# CHEMICAL COMPOSITIONS OF RED GIANT STARS IN THE OLD OPEN CLUSTER NGC 7789

Neel Nagarajan Department of Astronomy and McDonald Observatory, The
University of Texas, Austin, TX 78712, USA<EMAIL_ADDRESS>Christopher Sneden Department of Astronomy and McDonald Observatory, The
University of Texas, Austin, TX 78712, USA<EMAIL_ADDRESS>Melike
Afşar Department of Astronomy and Space Sciences, Ege University, 35100
Bornova, İzmir, Turkey<EMAIL_ADDRESS>Catherine A. Pilachowski
Department of Astronomy, Indiana University, Indiana University, Bloomington,
IN 47405, USA<EMAIL_ADDRESS>
###### Abstract

We have gathered optical-region spectra, derived model atmosphere parameters,
and computed elemental abundances for 15 red giant stars in the open cluster
NGC 7789. We focus on the light element group CNOLi that provides clues to
evolutionary changes associated with internal fusion events and chemical
mixing. We confirm and extend an early report that NGC 7789 stars 193 and 301
have anomalously large Li abundances, and that these values are apparently
unconnected to any other elements’ abundances in these stars. A companion
study of He i $\lambda$10830 lines in both field stars and cluster members
shows that star 301 has a strong He feature while star 193 does not. Possible
explanations for the large Li abundances of these stars include helium flash-
induced mixing events and binary interactions at some past or present times.
In either case an internal eruption of energy could cause fresh synthesis of
lithium via the Cameron-Fowler Berillyum transport mechanism. Rapid transport
of lithium to the outer layers may have created significant chromospheric
transient disturbances, producing enough helium ionization to allow for the
strong $\lambda$10830 absorption in star 301.

††software: linemake (https://github.com/vmplacco/linemake), MOOG (Sneden
1973), IRAF (Tody 1986, Tody 1993), SPECTRE (Fitzpatrick & Sneden 1987)

##  1 INTRODUCTION

Red giant stars typically exhibit low lithium abundances relative to those of
main sequence stars. During main sequence stars’ lives, Li is easily destroyed
in high temperature interior regions as part of the proton-proton cycle of
hydrogen fusion: 7Li($p,\alpha$)$\rightarrow$4He. As main sequence stars
evolve to become subgiants and then red giants, their deepening convective
envelopes mix surface and interior layers, effectively cleaning the stellar
atmospheres of their natal Li contents.

Because of this lithium-depletion process we typically observe surface Li
abundances to be large in main sequence stars (as much as log $\epsilon$(Li)
$\sim$ 3.3)111 We adopt the standard spectroscopic notation (Wallerstein &
Helfer, 1959) that for elements A and B, [A/B] $\equiv$ log10(NA/NB)⋆ $-$
log10(NA/NB)⊙. We use the definition log $\epsilon$(A) $\equiv$ log10(NA/NH) +
12.0, and equate metallicity with the stellar [Fe/H] value. but much lower in
red giant stars (log $\epsilon$(Li) $\lesssim$ 1.5). However, around one
percent of red giant stars in our Galaxy have unusually high Li abundances, as
evidenced by a strong Li $\lambda$6707 absorption. Many papers have reported
discoveries of Li-rich giants. Lithium excess is defined in various ways, but
most papers suggest that red giants have anomously large Li abundances if log
$\epsilon$(Li) $\gtrsim$ 1.5.

Two main sources have been suggested for such Li overabundances in red giants.
First, it might be possible for Li to be added as an expanding stellar
envelope engulfs a companion, typically thought to be a terrestrial companion
(e.g., Alexander 1967). Second, fresh Li could be produced in the interiors of
a star in the so-called beryllium transport mechanisim (Cameron & Fowler,
1971), in which Li is synthesized via 3He($\alpha,\gamma$)7Be(e-,$\nu$)7Li and
then dredge up to the surface before it can be destroyed in normal pp-chain
hydrogen fusion. The assets and liabilities of these ideas are discussed in
more detail in many papers, e.g., Casey et al. (2019), Deepak & Reddy (2019),
Martell et al. (2021). Recently, Sneden et al. (2022) (hereafter Paper 1)
conducted a large survey of He i 10830 Å chromospheric lines, discovering a
correlation between high Li abundances and very large $\lambda$10830
absorption strengths. Apparently the generation of Li in red giant envelopes
is often accompanied by major disturbances in their outer atmospheres.

Some recent studies have begun to search for Li-rich stars in open clusters,
where masses and ages can be estimated; e.g., Anthony-Twarog et al. (2013),
Carlberg et al. (2016), Magrini et al. (2021), and Sun et al. (2022). In the
present work we revisit the intermediate age open star cluster NGC 7789, which
contains the first two red giants reported to have anomalously large Li
abundances (log $\epsilon$(Li) $\sim$ +2.4, Pilachowski 1986). This cluster is
relatively bright and has an extensive literature history. At age 1.5 Gyr
(Gao, 2018) NGC 7789 has a well-developed red giant branch, and has been
subjected to several high-resolution spectroscopic analyses, but apparently
without focus on its Li abundance anomaly since the Pilachowski study. Basic
data for NGC 7789 are summarized in Table 1.

Pilachowski (1986) reported significant Li abundance excesses in NGC 7789 193
and NGC 7789 301 (hereafter labeled stars 193 and 301). Paper 1 included these
stars and 8 other cluster members from the Pilachowski sample. However, the
optical spectra from the paper over 3 decades ago were of modest spectral
resolution and signal-to-noise ratio (SNR); these stars lack recent accessible
optical high-resolution spectra in the Li i $\lambda$6707 region. Therefore we
have gathered and analyzed new optical echelle spectra for these NGC 7789
stars. In this paper we report Li abundances and discuss the Li-He connection
for NGC 7789 red giants.

In §2 we describe the stellar sample, our observations, and the reduction
steps to produce the final spectra. The model atmosphere and abundance
analyses are outlined in §3. We discuss the results in §4 with particular
emphasis on the LiCNO abundance group and relation to the $\lambda$10830 data
from Paper 1, and summarize our conclusions in §5.

##  2 OBSERVATIONS AND REDUCTIONS

###  2.1 The Stellar Sample

We observed and analyzed 15 NGC 7789 red giants. The program stars are listed
in Table 2 along with their astrometric and photometric properties of interest
to our work. Inspection of the parallax and proper motion data in this table
reveals that star 193 is discordant and likely not a member of NGC 7789. The
possibility that star 193 is a non-member arose in Pilachowski (1986), but the
star was not discarded because its high Li abundance would make the star
interesting regardless of its membership status. We also retained star 193 in
our work, and will comment on it later.

We began with a cluster membership analysis. We considered stars in a search
radius roughly double the radius of the WEBDA cluster chart for NGC 7789. We
gathered the astrometric data for stars in this area from the GAIA (Gaia
Collaboration et al., 2016) data archive Gaia Collaboration et al. (2022). We
retained stars with GAIA parameters parallax $p$ and proper motions
$\mu_{RA}$, $\mu_{Dec}$ in the following ranges respectively: 0.4 to 0.5,
$-$1.2 to $-$0.6, and $-$2.3 to $-$1.65. All these ranges are centered on the
cluster mean values as presented in SIMBAD.

After elimination of non-member stars, we then constructed an extinction-
corrected Gaia-based color magnitude diagram (CMD) for NGC 7789. We converted
$E(B-V)$ from Table 1 to $E(BP-RP)$ = 0.39 using two equations: $E(BP-RP)$
$\simeq$ 4.507$EW_{862}$ (developed by Gaia collaboration et al. 2022), and
$E(B-V)$ $\simeq$ 3.1$EW_{862}$ (from table 3 of Gaia collaboration et al.
2022). Here, $EW_{862}$ is the equivalent width of the diffuse interstellar
band at 862 nm and $E(B-V)$ $\simeq$ 3.1$EW_{862}$ is a relationship between
the band strength and the $UBV$ standard reddening value.

We converted the total extinction to Gaia units via the relation $A_{G}/A_{V}$
$\simeq$ 0.95 from Jordi et al. (2010). Application of these formulae yielded
the $(BP-RP)_{0}$ and $M_{G0}$ values employed to create the CMD shown in
Figure 1. This is similar to other CMDs for NGC 7789 in the literature, e.g.
Overbeek et al. (2015), Gao (2018), and Cantat-Gaudin et al. (2018). We have
included star 193 in this plot even though it is not a probable NGC 7789
member, and we have arbitrarily used the cluster reddening values in placing
it in the figure. Its position in Figure 1 should not be used to infer its
evolutionary history. Furthermore, star 765’s membership may also be in
question due to it’s relatively low parallax (Table 2), but its proper motions
are consistent with NGC 7789 membership. The dereddened CMD presented in
Figure 5 of Sandquist et al. (2020) for multiple open clusters shows that the
approximate red clump center is at [$(BP-RP)$,$M_{G0}$] = [1.1,0.4], which is
essentially coincident to the red clump center for NGC 7789.

In Figure 1 the stars with $(BP-RP)_{0}$ $\lesssim$ 0.4 are easily identified
as blue straggler stars (BSS). Such members of NGC 7789 have been known for
many decades since the pioneering CMD study of Burbidge & Sandage (1958). Blue
stragglers probably result from mass transfer in binary or triple systems
(e.g., McCrea 1964, Hills & Day 1976, Perets & Fabrycky 2009), suggesting that
NGC 7789 has a relatively large set of present or past multiple star systems.

Figure 1:  An extinction-corrected color magnitude diagram (CMD) for NGC 7789
using GAIA photometric data. As identified in the plot legend, the star-shaped
symbols show the red giants of this study. The reddening vector shows the
direction and length by which we shifted the stars on the CMD to correct for
interstellar reddening effects.

###  2.2 High-Resolution Spectra

We gathered high resolution spectra of the target stars with the McDonald
Observatory 2.7m Harlan J. Smith Telescope and Tull Echelle Spectrograph Tull
et al. (1995). The combination of slit, dispersing elements, and detector
yielded spectral resolution of $R~{}\equiv\lambda/\Delta\lambda$ $\simeq$
45,000. The observed spectral range covered 3500$-$9500 Å. However, our red
giants have sharply decreasing fluxes and substantial spectral line blanketing
toward shorter wavelengths, so in consequence our analyses considered only
wavelength regions with $\lambda$ $>$ 5300 Å. Typical signal-to-noise ratios
near 6000 Å were SNR $\sim$ 100 per resolution element. The obtained spectra
included star 971, a red giant tip member of NGC 7789. This star is by far the
reddest one of our sample (Table 2). Its low temperature along with inclement
weather during observation resulted in a very noisy reduced spectrum, and so
we dropped it from further consideration in this paper.

Figure 2:  Spectra of the Li i $\lambda$6707 resonance transition in all NGC
7789 red giants observed with the Tull Spectrograph (left panel) and of the He
i $\lambda$10830 transition in all of this cluster’s stars observed with
HET/HPF by Sneden et al. (2022) (right panel). Attention is called to Li-rich
stars 193 (in blue) and 301 (in red). All spectra of NGC 7789 giants observed
with either instrument are shown in this figure.

All spectra were obtained in a single observing run covering 7 nights. Total
integration times for each program star were 2$-$3 hours, obtained in
individual half-hour exposures. The program stars’ observations were
accompanied by bias, dark, flat-field incandescent, and Th-Ar lamp
integrations. We also observed several rapidly rotating hot stars to
facilitate cancellation of telluric features in the program stars.

###  2.3 Reductions and Equivalent Widths

We performed all spectral reductions using the IRAF222 http://iraf.noao.edu/
(Tody, 1986, 1993) facility. The reduction tasks included bias subtraction,
cosmic ray excision, echelle order extraction, continuum normalization, pixel
to wavelength conversion, and velocity correction to rest wavelength scales.

One of the chief goals of our study is to investigate the relationship between
Li abundance and He i $\lambda$10830 absorption strength. In the left-hand
panel of Figure 2 we show spectra of the Li i $\lambda$6707 resonance feature
in all stars observed here, and in the right-hand panel we show the He i line
in spectra collected for Paper 1.

We measured EWs of unblended atomic lines, beginning with the line lists used
by Böcek Topcu et al. (2015) for their study of open cluster NGC 752 and by
Böcek Topcu et al. (2016) for NGC 6940. We calculated the EWs with the
SPECTRE333 https://www.as.utexas.edu/ chris/spectre.html spectrum analysis
code (Fitzpatrick & Sneden, 1987). The line profiles were usually modeled by
Gaussian functions, with occasional Voigt functions for the strongest lines.
EWs could not be reliably measured for some lines below 5300 Å due to sharply
decreasing fluxes and substantial spectral line blanketing; those lines were
excluded from our analyses. The EWs are listed in Table 3.

##  3 ABUNDANCE ANALYSES

In this section we describe the steps in analyses that led to final abundances
of elements in our NGC 7789 program stars. Our focus is on model atmospheric
parameters and abundances of the LiCNO element group.

###  3.1 Model Atmospheres

We determined model atmospheric parameters and abundances from the line data
information in the NGC 7789 spectra. Initially we attempted to use the
measured EWs to estimate model parameters in a standard manner: (1) $T_{\rm
eff}$, from requiring low and high excitation Fe i lines to yield the same Fe
abundance on average; (2) log g, from requiring abundance agreement between
neutral and ionized lines of Fe and Ti; (3) $V_{mic}$, by forcing weak and
strong Fe i lines to yield similar abundances; and (4) [M/H] by asking for
stellar model metallicity to be consistent with derived abundances of most
elements. We tried to follow the procedures outlined in Böcek Topcu et al.
(2015, 2016) for NGC 752 and NGC 6940.

However, NGC 7789 red giant stars are substantially redder/cooler than those
in the other two clusters, and its lines have significantly larger EWs.
Unfortunately, almost all Fe i lines employed for the other clusters lie on
the flat part of the curve-of-growth in NGC 7789 red giant stars. In general,
the strongest Fe i lines have low excitation energies. The relative abundance
insensitivity of these lines, combined with the line strength and excitation
correlation makes it very difficult to determine $V_{mic}$ and $T_{\rm eff}$
in the traditional manner.

Therefore we chose to derive $T_{\rm eff}$ values purely from atomic line
depth ratios (LDRs). This method was pioneered by Gray & Johanson (1991) and
involves identification of pairs of absorption lines whose strengths have very
different responses to variations in $T_{\rm eff}$. Gray & Johanson considered
G-K stars, and concentrated on the 6200 Å spectral domain that has many low-
excitation V i transitions that become much stronger with decreasing $T_{\rm
eff}$, usually pairing them with higher excitation Fe i transitions that are
much less sensitive to $T_{\rm eff}$ changes. The physical principles involved
are basic Boltzmann/Saha relations that describe the populations of
ionization/excitation states in stellar atmospheres (see extended discussions
of these issues in Gray 2008). The work on LDRs has expanded to include other
species in other spectral regions for stars in different temperature-gravity
domains, e.g., Strassmeier & Schordan (2000), López-Valdivia et al. (2019).
Here we follow Böcek Topcu et al. (2015, 2016) in using the LDR calibrations
of Biazzo et al. (2007) which are based on the original Gray & Johanson work.

With $T_{\rm eff}$ established the other parameters log g, $V_{mic}$, and
[M/H] were derived from line EWs as outlined above. In particular the
microturbulent velocity was determined from forcing weak and strong Fe i lines
to yield the same average abundances. The gravity was set exclusively through
the requirement to force abundance agreement between these two Fe ionization
states. We employed interpolated models from the ATLAS444 Available at
http://kurucz.harvard.edu/grids.html; model interpolation software was kindly
provided by A. McWilliam and I. Ivans. grid (Kurucz, 2011, 2018). To compute
abundances we used the current version of the LTE plane-parallel line analysis
code MOOG (Sneden, 1973).555 Available at
http://www.as.utexas.edu/$\sim$chris/moog.html The derived atmospheric
parameters are listed in Table 4, along with the abundances of Fe i and Fe ii.

The Fe abundances were determined from on average 49 Fe i and 10 Fe ii
species. From these slightly sub-solar values we set a uniform metallicity of
[M/H] = $-$0.1. The mean metallicity for NGC 7789, $<$[Fe/H]$>$ = $-$0.02 with
$\sigma$ = 0.05, is in reasonable accord with those determined in previous
high-resolution spectroscopic studies: [Fe/H] = $-$0.04, $\sigma$ = 0.05
Tautvaišienė et al. (2005); [Fe/H] = $+$0.04, $\sigma$ = 0.07 Pancino et al.
(2010); [Fe/H] = $+$0.02, $\sigma$ = 0.05 Jacobson et al. (2011); [Fe/H] =
$+$0.03, $\sigma$ = 0.07 Overbeek et al. (2015).

Our study is not a comprehensive abundance survey of NGC 7789. It focuses
mainly on the LiCNO element group and its relationship to He i $\lambda$10830
absorption strengths. We derived abundances from EWs of a few elements that
have been featured in past studies. They are listed in Table 5 and in Figure 3
we compare some [X/Fe] values to those in previous papers. Our abundance
uncertainties are larger than those reported in the comparison studies, but in
most part they are in accord for Mg, Si, and Ca. The NGC 7789 abundances for
Ti, Na, and Ni should be explored further in a study that concentrates on
warmer, less line-rich stars than we have gathered for our work.

Figure 3:  Comparison of NGC 7789 abundances in this paper to results from
some previous papers. The various studies are identified by color types, as
defined in the plot legend and are as follows: Overbeek et al. (2015),
Jacobson et al. (2011), Pancino et al. (2010), Tautvaišienė et al. (2005)

###  3.2 LITHIUM, CARBON, NITROGEN, AND OXYGEN

Figure 4:  Observed and synthetic spectra for the two NGC 7789 stars with very
high Li abundances. The $\lambda$6707 resonance lines shown in panels (a) and
(c) are so saturated that they cannot yield reliable abundances. The weaker,
higher-excitation $\lambda$6103 lines in panels (b) and (d) are better Li
features for these two stars. (Sneden et al., 2022). In the legend of each
panel, ”best’” indicates the synthesis using the abundance derived for the
named element, and ”no” indicates the synthesis computed without any
contribution of this element. The other syntheses have abundance offsets (in
dex) as indicated in the legends.  Figure 5:  Observed and synthetic spectra
for the features used to determine C, N, and O abundances in star 501, a
typical program star. The observed points are shown with open circles. The
meanings of the legends in each panel are the same as in Figure 4.

Light elements Li, C, N, and O may be altered during a star’s lifetime through
interior proton fusion reactions and envelope mixing. We derived abundances
for these elements via matches of observed and synthetic spectra.

  * •

Lithium: Nearly all Li abundances are derived from the Li i resonance doublet
at 6707.8 Å, but in a small percentage of cool stars with very large Li
abundances the excited-state transition at 6103.6 Å becomes detectable. We
used both of these features in our analysis for stars 193 and 301.

  * •

Carbon: the best available spectral features for our stars were lines near the
C2 Swan system (d${}^{3}\Pi_{g}-a^{3}\Pi_{u}$) (0$-$1) bandhead at 5635 Å. The
stronger (0$-$0) bandhead near 5165 Å is extremely crowded with other atomic
and molecular features, and our spectra in the region of the comparable-
strength (1$-$0) bandhead, 4725 Å, are very noisy. Our spectra are even less
reliable in the 4300 Å CH G-band wavelength domain. C i high-excitation lines
can be detected but are not reliable abundance indicators for NGC 7789 cool
giants.

  * •

Nitrogen: CN A${}^{2}\Pi-X^{2}\Sigma^{+}$ red system lines can be detected in
all wavelength regions beyond 6000 Å. We concentrated on strong (2$-$0) band
features near 8000 Å to derive N abundances and 12C/13C ratios.

  * •

The [O i] $\lambda$6300 ground-state line is the single reliable O abundance
indicator in our spectra. Care must be taken with this transition since it is
beset with blending by Ni i and CN contaminants, and can potentially be
compromised by telluric O2 absorption and night-sky [O i] emission.

To create atomic/molecular line lists for the synthetic spectrum calculations
we used the $linemake$ facility (Placco et al., 2021)666
https://github.com/vmplacco/linemake. The code creates lists of transitions
starting with the Kurucz (2011, 2018)777
http://kurucz.harvard.edu/linelists.html line compendium, and updating them
with transition data from recent laboratory studies mainly by the Wisconsin-
Madison atomic physics group (Den Hartog et al. 2021 and references therein)
and by the Old Dominion University molecular physics group (e.g., Brooke et
al. 2016 and references therein).

We estimated LiCNO abundances by comparing our observed spectra with synthetic
spectra computed with our model atmospheres (Table 4) and these line lists.
Full molecular equilibrium calculations were performed as part of the
computations. In Figure 4 we illustrate the observed/synthetic spectrum
matches for the Li i $\lambda$6707,6103 lines in the two Li-rich stars 193 and
301. Other program stars have Li abundances that are 2-3 orders of magnitude
smaller than those of stars 193 and 301, rendering weak $\lambda$6707 lines
(see also Figure 2) and invisible $\lambda$6103 lines.

For C and O, the formation of CO molecules can significantly modify the C2 and
O i number densities, so we derived these abundances iteratively. In Figure 5
we show the observed/synthetic matches for NGC 7789 star 501. Note that
although the C2 band near 5635 Å is always weak, it is very sensitive to
carbon abundance changes because of its double-carbon molecular structure (see
panel (a) of this figure). Second, we repeat the caution from above that while
the [O i] transition at 6300.3 Å is the only reliable oxygen abundance
indicator in optical spectra of red giant stars, it has significant blending
by other lines (panel (b) of the figure), and results from it should be
treated with caution. Finally, there are many 12CN features in the
$\lambda$8000 region, making total nitrogen abundance derivation
straightforward. However, in all of our NGC 7789 red giants the 13CN lines are
much weaker (Figure 5 panel c), and (like many other studies) we are forced to
estimate 12C/13C ratios mostly from the blended 13CN feature at 8004.5 Å. Our
derived LiCNO abundances are listed in Table 6.

###  3.3 Abundance Uncertainties

Our abundances depend directly on uncertainties in EW measurements and in
synthetic/observed spectrum comparisons. To assess EW uncertainties we
repeated EW and synthetic spectrum computations for multiple lines, varying
line parameters and continuum placement. We concluded that typical measurement
uncertainties are $\pm$5%, contributing $\sim$ $\pm$0.03 dex to the abundances
of individual spectroscopic features. The abundances also depend on the model
atmosphere parameter choices. Earlier in §3.1 we discussed the particular
analytical problem of NGC 7789 red giants. All of our program stars have
strong-lined spectra, and the majority of our measured transitions are on or
near the flat part of the curve of growth, where derived abundances become
sensitive to assumed/derived microturbulent velocity $V_{mic}$. The
difficulties involved in using strong lines for model parameter derivation
should be kept in mind. In Table 7 we list the responses of each species
abundance to changes in model parameters that cover the range of expected
uncertainties in $T_{\rm eff}$, logg, [M/H], and $V_{mic}$. For this table we
have chosen star 765, whose parameters are roughly in the middle of our
stellar sample: $T_{\rm eff}$ = 4515 K, log g = 2.3, [M/H] = $-$0.1, $V_{mic}$
= 1.4 km s-1. We present changes in abundances [X/H] (equivalent to changes in
log $\epsilon$) to show how each species responds to model parameter changes.
Lines of neutral species yield higher elemental abundances from $T_{\rm eff}$
increases due to larger ionization, and abundances from lines of ionized
species increase with increasing log g due to decreased ionization in higher-
gravity stars. Note that V i abundances make the largest changes as $T_{\rm
eff}$ varies; this is why the LDR method works very well for $T_{\rm eff}$
estimation for our cool giant stars. Additionally, the Li abundance
sensitivity to $T_{\rm eff}$ and no other atmospheric parameter in Table 7
confirms what has been known in the literature about this species for decades.
Finally, in general the changes in [X/Fe] will be smaller than those of [X/H]
when comparisons of neutral to neutral and ion to ion are performed as has
been done in this paper.

##  4 DISCUSSION

Figure 6:  Carbon isotopic ratios (panel a) and CNO abundances (panels b$-$d)
as functions of Li abundance in all NGC 7789 program stars, including star
193, which is a probable non-member of the cluster. The dotted lines in panels
(b$-$d) indicate the solar abundance values. Star 193 is the blue dot and Star
301 is the red dot.

In this section, we discuss our main observational results on the light
elements in NGC 7789 red giants, and consider their relationships to their
probable evolutionary states.

###  4.1 Li and CNO Abundances

These elements are participants in the various $p-p$ and CNO cycles of
hydrogen fusion, and we will consider them as a group. In Figure 6 we plot the
derived 12C/13C and CNO abundance ratios as functions of the Li abundances.
Inspection of this figure, as well as the mean values listed in Table 6
suggest that there are few indications of significant star-to-star CNO
variations, and little evidence of their correlations with Li. For the CNO
elements the abundance scatter is $\sigma$ $\simeq$ 0.1, suggesting that all
of our program stars have had similar abundance changes in the first dredge-up
near the base of the red giant branch (RGB). Our derived C abundances in NGC
7789 do not appear to be much depleted, $\langle$[C/Fe]$\rangle$ = 0.01. In
contrast, the classic studies of thin disk population giants derived lower C
abundances. For example, Lambert & Ries (1981) found $\langle$[C/Fe]$\rangle$
= $-$0.24$\pm$0.13. We suspect that the difference stems from our exclusive
use of the C2 $\lambda$5635 bandhead for C abundances. Our previous study of
NGC 6940 (Böcek Topcu et al., 2016) found that the $\lambda$5635 bandhead
yields C abundances on average 0.17 dex higher than those derived from the
much stronger $\lambda$5165 bandhead. As stated above, our spectra are not
trustworthy in wavelength regions less than about 5300 Å. However, if we were
to shift our derived [C/Fe] abundance ratios downward by $\simeq$0.15, then in
consequence our CN-based N abundances would increase by the same amount, from
$\langle$[N/Fe]$\rangle$ = 0.35 (Table 6) to $\sim$0.50, compared with the
Lambert & Ries value of $\langle$[N/Fe]$\rangle$ = 0.38$\pm$0.11. Thus we
regard our C and N abundances as probably consistent with first dredge-up
expectations.

Our carbon isotopic ratios are in the range 12C/13C = 12-40, with
$\langle$12C/13C$\rangle$ = 21 ($\sigma$ = 9), again consistent with values
that have been reported for field thin disk giants for decades (e.g., Tomkin
et al. 1976 and references therein).

Most stars in our sample have similar Li abundances. From Table 6, excluding
stars 193, 301 and 765, $\langle$log $\epsilon$(Li)$\rangle$ = 1.08 $\pm$ 0.04
($\sigma$ 0.15). This value is shared by 11 of our 14 program stars. If stars
353 and 468 with slightly smaller Li abundances (log $\epsilon$(Li) = 0.8) are
excluded the average becomes more uniform but does not change much:
$\langle$log $\epsilon$(Li)$\rangle$ = 1.14 $\pm$ 0.02 ($\sigma$ 0.05). In
total, the Li and CNO abundances of these stars probably result from standard
first dredge-up episodes that acted on a set of normal stars of similar masses
and initial chemical compositions. Accurate mass measurements for individual
stars would be welcome but would require astroseismology.

To help clarify the evolutionary states of our stars, in Figure 7 we repeat
Figure 1 except that different colors are used to indicate derived lithium
abundances. For the present purpose we do not plot star 193, the probable non-
member of NGC 7789. We suggest that all stars with log $\epsilon$(Li) $\sim$ 1
are first-ascent red giant branch (RGB) members. Star 468 [$(BP-
RP)_{0}$,$M_{G0}$] = [1.44,$-$1.85] is slightly brighter than the first-ascent
sequence and might be an asymptotic branch (AGB) star, but the separation is
small. A more convincing AGB assertion can be made for star 765
([1.29,$-$1.59]), which is nearly a magnitude brighter than RGB stars with
similar colors, and its Li abundance log $\epsilon$(Li) = 0.0 is much smaller
than any other program star. Star 1101 ([1.08,0.34]) is our sole
representative of the clump in NGC 7789, and being a clump star means it has
evolved beyond the first-ascent RGB.

Stars 193 (not plotted in Figure 7) and 301 (the blue point in the figure) may
or may not require separate interpretations. Although we agree with previous
work suggesting that star 193 is not a member of NGC 7789, that fact may have
little impact on the interpretation of its light element abundances. Stars 193
and 301 are solar-metallicity K giants with similar very high Li abundances.
It is likely that both of them underwent similar evolutionary events that
resulted in their very high surface Li contents. Star 193 departs from NGC
7789 cluster giants mainly in its relatively low [N/Fe] and consequently large
[C/N] (Table 6). Ejection of star 193 from NGC 7789 membership simply puts it
in good company with the many field red giants of Paper 1 that have enhanced
Li while showing no evidence for chromospheric activity anomalies.

Figure 7:  Another CMD for NGC 7789 with the data from Figure 1, using colors
to indicate the Li abundances for NGC 7789 members.

###  4.2 COMPARISON WITH HELIUM

Paper 1 reported He i $\lambda$10830 EWs for 10 NGC 7789 red giants; see Table
6. Five of the present program stars were not observed for Paper 1, and star
225 was included in Paper 1 but not included in our optical survey. In Figure
8 we plot Li abundances versus He i reduced widths log(RWHe).888 log(RWHe)
$\equiv$ log10(EW${}_{\rm He}/\lambda$) = log10(EWHe/10830) We highlight the
NGC 7789 stars and add in the field stars from Paper 1 for comparison.
Following the discussion in Paper 1, a yellow horizontal line is placed at log
$\epsilon$(Li) = 1.25 to approximately separate Li-rich giants from the
majority of normal Li-poor stars. A vertical line set at log(RWHe) = $-$4.85
provides a suggested separation between the vast majority of stars with weak
He i $\lambda$10830 absorption and the relatively few with anomalously strong
He lines. These lines are rough guides and their values should not be
interpreted rigidly.

In Figure 8 all program stars with log $\epsilon$(Li) $\leq$ 1.15 have weak
$\lambda$10830 absorption lines, as do the vast majority of field red giants.
Probable non-member star 193 has a large Li abundance without excess He
absorption, but many other field stars share this characteristic. Member star
301 has both a high Li and a large $\lambda$10830 feature. In fact, this star
has one of the largest He i lines found in the Paper 1 survey. All other stars
with log(RWHe) $\gtrsim$ $-$4.2 exhibit significant rotation spectral line
broadening, with derived velocities in the range $V$sin($i$) = 9$-$140 km s-1.
Paper 1 did not report rotational line broadening for star 301 but the derived
Gaussian smoothing was larger than for most stars. Therefore we have carefully
re-examined its $\lambda$10830 region with synthetic spectra. We now
tentatively can assign a small rotational value $V$sin($i$) = 6$\pm$1 km s-1
for star 301. However, this velocity broadening is nearly at the limit of our
ability to isolate rotation from the other broadening sources (thermal,
microturbulent, macrotrubulent, and instrumental), so caution should be
exercised in interpretation of this value. What is not in dispute is the
relatively small total line broadening of star 301 in comparison to many other
Li-rich red giants.

The vast majority of the 300-star field red giant sample in Paper 1 proved to
be mostly stars with colors and magnitudes consistent with the red clump and
general red horizontal branch. Stars with large Li abundances and He i
$\lambda$10830 line strengths both weak and strong were as likely to occupy
this CMD domain as those with neither Li nor He anomalies. The interpretation
in Paper 1 centered on assignment of fresh Li production to the helium flash
that preceded appearance of the stars on the red horizontal branch, as
outlined in §1 (see also Mallick et al. 2023). This idea may be applicable for
stars 301 and 193 as well. Star 301 has high Li and its strong $\lambda$10830
line argue for recent helium-flash disturbance of its outer envelope.

However, star 301’s evolutionary state is still uncertain. Based on its CMD
position, certain possibilities arise. Star 301 could be ascending the RGB, in
which case it would not have undergone any helium flash yet, and so both Li
and He features would likely be due to interactions with a binary companion.
We did not detect any binary companion in star 301’s spectrum. In this case,
this star’s apparent excess rotation suggests that it could have an unseen
companion like a white dwarf or small main sequence star that is simply too
dim to contribute any significant flux or spectral features. Figure 8 of Casey
et al. (2019) presents three models for how the presence of a binary
companion(s) like star(s) or planet(s) can explain the lithium-richness of
certain red giant stars. If star 301 is still ascending the RGB then tidal
interactions with the putative companion are to blame. A radial velocity
monitoring program might be useful for detecting the companion. More directly,
speculation on star 301’s “late” arrival on the RGB probably requires it to
have recently evolved from its previous status as a cluster BSS. But such a
scenario must account for star 301’s carbon istopic ratio, which is among the
smallest of any star in our sample (Table 6). Blame for its small 12C/13C
would shift to the companion’s transferred material.

Star 301 might be a normal red clump giant, even though its absolute magnitude
($M_{G0}$) appears to be $\sim$0.5 mag brighter than the NGC 7789 clump
(Figure 7). No reasonable amount of (unproven) differential reddening
corrections would place star 301 in the NGC 7789 clump (see the reddening
vector in Figure 1). Finally, NGC 7789 may have a more complex member
population than most open clusters. From investigation of CN and CH molecular
band strengths of NGC 7789 giant stars, Carrera & Martínez-Vázquez (2013)
suggested that this cluster may have a CN intracluster spread of the sort seen
in typical globular clusters. They urged caution in interpretation of their
observations pending a much larger CN survey, but it is possible the star 301
is a representative of an anomalous sub-population of NGC 7789. If this is
true, then star 301 should be dealt with as a special evolutionary case.
Asteroseismological data for this star would help its interpretation.

Star 193 has an absolute magnitude more consistent with a red clump star but
it appears to be too red (Figure 1). This could be alleviated by assignments
of a larger reddening value for this star than what we have assumed for the
cluster. Though star 193 is a helium flash candidate, its lithium could too be
explained by a binary companion as one of the binary companion models in Casey
et al. (2019) involve lithium richness in post-he-flash stars but due to a
binary companion rather than the helium flash.

Further progress in this area requires increasing the number of Li-rich red
giants that have He i $\lambda$10830 data and other common characteristics.
Efforts by our group are underway to gather such observations for Kepler field
giants, whose evolutionary states can be clearly defined through
asteroseismology. We are also collecting spectra of red giants with known
rotational velocities to test that possible connection. Additionally the
number of Li-rich red giant members in other open clusters is slowly growing
(Magrini et al. 2021 and references therein). Observational campaigns on these
clusters should be undertaken in the future.

Figure 8:  Correlation between lithium abundances and He absorption line
strengths in Paper 1 and the present work. Identification of the points are in
the figure legend.

##  5 CONCLUSIONS

In this paper we have described the derivation of model atmospheric parameters
and LiCNO abundances for 15 red giants of the open cluster NGC 7789. We have
explored the relationship between Li abundances, CNO abundances, and He
$\lambda$10830 absorption strengths. We find that the majority of observed NGC
7789 stars have consistent Li abundances, log $\epsilon$(Li) $\simeq$ 1.0. Two
anomalous stars 193 and 301 have log $\epsilon$(Li) of 3.5 and 3.0
respectively; they are unambiguous examples of the Li-rich red giant subclass.
The probable non-membership of star 193 hinders a detailed investigation of
its evolutionary state, but its Li abundance is well determined. The He i
$\lambda$10830 transition does not lead to an easy interpretation of the Li
richness of stars 193 and 301, given the feature is present in one star but
not the other. In concert with other papers on Li abundance enhancements in
field red giants, we suggest (but cannot prove) that the Li and He
peculiarities detected in a small percentage of giants are due to recent or
long-ago envelope turbulence that accompanied the helium flashes in these
stars. Li abundances and He line strength measurements for many other NGC 7789
red giants should go a long way to clarifying the overall stellar evolution
story of this cluster.

We thank the referee for helpful comments that improved the paper These
results are based on observations obtained with the Tull Spectrograph aboard
the 2.7 meter Harlan J Smith Telescope at McDonald Observatory. We acknowledge
support from NSF grant AST-1616040 (CS). We also appreciate additional
financial support and resources from the University of Texas at Austin’s
Astronomy Department and the College of Natural Sciences. This work has made
use of data from the European Space Agency (ESA) mission Gaia
(https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and
Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has
been provided by national institutions, in particular the institutions
participating in the Gaia Multilateral Agreement. Smith (Tull Spectrograph),
HET (HPF)

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Table 1: Basic Data for NGC 7789 Quantity | Valueaafrom WEBDA or SIMBAD except in the case of the Reddening and the age (age is from (Gao, 2018)). For the reddening, we used the value found in (Jacobson et al., 2011) and (Overbeek et al., 2015) quoting (Gim et al., 1998) and (Tautvaišienė et al., 2005)
---|---
RA | 23:57:24
Dec | +56:42:30
$l$ (Galactic) | 115.53∘
$b$ (Galactic) |  $-$5.53∘
Parallax | 0.45 mas
Distance | 2337 pc
Reddening | 0.28 mag
Distance Modulus | 11.72
Age | 1.5 Gyr
$\mu$ (RA) |  $-$0.92 mas
$\mu$ (Dec) |  $-$1.93 mas
Radial Velocity |  $-$54.5 km s-1
Table 2: Basic Data for Target Starsaaadopted from SIMBAD StarbbKüstner (1923) | parallax |  $\mu$ (RA) |  $\mu$ (Dec) | $G$ | $BP-RP$ | $V$ | $B-V$
---|---|---|---|---|---|---|---
| mas | mas | mas |  |  |  |
72 | 0.5040 |  $-$0.775 |  $-$2.071 | 10.22 | 2.06 | 11.05 | 1.74
193 | 0.6742 |  $+$8.491 |  $-$2.338 | 12.12 | 1.63 | 12.61 | 1.42
301 | 0.5144 |  $-$1.123 |  $-$2.444 | 11.72 | 1.60 | 12.27 | 1.34
329 | 0.4949 |  $-$0.883 |  $-$1.890 | 11.80 | 1.66 | 12.31 | 1.38
353 | 0.4855 |  $-$0.820 |  $-$2.017 | 12.10 | 1.65 | 12.59 |  $\cdots$
461 | 0.4793 |  $-$0.968 |  $-$1.933 | 10.69 | 1.88 | 11.35 | 1.63
468 | 0.4829 |  $-$0.886 |  $-$1.914 | 10.53 | 1.83 | 11.61 | 1.14
494 | 0.4898 |  $-$0.825 |  $-$1.985 | 9.86 | 2.22 | 10.74 | 1.68
501 | 0.4749 |  $-$0.830 |  $-$2.037 | 10.51 | 1.95 | 11.26 | 1.71
637 | 0.4882 |  $-$0.983 |  $-$1.716 | 11.90 | 1.65 | 12.40 | 1.42
669 | 0.4830 |  $-$0.775 |  $-$2.071 | 10.85 | 1.83 | 11.46 | 1.58
765 | 0.4175 |  $-$1.065 |  $-$2.033 | 11.11 | 1.68 | 11.59 | 1.21
970 | 0.4800 |  $-$0.970 |  $-$2.110 | 11.29 | 1.80 | 11.89 |  $\cdots$
971 | 0.5398 |  $-$1.029 |  $-$2.105 | 9.97 |  | 11.05 | 1.90
1066 | 0.4867 |  $-$0.979 |  $-$2.055 | 11.43 | 1.73 | 11.99 |  $\cdots$
1101 | 0.5004 |  $-$0.876 |  $-$2.013 | 12.65 | 1.48 | 13.05 |  $\cdots$
$<>$cccomputed without star 193 | 0.4881 |  $-$0.919 |  $-$2.046 |  |  |  |
$\sigma$cccomputed without star 193 | 0.0256 | 0.107 | 0.154 |  |  |  |
Table 3: Equivalent Widths

Species | $\lambda$ | $\chi$ | log($gf$) | 72 | 193 | 301 | 329 | 353 | 461 | 468 | 494 | 501 | 637 | 669 | 765 | 978 | 1066 | 1101
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
Na I | 5682.64 | 2.101 | -0.70 | 216 | 194 | 173 | 181 | 175 | 210 | 209 | 212 | 217 | 186 | 197 |  $\cdots$ | 194 | 180 | 146
Na I | 6154.23 | 2.101 | -1.56 | 139 | 109 | 94 | 98 | 97 | 128 | 138 | 144 | 133 | 99 | 118 | 105 | 112 | 105 | 71
Na I | 6160.75 | 2.103 | -1.26 | 154 | 130 | 112 | 120 | 114 | 142 | 144 | 195 | 147 | 118 | 136 | 121 | 132 | 125 | 93
Mg I | 5528.41 | 4.343 | -0.62 | 290 | 321 | 275 | 266 | 261 | 299 | 290 | 291 |  $\cdots$ | 236 | 285 |  $\cdots$ | 272 | 279 | 209
Mg I | 5711.08 | 4.343 | -1.83 | 154 | 166 | 143 | 144 | 146 | 161 | 141 | 148 | 157 | 152 | 161 | 154 | 151 | 151 | 113
Mg I | 7811.11 | 5.941 | -0.95 | 99 | 143 | 80 | 79 | 81 | 94 |  $\cdots$ | 101 | 93 | 86 | 91 | 85 | 81 | 92 |  $\cdots$
Al I | 6696.02 | 3.140 | -1.35 | 140 | 212 | 99 | 131 | 123 | 163 | 121 | 150 | 126 | 102 | 148 | 101 | 109 | 130 | 72
Al I | 6698.67 | 3.140 | -1.64 | 90 | 88 | 54 | 76 | 66 | 91 | 88 | 99 | 88 | 74 | 86 | 73 | 78 | 75 | 41
Al I | 7835.30 | 4.018 | -0.65 | 111 | 125 | 85 | 81 | 74 | 103 | 93 | 117 | 109 | 79 | 98 | 81 | 97 | 85 | 64
Si I | 5488.98 | 5.614 | -1.90 | 29 | 54 | 33 | 44 | 42 | 33 | 36 | 25 | 31 | 46 | 40 | 31 | 41 | 41 | 34
Si I | 5517.53 | 5.082 | -2.61 | 22 | 35 | 33 | 27 | 23 | 32 |  $\cdots$ | 35 |  $\cdots$ | 36 | 34 | 31 |  $\cdots$ | 28 | 25

Note. — All EWs are in units of Å (This table is available in its entirety in
machine-readable form.)

Table 4: Model atmosphere parameters Star | $T_{\rm eff}$ | [M/H] | log g | $V_{mic}$ | [Fe/H] | $\sigma$ | #lines | [Fe/H] | $\sigma$ | #lines
---|---|---|---|---|---|---|---|---|---|---
| (K) |  |  | km s-1 | (I) | (I) | (I) | (II) | (II) | (II)
72 | 4240 |  $-$0.1 | 2.00 | 1.00 | 0.01 | 0.18 | 50 | 0.01 | 0.26 | 12
193 | 4590 |  $-$0.1 | 2.50 | 1.60 |  $-$0.02 | 0.10 | 49 |  $-$0.02 | 0.16 | 12
301 | 4730 |  $-$0.1 | 2.60 | 1.55 |  $-$0.06 | 0.13 | 49 |  $-$0.07 | 0.17 | 11
329 | 4625 |  $-$0.1 | 2.40 | 1.40 | 0.04 | 0.12 | 52 |  $-$0.03 | 0.17 | 12
353 | 4675 |  $-$0.1 | 2.40 | 1.60 |  $-$0.13 | 0.12 | 50 |  $-$0.13 | 0.22 | 10
461 | 4315 |  $-$0.1 | 2.10 | 1.50 |  $-$0.01 | 0.15 | 48 |  $-$0.01 | 0.14 | 7
468 | 4370 |  $-$0.1 | 2.10 | 1.60 |  $-$0.05 | 0.14 | 44 |  $-$0.07 | 0.14 | 10
494 | 4170 |  $-$0.1 | 1.90 | 1.45 | 0.02 | 0.25aaStar 494 had an exceptionally low signal to noise (SNR) which is why we ended up with an unusually large standard deviation. In addition, the reason there is such a wide range of standard deviations for Fe (II) is that many Fe (II) lines lay in noisy regions of our spectra, making them harder to measure and depending on the star. | 49 | 0.02 | 0.19 | 9
501 | 4295 |  $-$0.1 | 2.10 | 1.50 | 0.00 | 0.15 | 46 | 0.00 | 0.11 | 7
637 | 4620 |  $-$0.1 | 2.60 | 1.50 |  $-$0.06 | 0.10 | 51 |  $-$0.06 | 0.08 | 9
669 | 4390 |  $-$0.1 | 2.00 | 1.45 |  $-$0.03 | 0.14 | 50 |  $-$0.02 | 0.33 | 10
765 | 4515 |  $-$0.1 | 2.30 | 1.40 | 0.07 | 0.15 | 50 | 0.06 | 0.12 | 10
970 | 4455 |  $-$0.1 | 2.30 | 1.55 |  $-$0.02 | 0.13 | 49 |  $-$0.02 | 0.08 | 9
1066 | 4530 |  $-$0.1 | 2.30 | 1.50 |  $-$0.03 | 0.13 | 53 |  $-$0.02 | 0.15 | 9
1101 | 5050 |  $-$0.1 | 2.65 | 1.40 |  $-$0.08 | 0.12 | 47 |  $-$0.08 | 0.13 | 7
$<>$ |  |  |  |  |  $-$0.02 | 0.14 | 49 |  $-$0.03 | 0.16 | 10
$\sigma$ | 225 |  $\cdots$ | 0.24 | 0.07 | 0.05 | 0.04 |  $\cdots$ | 0.05 | 0.07 |  $\cdots$
Table 5: Output Abundances in bracket ([X/Fe]) form Star | Na | Mg | Al | Si | Ca | Sc2aaFor these 5 elements, the abundances presented were found using the ”blends” task instead of ”abfind” due to the presence of hyperfine splitting. All other abundances were found using ”abfind”. | Ti1 | Ti2 | VaaFor these 5 elements, the abundances presented were found using the ”blends” task instead of ”abfind” due to the presence of hyperfine splitting. All other abundances were found using ”abfind”. | Cr1 | Cr2 | MnaaFor these 5 elements, the abundances presented were found using the ”blends” task instead of ”abfind” due to the presence of hyperfine splitting. All other abundances were found using ”abfind”. | Fe1 | Fe2 | Ni | LaaaFor these 5 elements, the abundances presented were found using the ”blends” task instead of ”abfind” due to the presence of hyperfine splitting. All other abundances were found using ”abfind”. | EuaaFor these 5 elements, the abundances presented were found using the ”blends” task instead of ”abfind” due to the presence of hyperfine splitting. All other abundances were found using ”abfind”.
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
72 | 0.61 | 0.01 | 0.33 | 0.29 | 0.24 | 0.48 | 0.44 | 0.51 | 0.49 | 0.31 | 0.26 | 0.29 | 0.00 | 0.00 | 0.47 | 0.35 | 0.22
193 | 0.46 | 0.27 | 0.83 | 0.50 | 0.17 | 0.15 | 0.20 | 0.23 | 0.22 | 0.18 | 0.32 | 0.20 | 0.00 | 0.00 | 0.24 | 0.18 | 0.55
301 | 0.38 |  $-$0.03 | 0.12 | 0.20 | 0.32 | 0.13 | 0.20 | 0.33 | 0.20 | 0.31 | 0.23 | 0.05 | 0.00 | 0.00 | 0.12 | 0.24 | 0.17
329 | 0.45 |  $-$0.04 | 0.35 | 0.35 | 0.17 | 0.09 | 0.14 | 0.33 | 0.14 | 0.18 | 0.15 | 0.13 | 0.00 | 0.00 | 0.22 | 0.25 | 0.28
353 | 0.45 | 0.03 | 0.29 | 0.38 | 0.16 | 0.10 | 0.15 | 0.22 | 0.21 | 0.18 | 0.20 | 0.11 | 0.00 | 0.00 | 0.15 | 0.29 | 0.32
461 | 0.50 | 0.04 | 0.43 | 0.39 | 0.18 | 0.20 | 0.27 | 0.53 | 0.34 | 0.18 | 0.29 | 0.22 | 0.00 | 0.00 | 0.33 | 0.27 | 0.29
468 | 0.64 |  $-$0.13 | 0.22 | 0.42 | 0.20 | 0.16 | 0.28 | 0.16 | 0.32 | 0.15 | 0.07 | 0.21 | 0.00 | 0.00 | 0.35 | 0.37 | 0.48
494 | 0.77 | 0.01 | 0.43 | 0.39 | 0.18 | 0.11 | 0.47 | 0.69 | 0.58 | 0.30 | 0.53 | 0.12 | 0.00 | 0.00 | 0.46 | 0.36 | 0.19
501 | 0.57 | 0.19 | 0.26 | 0.29 | 0.18 | 0.16 | 0.33 | 0.41 | 0.41 | 0.33 | 0.20 | 0.13 | 0.00 | 0.00 | 0.35 | 0.38 | 0.46
637 | 0.41 |  $-$0.04 | 0.16 | 0.42 | 0.12 | 0.20 | 0.13 | 0.26 | 0.19 | 0.15 | 0.37 | 0.16 | 0.00 | 0.00 | 0.25 | 0.35 | 0.50
669 | 0.52 | 0.07 | 0.41 | 0.28 | 0.20 | 0.16 | 0.32 | 0.42 | 0.32 | 0.23 | 0.31 | 0.30 | 0.00 | 0.00 | 0.31 | 0.24 | 0.27
765 | 0.27 | 0.10 | 0.01 | 0.24 | 0.09 | 0.14 | 0.13 | 0.19 | 0.13 | 0.12 | 0.07 | 0.17 | 0.00 | 0.00 | 0.25 | 0.21 | 0.36
970 | 0.44 |  $-$0.07 | 0.16 | 0.40 | 0.12 | 0.13 | 0.14 | 0.29 | 0.21 | 0.10 | 0.17 | 0.15 | 0.00 | 0.00 | 0.25 | 0.31 | 0.35
1066 | 0.41 | 0.02 | 0.27 | 0.38 | 0.19 | 0.10 | 0.13 | 0.45 | 0.15 | 0.15 | 0.12 | 0.00 | 0.00 | 0.00 | 0.24 | 0.25 | 0.26
1101 | 0.43 | 0.03 | 0.06 | 0.19 | 0.05 |  $-$0.07 | 0.07 |  $-$0.03 | 0.04 | 0.02 |  $-$0.21 |  $-$0.07 | 0.00 | 0.00 | 0.05 | 0.21 | 0.45
Table 6: LICNO Abundances, Carbon Isotopic Ratios, He EWs Star |  12C/13C | [Fe/H] | [C/Fe] | [N/Fe] | [O/Fe] | [C/N] | log $\epsilon$(Li) | EW(He)aaTaken from Paper 1
---|---|---|---|---|---|---|---|---
72 | 15 | 0.01 | 0.09 | 0.29 | 0.22 |  $-$0.20 | 1.15 | 95
193 | 16 |  $-$0.02 | 0.12 | 0.12 | 0.07 | 0.00 | 3.50 | 143
301 | 12 |  $-$0.07 |  $-$0.10 | 0.27 | 0.07 |  $-$0.37 | 3.00 | 1040
225 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | 102
329 | 35 |  $-$0.04 |  $-$0.01 | 0.39 | 0.14 |  $-$0.40 | 1.20 | 98
353 | 40 |  $-$0.13 | 0.03 | 0.38 | 0.18 |  $-$0.35 | 0.80 | 102
461 | 18 |  $-$0.01 | 0.01 | 0.41 | 0.21 |  $-$0.40 | 1.10 | 129
468 | 16 |  $-$0.06 | 0.11 | 0.51 | 0.36 |  $-$0.40 | 0.80 |  $\cdots$
494 |  $\cdots$ | 0.02 | 0.16 | 0.23 | 0.28 |  $-$0.07 | 1.05 | 39
501 | 22 | 0.00 |  $-$0.03 | 0.55 | 0.05 |  $-$0.58 | 1.15 | 145
637 | 25 |  $-$0.06 | 0.03 | 0.31 | 0.11 |  $-$0.28 | 1.20 | 110
669 | 22 |  $-$0.03 |  $-$0.07 | 0.28 | 0.08 |  $-$0.35 | 1.10 |  $\cdots$
765 | 15 | 0.06 |  $-$0.06 | 0.34 | 0.09 |  $-$0.40 |  $\lesssim$0.0 |  $\cdots$
970 | 12 |  $-$0.02 | 0.00 | 0.47 | 0.17 |  $-$0.47 | 1.10 |  $\cdots$
1066 | 30 |  $-$0.03 |  $-$0.12 | 0.38 | 0.03 |  $-$0.50 | 1.20 |  $\cdots$
1101 | 15 |  $-$0.08 | 0.08 | 0.33 | 0.08 |  $-$0.25 |  $\lesssim$0.5 |  $\cdots$
$<>$ | 21 |  $-$0.03 | 0.02 | 0.35 | 0.14 |  $-$0.33 |  |
$\sigma$ | 9 | 0.05 | 0.10 | 0.11 | 0.10 |  $-$0.16 |  |
Table 7: Abundance Dependencies on Model Atmosphere Parameters Quantity |  $\Delta$(log $\epsilon$) |  $\Delta$(log $\epsilon$) |  $\Delta$(log $\epsilon$) |  $\Delta$(log $\epsilon$)
---|---|---|---|---
parameter |  $\Delta$$T_{\rm eff}$ |  $\Delta$log g |  $\Delta$[M/H] |  $\Delta$$V_{mic}$
change |  $+$100 K |  $+$0.2 |  $+$0.3 |  $-$0.2km s-1
Fe i | 0.02 | 0.02 | 0.06 | 0.10
Fe ii |  $-$0.14 | 0.16 | 0.13 | 0.04
Na i | 0.07 |  $-$0.02 | 0.01 | 0.08
Mg i | 0.02 |  $-$0.03 | 0.03 | 0.05
Al i | 0.06 |  $-$0.01 |  $-$0.02 | 0.06
Si i |  $-$0.06 | 0.06 | 0.07 | 0.04
Ca i | 0.10 |  $-$0.06 | 0.01 | 0.19
Sc ii |  $-$0.03 | 0.11 | 0.10 | 0.13
Ti i | 0.14 | 0.00 |  $-$0.01 | 0.10
Ti ii |  $-$0.04 | 0.11 | 0.10 | 0.11
V i | 0.17 | 0.10 | 0.00 | 0.20
Cr i | 0.11 |  $-$0.01 | 0.01 | 0.09
Cr ii |  $-$0.09 | 0.11 | 0.09 | 0.06
Mn i | 0.08 |  $-$0.08 | 0.09 | 0.11
Ni i | 0.12 | 0.06 | 0.08 | 0.13
La ii | 0.01 | 0.11 | 0.10 | 0.15
Eu ii |  $-$0.01 | 0.12 | 0.10 | 0.15
Li i | 0.15 | 0.00 | 0.00 | 0.00
C2 |  $-$0.02 | 0.05 | 0.05 | 0.00
CN |  $-$0.03 | 0.15 | 0.15 | 0.05
$[$O i$]$ | 0.02 | 0.13 | 0.12 | 0.01