
21 5024404 47.09 4.78 242.00 3.54 yes RC yes
22 5024405 98.89 8.29 yes RGB
23 5024414 78.17 6.46 280.70 6.16 yes RC no
24 5024456 3.86 0.70 yes RGB
25 5024476 65.94 5.74 298.00 3.79 yes RC no
26 5024512 72.97 6.70 yes RGB
27 5024517 50.13 4.94 319.20 5.11 yes RGB non-photometric member
28 5024582 46.30 4.82 323.50 4.76 yes RC yes
29 5024583 37.89 3.91 yes RGB
30 5024601 32.30 3.68 yes RC Very low l=1
31 5024750 12.74 1.80 yes RGB
32 5024851 4.06 0.75 yes RGB
33 5024870 RGB
34 5024967 44.97 4.71 yes RC Very low l=1
35 5024984 RGB
36 5111718 135.49 10.59 87.53 1.06 yes RGB
37 5111820 RGB
38 5111940 52.79 5.20 yes RGB
39 5111949 47.35 4.81 317.00 4.76 yes RC yes
40 5112072 125.27 10.08 92.40 0.91 yes RGB
41 5112288 46.94 4.77 yes RC no Very low l=1
42 5112361 67.63 6.19 91.40 0.90 yes RGB
43 5112373 43.84 4.61 239.60 2.57 yes RC yes
44 5112387 44.91 4.70 267.20 3.21 yes RC yes
45 5112401 36.50 4.05 311.00 9.00 yes RC yes presence of glitch
46 5112403 141.73 11.18 86.80 1.07 yes RGB
47 5112467 45.15 4.75 285.20 3.67 yes RC yes
48 5112481 5.18 0.92 yes RGB
49 5112491 44.25 4.68 324.20 4.65 yes RC yes
50 5112558 RGB
51 5112730 43.60 4.56 320.00 4.46 yes RC yes
52 5112734 40.65 4.16 yes RGB
53 5112741 RGB
54 5112744 43.97 4.44 yes RGB
55 5112751 1.32 0.39 yes RC no Very low l=1
56 5112786 7.70 1.17 yes RGB
57 5112880 25.43 2.82 yes RGB
58 5112938 44.54 4.73 321.00 4.59 yes RC yes
59 5112948 42.28 4.31 yes RGB
60 5112950 41.59 4.35 319.50 4.26 yes RC yes
61 5112974 40.08 4.32 310.60 3.87 yes RC yes presence of glitch
62 5113041 37.13 4.01 yes RGB
63 5113061 4.53 0.84 yes RGB
64 5113441 154.68 11.76 89.65 1.24 yes RGB
65 5199859 0.70 0.15 yes RGB
66 5200088 RGB
67 5200152 45.71 4.74 327.20 4.89 yes RC yes
Table 1: NGC6819 stellar catalogue. Stellar classes were chosen considering the classification in Stello et al. (2011) and  in Vrard et al. (2016).f
N KIC Vrard et al. class selected notes
(Hz) (Hz) (s) (s) (2016) RC stars
1 2297384 30.49 3.75 313.78 3.00 yes RC yes
2 2297574 RGB
3 2297793 RGB
4 2297825 30.43 3.77 301.10 2.76 yes RC yes
5 2298097 RGB
6 2435987 38.07 4.22 yes RGB
7 2436097 42.06 4.54 yes RGB
8 2436209 57.01 5.76 67.30 2.46 yes RGB
9 2436291 RGB
10 2436332 28.29 3.40 yes RGB
11 2436376 RGB
12 2436417 27.07 3.40 305.00 5.00 yes RC yes
13 2436458 37.08 4.17 yes RGB
14 2436540 57.76 5.83 yes RGB
15 2436593 24.96 3.56 yes RGB binary star
16 2436676 131.86 11.35 80.60 1.16 yes RGB
17 2436688 76.01 7.28 yes RGB
18 2436715 RGB
19 2436732 30.27 3.66 259.76 2.04 yes RC yes
20 2436759 32.63 3.73 yes RGB
21 2436804 RGB
22 2436814 24.51 3.13 yes RGB
23 2436818 97.32 8.84 76.10 0.56 yes RGB
24 2436824 34.03 3.87 yes RGB
25 2436842 RGB
26 2436900 35.62 4.07 yes RGB
27 2436912 29.79 3.73 yes RC no suppressed l=1
28 2436944 30.86 3.72 271.90 2.28 yes RC yes
29 2436954 34.48 4.16 yes RGB
30 2436981 RGB
31 2437033 RGB
32 2437040 25.49 3.08 yes RGB
33 2437103 28.46 3.72 276.00 2.17 yes RC yes presence of glitch
34 2437112 RGB
35 2437171 RGB
36 2437178 RGB
37 2437209 RGB
38 2437240 45.56 4.86 63.60 1.68 yes RGB
39 2437261 RGB
40 2437270 69.36 6.54 62.60 2.75 yes RGB
41 2437296 RGB
42 2437325 93.60 8.54 75.30 0.53 yes RGB
43 2437340 8.44 1.30 yes RGB
44 2437353 31.25 3.80 297.00 2.73 yes RC yes
45 2437394 159.70 12.99 yes RGB
46 2437402 46.41 4.84 yes RGB
47 2437413 RGB
48 2437443 RGB
49 2437444 18.83 2.48 yes RGB
50 2437488 64.77 6.30 yes RGB
51 2437496 4.43 0.86 yes RGB
52 2437507 20.41 2.62 yes RGB
53 2437564 32.02 3.82 292.10 2.69 yes RC yes presence of glitch
54 2437589 46.11 4.63 yes unclear Uncertain . probably RGB.
55 2437624 RGB
56 2437630 RGB
57 2437653 74.58 7.07 yes RGB
58 2437698 29.78 3.70 287.50 6.01 yes RC yes
59 2437781 85.46 7.88 yes RGB
60 2437792 RGB
61 2437804 26.71 3.34 247.30 1.63 yes RC yes
62 2437805 31.93 3.76 261.50 2.16 yes RC yes
63 2437816 17.78 2.37 yes RGB
64 2437851 12.37 1.90 yes RGB
65 2437897 RGB
66 2437930 RGB
67 2437933 107.46 9.45 76.80 0.63 yes RGB
68 2437957 93.15 8.57 yes RGB
69 2437965 7.20 1.28 yes RGB
70 2437972 85.43 7.89 69.10 0.41 yes RGB
71 2437976 89.62 8.21 75.50 0.51 yes RGB
72 2437987 29.95 3.72 278.00 15.00 yes RC yes presence of glitch
73 2438038 62.25 6.18 66.40 2.65 yes RGB
74 2438051 30.46 3.66 250.20 1.93 yes RC yes presence of glitch
75 2438053 RGB
76 2438140 70.96 6.79 67.30 3.20 yes RGB
77 2438192 RGB
78 2438333 61.09 6.11 65.00 2.48 yes RGB
79 2438421 0.67 0.21 yes RGB
80 2568519 RGB
81 2568916 0.45 0.23 yes RC no very poor SNR
82 2569055 30.49 3.69 268.00 2.22 yes RC yes
83 2569126 RGB
84 2569360 21.32 2.76 yes RGB
85 2569488 0.54 0.23 yes RC no very poor SNR
86 2569618 56.41 5.70 yes RGB
87 2569650 RGB
88 2569673 RGB
89 2569712 RGB
90 2569752 RGB
91 2569912 RGB
92 2569935 5.20 0.98 yes RGB
93 2569945 30.13 3.78 297.30 2.66 yes RC yes
94 2570094 67.83 6.55 71.05 0.34 yes RGB
95 2570134 RGB
96 2570144 RGB
97 2570172 75.15 7.09 yes RGB
98 2570214 28.05 3.54 245.20 1.69 yes RC yes presence of glitch
99 2570244 105.70 9.28 76.85 0.62 yes RGB
100 2570263 RGB
101 2570384 47.54 4.84 yes RGB
102 2570518 46.45 4.98 yes RGB
103 2570519 RGB
104 2579142 RGB
105 2582664 RGB
106 2585397 RGB
Table 2: NGC6791 stellar catalogue. Stellar classes were chosen considering the classification in Stello et al. (2011) and  in Vrard et al. (2016).
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