The Astrophysical Journal, 965:189 (15pp), 2024 April 20

https://doi.org/10.3847/1538-4357/ad293e© 2024. The Author(s). Published by the American Astronomical Society.

A JWST Project on 47 Tucanae. Overview, Photometry, and Early Spectroscopic Resultsof M Dwarfs and Observations of Brown Dwarfs*
A. F. Marino1,2 , A. P. Milone1,3 , M. V. Legnardi3 , A. Renzini1 , E. Dondoglio1 , Y. Cavecchi4 , G. Cordoni5 ,A. Dotter6 , E. P. Lagioia7 , T. Ziliotto3 , M. Bernizzoni3, E. Bortolan3, M. G. Carlos8 , S. Jang9 , A. Mohandasan3 ,F. Muratore3, and M. Tailo1,101Istituto Nazionale di Astroﬁsica - Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, Padova IT-35122, Italy; anna.marino@inaf.it
2Istituto Nazionale di Astroﬁsica - Osservatorio Astroﬁsico di Arcetri, Largo Enrico Fermi, 5, Firenze IT-50125, Italy
3
Dipartimento di Fisica e Astronomia “Galileo Galilei” - Università di Padova, Vicolo dell’Osservatorio 3, Padova IT-35122, Italy
4
Departament de Fisica, EEBE, Universitat Politecnica de Catalunya, Av. Eduard Maristany 16, E-08019 Barcelona, Spain
5
Research School of Astronomy and Astrophysics, Australian National University, Canberra, ACT 2611, Australia
6
Department of Physics and Astronomy, Dartmouth College, 6127 Wilder Laboratory, Hanover, NH 03755, USA
7South-Western Institute for Astronomy Research, Yunnan University, Kunming 650500, Peopleʼs Republic of China
8
Theoretical Astrophysics, Department of Physics and Astronomy, Uppsala University, Box 516, SE-751 20 Uppsala, Sweden
9
Center for Galaxy Evolution Research and Department of Astronomy, Yonsei University, Seoul 03722, Republic of Korea
10Dipartimento di Fisica e Astronomia Augusto Righi, Università degli Studi di Bologna, Via Gobetti 93/2, 40129 Bologna, Italy
Received 2024 January 12; revised 2024 February 9; accepted 2024 February 12; published 2024 April 19

Abstract
James Webb Space Telescope (JWST) observations have been demonstrated to be efﬁcient in detecting multiple
stellar populations in globular clusters (GCs) in the low-mass regime of M dwarfs. We present an overview, andﬁrst results, of different projects that can be explored by using the JWST observations gathered under program
GO2560 for 47 Tucanae, the ﬁrst program entirely devoted to the investigation of multiple populations in verylow-mass stars, which includes spectroscopic data for the faintest GC stars for which spectra are available. Ourcolor–magnitude diagram (CMD) shows some substructures for ultracool stars, including gaps and breaks in slope.
In particular, we observe both a gap and a minimum in the F322W2 luminosity function less than 1 mag apart, anddiscuss which it could be associated with the H-burning limit. We detect stars fainter than this minimum, very
likely brown dwarfs. We corroborate the ubiquity of the multiple populations across different masses, from
∼0.1 Me up to red giants (∼0.8 Me). The oxygen range inferred for the M dwarfs, both from the CMD and fromthe spectra of two M dwarfs associated with different populations, is similar to that observed for giants. We havenot detected any difference between the fractions of stars in distinct populations across stellar masses  0.1 Me.
This work demonstrates the JWSTʼs capability in uncovering multiple populations within M dwarfs and illustrates
the possibility to analyze very-low-mass stars in GCs approaching the H-burning limit and the brown-dwarf
sequence.
Uniﬁed Astronomy Thesaurus concepts: M dwarf stars (982); Population II stars (1284); CCD photometry (208)
multiple populations (MPs) in GCs and several physicalprocesses have been proposed for their formation. However,a consensus regarding their origin is still missing (see thereviews by Renzini et al. 2015; Bastian & Lardo 2018; Milone& Marino 2022, and references therein).
The current understanding of the “observed framework”
provides two main competing scenarios. According to themultiple generations scenarios, MPs are attributed to distinct
bursts of star formation at different epochs. 2P stars are bornfrom material processed and ejected by 1P stars (e.g., Cottrell& Da Costa 1981; Dantona et al. 1983; Renzini et al. 2022).
Various kinds of 1P polluters have been proposed, including
massive binary, asymptotic giant branch (AGB), rotating, andsupermassive stars (e.g., Ventura et al. 2001; Decressin et al.
2007; de Mink et al. 2009; Krause et al. 2013; Denissenkov &Hartwick 2014; D’Antona et al. 2016; Calura et al. 2019;
Lacchin et al. 2024). However, regardless of proposed
polluters, the amount of available material for 2P star formationis only a small fraction of the total present-day mass of 1P stars,
a challenge known as the “mass budget problem.” To cope with
this, 1P stars are required to be substantially more massive attheir formation and the progenitors of GCs should have lost alarge fraction of them, thus providing a signiﬁcant contribution
to the mass of the Galactic halo. In this context, it has been

1. Introduction
The majority of globular clusters (GCs) harbor diverse stellar
populations characterized by distinct chemical compositions.
This includes a ﬁrst population (1P) of stars with compositions
akin to ﬁeld stars exhibiting comparable metallicity, alongside
one or more “second” populations (2Ps) enriched in helium,
nitrogen, aluminum, and sodium, while being depleted incarbon and oxygen (e.g., Kraft 1994; Bastian & Lardo 2018;
Gratton et al. 2019; Milone & Marino 2022).
Since the ﬁrst discoveries, the intricate complexities underlying the coexistence of distinct stellar populations in GCs have
posed a challenge to our understanding of both stellar evolution
and star formation in the early Universe. In the past decades, ahuge observational effort has been devoted to the study of* The main data set presented in this article was obtained from the Mikulski
Archive for Space Telescopes (MAST) at the Space Telescope Science
Institute. The speciﬁc observations analyzed can be accessed via doi:10.17909/
6t82-4360.
Original content from this work may be used under the termsof the Creative Commons Attribution 4.0 licence. Any further
distribution of this work must maintain attribution to the author(s) and the titleof the work, journal citation and DOI.

1

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

investigated including NGC 6121 (M4; Milone et al. 2014),
Omega Centauri (Milone et al. 2017a), and NGC 6752 (Milone
et al. 2019). These works highlight the presence of similar
chemical variations among low- and higher-mass stars,
supporting the multiple generations scenario. Dondoglio et al.
(2022) present an analysis of a sample of GCs below the MS
knee, and provided for the ﬁrst time, the MFs of MPs in
NGC 2808 and M4 along a wide range of stellar masses, from
∼0.2 to ∼0.8 Me, ﬁnding that the fraction of MPs does notdepend on stellar mass, thus further challenging the accretion
scenario. First evidence of multiple stellar populations in verylow-mass stars from JWST data has been reported for
47 Tucanae (Milone et al. 2023a), M92 (Ziliotto et al. 2023),
and NGC 6440 (Cadelano et al. 2023). In this paper we exploit
deep JWST data to reach well below the hydrogen-burning
limit in 47 Tucanae, and into the brown-dwarf regime.
This work provides an overview of a project, mainly basedon observations gathered under the JWST program GO2560.
The project aims at the analysis of multiple stellar populations
among very-low-mass stars, in the domain of M dwarfs, in the
GC NGC 104 (47 Tucanae). Since the 1970s, the presence ofmultiple stellar populations with different light-element
abundances in this GC has been strongly established from theanalysis of bright RGB stars (e.g., Dickens et al. 1979; Bellet al. 1983; Brown et al. 1990; Briley et al. 1991). Nowadays,
we know that 47 Tucanae hosts two main stellar populations of
1P and 2P stars, which have been photometrically detectedalong the main evolutionary phases, including the MS,
supergiant branch, RGB, horizontal branch (HB), and AGB(e.g., Anderson et al. 2009; Milone et al. 2012a; Tailo et al.
2020; Lagioia et al. 2021; Jang et al. 2022; Lee 2022), andcorrespond to stars with different light-element abundances(Marino et al. 2019). A range ∼0.4 dex in [C/Fe], ∼0.5 dex in[O/Fe], and ∼1.0 dex in [N/Fe] (Carretta et al. 2009;
Dobrovolskas et al. 2014; Marino et al. 2016) has been
observed, whereas helium spans an interval of δY = 0.05 inmass fraction (Milone et al. 2018b). Astrometric plus spectroscopic analysis, mostly based on RGB stars or bright MS stars,
shows that these stellar populations exhibit different radial
distributions and internal kinematics (e.g., Richer et al. 2013;
Milone et al. 2018a; Cordoni et al. 2020).
In this work, we exploit the JWST’s capabilities to explore
the regime of low-mass stars in 47 Tucanae, starting from
∼0.1 Me and reaching the H-burning limit and beyond. First
photometric data for this program, observed on 2022 July 13,
have recently demonstrated the JWSTʼs capability in uncovering MPs within the low-mass realm of M dwarfs, marking
indeed a pivotal advancement in the possibility of understanding the formation and evolution of GCs (Milone et al.
2023a).
Here, we present the latest observations collected in 2023
September, including both photometric and spectroscopic data,
which allowed us to explore the faintest stars of 47 Tucanae,
encompassing a number of astrophysics issues, from theanalysis of multiple stellar populations to the observation ofbrown dwarfs, and to the spectral detection of chemicalvariations. For the ﬁrst time, we provide direct spectroscopic
evidence of MPs among M dwarfs, obtained from the ﬁrst
spectra ever observed in the faint MS of a GC.
The overview of this work is as follows: Section 2 describes
the data set while Section 3 presents the color–magnitude
diagrams (CMDs) obtained from NIRCam images, and focusesargued that GCs would have formed inside dwarf galaxies as aresult of a cooling catastrophe, with a fraction of the dwarf
itself having contributed material for the formation of 2Ps(Renzini et al. 2022). As opposed to a compact, massive
progenitor that would not lose much mass, a dense GC
embedded in an extended envelope would easily survive tidal
interaction while losing most of the envelope itself, e.g., assuggested by N-body simulations (Lacchin et al. 2024).
Alternatively, it has been suggested that a fraction of stars
from the same generation successively accreted material
processed and ejected by massive stars of their own generation(Bastian et al. 2013; Gieles et al. 2018). This scenario was
proposed as an attempt to cope with the mass budget issue,
assuming that a much lesser amount of processed material
would be required to produce the 2Ps, compared to the multiple
stellar generations case. However, if only a small amount ofaccreted material was sufﬁcient to produce the 2Ps, then one
would expect this material to be mixed over the wholeenvelope as stars reach the red giant branch (RGB), thus
producing a difference in composition with respect to mainsequence (MS) stars, which is not observed. Even more
fundamentally, as 2P stars are enriched in helium, the putative
accreted material would have a higher mean molecular weightcompared to underlying layers, thus causing Rayleigh–Taylor
instabilities to mix the accreted material with most of the star.
Thus, accretion seems not to solve the mass budget problem, as
it hardly account for the discreteness of MPs (Renzini et al.
2015). More recently, the occurrence of stellar mergers within
forming binary-rich GCs has been suggested as an explanation
for the MPs observed in GCs (Wang et al. 2020).
The two scenarios (multiple stellar generations and accretion) result in different expectations for the population pattern
across stellar masses. Based on the multiple generations
scenario, the population pattern should be identical for stars
with different masses, which means that both the high-mass
and low-mass stars formed during each burst of star formation
should exhibit similar chemical compositions. On the other
hand, in the accretion scenario the amount of accreted material
depends on stellar mass, and it is proportional to the square of
the stellar mass (e.g ., µ 2 in the case of Bondi-likeaccretion; Bondi & Hoyle 1944). As very-low-mass stars areless efﬁcient in accreting polluted material, the differences inhelium, carbon, nitrogen, and oxygen abundances commonly
observed among red giants would systematically decrease in
the M-dwarf domain.
The predictions of the two scenarios also diverge concerning
the mass function (MF) of the MPs. If the 2Ps grew out of
Bondi accretion, the MF of 2Ps should be signiﬁcantly ﬂattercompared to that of 1Ps (Ballesteros-Paredes et al. 2015). In
contrast, if there were multiple star formation episodes, a muchless pronounced difference, if any, is expected between the MF
slopes of 1P and 2P stars. As the multiple stellar population
pattern across different masses is a key ingredient used toconstrain the formation scenarios, a comparison of stellar
population properties in the very-low-mass regime of M dwarfs
with those, already known, of red giants holds the promise of
shedding light on the origin of this enigmatic phenomenon.
First analyses of multiple stellar populations below the MS
knee have been conducted with data from the Hubble Space
Telescope (HST) for the GC NGC 2808, which already showed
the occurrence of different stellar populations among low-massstars (Milone et al. 2012b). Other GCs were successively
2

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

The footprints of these images are shown in Figure 1, where
we also show the images of ﬁeld B (R.A. ∼ 00h22m36s,
decl. ∼ −72d09m27s) studied by Milone et al. (2023a). The
main properties of the JWST images of ﬁelds A and C and the
HST images of ﬁeld C are summarized in Table 1, whereas werefer to Milone et al. (2023a) for details on the other images ofﬁelds A and B.11 For completeness, we note that the line-ofsight-based rotation axis of 47 Tucanae has inclination ∼ 136°
north to east (Bianchini et al. 2013; Cordoni et al. 2020).
To measure stellar ﬂuxes and positions we used a methodthat is based on two main steps. First-step photometry andastrometry is obtained by using the computer program
img2xym. Originally devised by Jay Anderson (e.g., Anderson& King 2006) for the reduction of HST images, themethodology involves independently measuring stellar ﬂuxes
and positions in each image. This is achieved by employing aspatially variable point-spread function (PSF) model (Anderson& King 2000) along with a “perturbation PSF” that ﬁne-tunes
the ﬁtting process to accommodate slight variations in thetelescope focus. The perturbation PSF is obtained throughunsaturated, bright, and isolated stars. All magnitude determinations derived from individual ﬁlters and cameras have beenstandardized to a common photometric zero-point, aligningwith the zero-point of the deepest exposure in the chosen ﬁlter,
which serves as the reference frame for constructing thephotometric master frame. This alignment was achieved byutilizing bright, unsaturated stars well ﬁtted by the PSF tocalculate the magnitude differences between the master frameand each exposure. The mean of these differences was thenused to transform star measurements in each exposure into thereference frame. For geometric distortion correction, solutions
provided by Bellini et al. (2011), Anderson (2022), and Milone
et al. (2023a) were applied to adjust the stellar positions in the
UVIS/IR, UVIS/WFC3, and NIRCam images. The coordinates of stars in all cluster images were transformed into acommon reference system based on Gaia Data Release 3
catalogs (Gaia Collaboration et al. 2021). This transformation
ensures alignment with the west and north directions for theabscissa and ordinate, respectively.
The second-step photometry and astrometry of all sources
has been carried out with the computer program KS2, which isdeveloped by Jay Anderson and is based on the program
kitchen_sync by Anderson et al. (2008; see Sabbi et al. 2016;
Bellini et al. 2017; Milone et al. 2023b, for details).
KS2 employs three distinct methods for measuring stars, andeach method provides optimal photometry for stars withvarying levels of brightness. In method I, which works wellfor bright stars, the stellar ﬂuxes and positions wereindependently derived in each individual exposure by using
the most suitable effective PSF model for their speciﬁc locationon the detector. The sky brightness is measured over an annularregion between 4 and 8 pixels from the center of the star.
Method II combines information from all exposures to derive
the magnitudes of stars in each exposure by means of aperture
photometry, after subtracting nearby stars. Speciﬁcally, weused a 5 × 5 pixel aperture and determined the sky as inmethod I. This method works well for faint stars, which often

Figure 1. Footprints of the images used in our work on 47 Tucanae. Cyan andazure colors represent the JWST and HST images, respectively, of ﬁeld A,
whereas the JWST and HST footprints of ﬁeld B are colored yellow andorange, respectively. Light and dark green colors respectively indicate JWSTand HST images of ﬁeld C. North is up, and east is to the left.

on very-low-mass stars, including M dwarfs and the browndwarfs, which are poorly explored in GCs. This section
includes an analysis of MPs among M dwarfs, based on the
CMDs. We present in Section 4 early results based on the
spectra of two M dwarfs associated with stellar populations
with very different chemical compositions, and Section 5compares the fraction of 1P and 2P stars among stars withdifferent masses. Finally, Section 6 discusses and summarizes
our results.
2. Data and Data Reduction
To investigate low-mass stars in 47 Tucanae with photometry, we used images collected with the Near-InfraRedcamera (NIRCam) on board JWST, the Wide Field Channel ofthe Advanced Camera for Surveys (WFC/ACS), and theultraviolet and visual channel and the infrared channel of the
Wide Field Camera 3 of HST (UVIS/WFC3 and IR/WFC3).
The data set consists of NIRCam images of a ﬁeld, denoted
as C and positioned approximately 11′ (corresponding to∼3.5 half-light radii, assuming a half-light radius of
3 17; Harris 1996) westward from the cluster center,(R.A. ∼ 00h21m16s, decl. ∼ −72d06m16s), which is observed
as part as GO2560 (PI: A. F. Marino). We acquired images ofﬁeld C simultaneously using the F115W ﬁlter of the shortwavelength channel and the F322W2 ﬁlter of the longwavelength channel. These images were captured using theDEEP8 readout pattern and were dithered to cover the gaps
between the A and B detectors of the short-wavelength channel
effectively. In addition, we used data of a different ﬁeld,
denoted as A (R.A. ∼ 00h22m37s, decl. ∼ −72d04m06s) andlocated about 5′ (∼1.5 half-light radii) westward from thecenter of 47 Tucanae. In particular, we analyzed three regionsof ﬁeld A that have been observed as part of three distinct visitsof JWST (GO2559, PI: I. Caiazzo).

11
In addition to the NIRCam data listed in Table 1, ﬁeld A comprises ACS/
WFC images in F606W and F814W, and IR WFC3 images in F105W, F110W,
F140W, and F160W. Field B data consist of NIRCam F115W and F322W2
images collected as part of GO2560, and ﬁeld B has been observed with HST
in the F110W and F160W ﬁlters of the IR WFC3 camera and in F606W of
UVIS/WFC3 (see Milone et al. 2023a, for further details).

3

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

Figure 2. Left: proper motions for stars in ﬁeld C brighter than mF322W2 = 24.0 mag. Stars in 47 Tucanae, within the red dotted circle, and SMC are clearly separated
in this plot. Right: mF322W2 vs. mF110W − mF160W + mF115W − mF322W2 pseudo-CMD for probable cluster members (black circles), which are located within the red
dotted circle in the left panel and for SMC stars (red crosses).

Table 1
Description of the Images Used in This PaperMISSIONCAMERAFILTERDATE

N × EXPTIME

GO

PI

16 × 857 s
16 × 857 s
16 × 857 s
16 × 857 s
16 × 857 s
16 × 857 s

2559, Visit 1
2559, Visit 1
2559, Visit 2
2559, Visit 2
2559, Visit 3
2559, Visit 3

I. Caiazzo
I. Caiazzo
I. Caiazzo
I. Caiazzo
I. Caiazzo
I. Caiazzo

38 × 1031 s
38 × 1031 s
2 × 50 s + 1347 s + 1398 s
102 s + 174 s + 2 × 1399 s
4 × 299 s + 4 × 1199 s

2560
2560
11677
11677
11677

A. F. Marino
A. F. Marino
H. B. Richer
H. B. Richer
H. B. Richer

Field A
JWST
JWST
JWST
JWST
JWST
JWST

NIRCam
NIRCam
NIRCam
NIRCam
NIRCam
NIRCam

F150W2
F322W2
F150W2
F322W2
F150W2
F322W22022 September 142022 September 142022 September 152022 September 152022 September 152022 September 15
Field C

JWST
JWST
HST
HST
HST

NIRCam
NIRCamUVIS/WFC3
IR/WFC3
IR/WFC3F115W
F322W2
F606W
F110W
F160W

2023 September 24
2023 September 242010 April 102010 April 102010 April 10

Note. For each data set, we indicate the mission (JWST or HST), the camera, the ﬁlter, the date, the exposure times, the GO program, and the principal investigator.
We also indicate the visit numbers of GO2559 observations. The observations collected during visit 1, 2, and 3 are centered around (R.A., decl.) = (0h22m6.0s,−72d03m51s), (0h22m13.s9, −72d05m44s), and (0h23m11.s2, −72d04m25s), respectively.

do not have enough photons to constrain their ﬂuxes andpositions in the individual exposures. Method III is similar tomethod II. The main difference is that the aperture photometry
is computed by considering a circular region with a radius of
0.75 pixels, and the background sky is determined from theannulus located between 2 and 4 pixels away from the initially
identiﬁed position. Subsequently, the multiple measurements
for each star are averaged to obtain the most accurate
estimations of their magnitudes and positions.
We used the parameters that are indicative of the astrometric andphotometric quality provided by the KS2 computer programs to
select the stars that are well ﬁtted by the PSF model. To do this, weused the procedure by Milone et al. (2023b, see their Section 2.4).

Photometry has been calibrated to the Vega system as in Milone
et al. (2023b) and by using the zero-points available in the Space
Telescope Science Institute webpage12 for WFC3 and NIRCam.
We veriﬁed that the photometry is not affected by signiﬁcantreddening variations across the ﬁeld of view (Legnardi et al.
2023). Hence, we did not correct the photometry for differential
reddening.
Proper motions are derived as in Milone et al. (2023a, see their
Section 2.3) by comparing the distortion-corrected positions of
12

https://www.stsci.edu/hst/instrumentation/wfc3/data-analysis/
photometric-calibration; https://jwst-docs.stsci.edu/jwst-near-infrared-camera/
nircam-performance/nircam-absolute-ﬂux-calibration-and-zero-points

4The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

Figure 3. mF115W vs. mF115W − mF322W2 CMD for stars in ﬁeld C with available proper motions (left panel). The middle panel shows F115W magnitude as a functionof proper motion relative to 47 Tucanae. The vertical aqua line separates the bulk of cluster members (black points in the left and middle panels) from ﬁeld stars (redcrosses). Stars below the aqua horizontal line, where the proper motions do not allow us to disentangle ﬁeld stars from cluster members, are represented with graycrosses. Probable white dwarfs, selected from the mF606W vs. mF606W − mF115W CMD are represented with blue triangles in the left and middle panels. In the rightpanel we show the mF322W2 vs. mF115W − mF322W2 CMD for all stars in the ﬁeld CMD (gray points) with the stars simulated by the trilegal Galactic model (red starredsymbols; Girardi et al. 2005) for a ﬁeld with the same area as ﬁeld C in the direction of 47 Tucanae.

stars measured at different epochs. The proper-motion diagram
derived for ﬁeld-C stars brighter than mF322W2 = 24.0 is plotted in
the left panel of Figure 2, showcasing two primary stellar clumps
encompassing the majority of cluster members and Small
Magellanic Cloud (SMC) stars. The red circle is used to
distinguish probable 47 Tucanae members from ﬁeld stars, which
are colored black and red, respectively, in the mF322W2 versus
mF110W − mF160W + mF115W − mF322W2 pseudo-CMD.
The mF115W versus mF115W − mF322W2 CMD of all stars withavailable proper motions is plotted in the left panel of Figure 3.
We computed for each star the proper motion relative to theaverage proper motion of 47 Tucanae, μR, and plotted thisquantity against the F115W magnitude. For magnitudes brighter
than mF115W = 26.1 mag the stars of 47 Tucanae and the SMC
exhibit distinct proper motions. Hence, we draw by eye thevertical aqua line to separate the bulk of cluster members from theﬁeld stars, which are colored black and red, respectively in the left
and middle panels of Figure 2. Due to the large observational
errors, the proper motions of stars fainter than mF115W = 26.1 mag
do not allow us to disentangle ﬁeld stars from cluster members.
To estimate the amount of faint ﬁeld stars that we expect
in ﬁeld C, we compare the observed mF322W2 versus
mF115W − mF322W2 CMD of all stars with available NIRCamphotometry and the simulated CMD calculated by the Trilegal
code (Girardi et al. 2005) for a Galactic ﬁeld with the same area
and the same Galactic coordinates as ﬁeld C. The result is plottedin the right panel of Figure 3 where we observe that a negligible
fraction of Milky Way ﬁeld star interlopers are expected tocontaminate the observed CMD sequences of 47 Tucanae.

2.1. Artiﬁcial Stars and CompletenessWe conducted artiﬁcial star (AS) experiments to estimate thephotometric uncertainties and the completeness level of oursample by using the procedure adopted in previous papers (e.g.,Anderson et al. 2008). These stars were distributed across theﬁeld of view similarly to the cluster stars and along the ﬁducial
lines of the MS of 47 Tucanae that we derived from theobserved CMDs. The ASs have been reduced by using exactly
the same procedure adopted for the real stars. The KS2computer program generates the same diagnostic measurementsof photometric and astrometric quality for ASs as it does forreal stars. In our analysis, we considered only a subset of ASs
that are relatively isolated, exhibit good PSF ﬁtting, and show
small rms values in magnitude. These ASs were selected using
the same criteria we applied to real stars in our investigation.
We partitioned the NIRCam ﬁeld for each ﬁeld into ﬁve
concentric annuli, centered on 47 Tucanae. Within eachannulus, we analyzed the results of the AS experiments in N
distinct 0.5 mag bins, spanning from the saturation limit toapproximately one F322W2 magnitude below the faintest starthat we detected. For each of the 5 × N grid positions, we
calculated the average completeness by comparing the numberof recovered ASs with the input ASs within that particular bin.
2.2. NIRSpec Spectra
NIRSpec has been used in the multiobject spectroscopy(MOS) mode, involving a microshutter assembly conﬁguration,
which allowed us to collect 29 source spectra of sufﬁcient
5

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

quality simultaneously, within a 3.¢ 4 ´ 3.¢ 6 ﬁeld of view
centered on ﬁeld B. We choose the G235M/F170LP disperser
and ﬁlter combination, which observes the wavelength range
1.66–3.07 μm, at a nominal resolving power ∼ 1000.
The improved reference sampling and subtraction (IRS2)
readout mode, with the NRSIRS2 pattern, which has ﬁve
frames averaged into a single group, has been employed for our
MOS observations. Each source has been observed with a twoshutter nod in slitlet pattern with two identical conﬁgurations
and each conﬁguration has been executed 10 times. Each
exposure was observed for 1182 s (with 16 groups per
integration), for a total on-target time of 47,280 s. However, adither of 20″ in the dispersion direction has been applied toclose the detector gap. In the end, of our 29 observed stars,
17 are covered by both dither positions, and hence have all the
40 exposures, while 12 have only one dither position(20 exposures).
We have analyzed the NIRSpec 1D extracted spectra
processed with the JWST Science Calibration Pipeline (Bushouse et al. 2023). While for a full analysis of the individual
stars we refer to a forthcoming paper devoted to oxygen
estimates, we present in this paper the results for spectra of
M dwarfs with extreme chemical compositions.

To investigate further the lower MS of 47 Tucanae, we show inthe left panel of Figure 5 a mF150W2 versus mF150W2 − mF322W2
CMD of stars in ﬁeld A. As highlighted by the Hess diagram
plotted in the inset, we conﬁrm the main gap and bend in the
CMD already evident in Figure 4. In particular, this diagram
clearly shows the change in the slope of the sequence of ultracool
stars at mF322W2 = 24.5 mag. The right panels of Figure 5
compare the mF322W2 versus mF150W2 − mF322W2 (top) and themF322W2 versus mF115W − mF322W2 CMDs (bottom) zoomed inaround the bottom of the MS, which are derived from GO2559
and GO2560 data. The dashed lines, derived by eye, enclose thebulk of stars along the MS and the brown-dwarf sequence, which
we marked with black circles. The colored symbols mark the stars
for which photometry from both GO2559 and GO2560 isavailable.
The luminosity functions and MFs may provide additional
information of low-mass stars in 47 Tucanae. Unfortunately,isochrones that reproduce both massive and low-mass stars
in 47 Tucanae are not available to us as shown in the left
panel of Figure 6, where we compare the mF322W2 versus
mF115W − mF322W2 CMD with three isochrones. We have marked
with black dots the probable cluster MS and brown dwarfs that weselected by eye, whereas the white dwarfs of 47 Tucanae and the
SMC are colored gray. The red and the blue isochrones provide
the best ﬁt of 1P and 2P stars with extreme chemical composition,
respectively, and comprise stars more massive than ∼0.1 Me.
These isochrones, which refer speciﬁcally to 47 Tucanae and aredescribed in detail by Milone et al. (2023a), are extracted fromthe Dartmouth database (Dotter et al. 2008) and share the sameage (13 Gyr), iron abundance ([Fe/H] = −0.75), and [α/Fe] =+0.4 dex. However, they have different abundances of He, C, N,
and O, in such a way that the blue isochrone is enhanced inhelium mass fraction by ΔY = 0.04 and in [N/Fe] by 1.2 dex,
and depleted in [C/Fe] and [O/Fe] by −0.35 and 0.50 dex,
respectively, relative to the red isochrone that has Y = 0.248,[O/Fe] = 0.4, and solar carbon and nitrogen abundances.
To derive the mass–luminosity relation for stars less massive
than ∼0.1 Me we tentatively used the tracks by Phillips et al.
(2020; brown isochrone plotted in the left panel of Figure 6),
which however are based on model atmospheres and evolution
models for brown dwarfs and giant exoplanets with solar
metallicity. Clearly, this isochrone, which is computed for anage of 10 Gyr, does not match the brown-dwarf sequence of
47 Tucanae, and differs from those derived by Gerasimov et al.
(2024) that qualitatively reproduce the observed CMDs. In a
recent paper, Nardiello et al. (2023) identiﬁed 10 candidate
brown dwarfs by using a mF322W2 versus mF115W − mF322W2
CMD of ﬁeld B stars in 47 Tucanae. These stars are distributed
along the isochrone of Phillips et al. (2020) and span F322W2
magnitudes between ∼25.2 and 25.8 mag, which are superimposed on the SMC MS, and may belong to the SMC rather
than to 47 Tucanae. Accurate stellar proper motions, which aremandatory to disentangle SMC stars and possible cluster
members, are not available for these very faint stars. Hence, weneither exclude nor conﬁrm the presence of brown dwarfs in
47 Tucanae following the Phillips et al. (2020) isochrone. In thefollowing, we only focus on the sequence of ultracool stars
analyzed in this paper.
The luminosity functions of 47 Tucanae stars are plotted in the
top-right and middle-right panels of Figure 6 for the stars
observed with detectors A and B of the long NIRCam channels
and for all stars together. Speciﬁcally, we plotted the logarithm of

3. The NIRCam Color–Magnitude Diagram of 47 Tucanae
The mF322W2 versus mF115W − mF322W2 CMD of stars in ﬁeld
C is illustrated in the lower panels of Figure 4, whereas theupper panels show portions of the NIRCam ﬁeld of view. We
notice that the mF115W − mF322W2 color broadening of theupper MS of 47 Tucanae (mF322W2  18.7) is comparable withthe broadening due to the observational errors alone. In
contrast, the MS color spread for stars fainter than the MS knee
is much wider than the observational errors. This phenomenon,
which is due to the presence of MPs among the M dwarfs, will
be further investigated in Section 3.3. In the following
Section 3.1 we investigate fainter stars at magnitudes
approaching the H-burning limit.
3.1. Ultracool Dwarfs
We analyze here the faintest 47 Tucanae stars appearing in the
mF322W2 versus mF115W − mF322W2 CMD of Figure 4. In
particular, we focus on the faintest portion of the CMD,
composed by a plume of stars extending from the bottom of thevertical sequence (at mF322W2 ∼ 23.3 mag) down tomF322W2 ∼ 27.0 mag. This sparsely populated group of stars
covers a wide color range (∼2 mag) and it appears as anextension of the bluest portion of the vertical sequence
composed by the cluster’s M dwarfs. A quick glance at thissequence suggests that its stars, which have masses less than∼0.1 Me, are not continuously distributed. In particular, we
notice the presence of compelling features, namely, possible
clumps at mF322W2 ∼ 23.3, a sharp gap ∼24.2 mag, and a distinct
change in slope at ∼24.5 mag (see also Figures 4 and 5).
The observed luminosity function approaches its minimum
value around mF322W2 ∼ 25.3, and rises up toward fainterluminosities, where we observe the bright portion of the browndwarf sequence. As an example of the analyzed ﬁeld of view,
the upper-right panels of Figure 4 show a zoom in view of thestacked F322W2 and F115W images that include two probablebrown dwarfs (red and blue circles).
6

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

Figure 4. Top: three-color (RGB) image of a portion of the studied ﬁeld of view (ﬁeld C). We used the stacked F115W and F322W2 images for the B and R channels,
respectively, whereas the G channel is obtained by averaging the F115W and F322W2 stacked images (left). The right panels show the F322W2 (top) and F115W(bottom) stacked images for the region within the yellow box highlighted in the ﬁeld of view on the left panel. Bottom: mF322W2 vs. mF115W − mF322W2 CMD of stars
in ﬁeld C (left). The arrows indicate the location of the possible gaps or discontinuities along the sequence of ultracool 47 Tucanae stars. Stars fainter thanmF322W2 ∼ 25.3 mag and with colors redder than mF115W − mF322W2 ∼ 1.5 mag are probable brown dwarfs. The blue and red circles in the CMD highlight two ofthem, which are also marked in the top-right panel. We choose these two stars, among the brown-dwarf sample, for illustrative purposes, as they are located in a
relatively small region of the ﬁeld of view. A zoom around the MS region of 47 Tucanae populated by M dwarfs is shown in the right panel.

into stellar masses for stars brighter than mF322W2 ∼ 22.5,
which have masses larger than ∼0.1 Me, In this mass interval,
we ﬁnd that the MF exhibits a peak at ∼0.5 Me and declines
toward lower masses.
Due to the lack of appropriate isochrones for stars with
masses lower than ∼0.1 Me, it is not possible to provide
accurate MFs for these stars. Nevertheless, we assumed that theluminosity corresponding to mF322W2 = 25.3 corresponds to astellar mass of 0.074 Me, which according to Gerasimov et al.
(2024) corresponds to the hydrogen-burning limit, whereasstars with mF322W2 = 27 have 0.06 Me, as predicted by
Gerasimov et al. (2024). The results are plotted in thebottom-right panel of Figure 6, which show an abrupt change

the numbers of stars per square arcminute, corrected forcompleteness, as a function of mF322W2 magnitude. The star
counts are calculated in 0.4 mag wide bins. The luminosity
function exhibits a peak around mF322W2 ∼ 20.2. As wemove toward fainter luminosities, the star count per unit
magnitude consistently decreases, reaching a minimum aroundmF322W2 ∼ 25.3 mag. Beyond this point, the number of stars
increases for fainter values of mF322W2.
The bottom-right panel of Figure 6 shows a tentative
determination of the MF, which is derived from the luminosity
function plotted in the middle-right panel. For simplicity, we
did not account for the MPs in 47 Tucanae, and we used the redisochrone plotted in the left panel to convert the luminosities
7

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

Figure 5. Left: mF322W2 vs. mF150W2 − mF322W2 CMD of stars in ﬁeld A (left). The inset shows a Hess diagram of the low MS of 47 Tucanae, where the main gaps anddiscontinuities are marked with colored arrows. Speciﬁcally the crimson arrows indicate the changes in slope at mF322W2 around 23.3 and 24.5 mag, while the tealarrows correspond to the sharp gap at mF322W2 ∼ 24.2, and the observed minimum in the luminosity function at mF322W2 ∼ 25.3 mag. Right: mF322W2 vs.
mF150W2 − mF322W2 (top) and mF322W2 vs. mF115W − mF322W2 CMDs (bottom) for faint MS stars and brown dwarfs. The stars observed in both CMDs are represented
with colored symbols; speciﬁcally, the aqua, orange, and red colors indicate the dwarfs that populate the main MS segments, whereas the likely brown dwarfs
observed in both data sets are colored brown.

limit in 47 Tucanae? Is it at mF322W2 ∼ 24.2, at 25.3, or at 26?
What would be needed is an extended set of isochrones,
possibly with options in the input physics. The data we present
here represent a major jump in depth and completeness forcoeval M and brown dwarfs all at the same distance andextending down to unprecedented faint luminosities. These
data demand a new, parallel effort in modeling M dwarfs andbrown dwarfs.
In the meantime, we can still speculate on the nature of of the
gap at mF322W2 ∼ 24.2 and the origin of the different
luminosities at the minimum of the luminosity function, justmentioned above. If the gap does not correspond to thehydrogen-burning limit, what else can cause it? These M
dwarfs are fully convective, so convection effects could hardly
result in such a sharp discontinuity (note that a drop is notevident in the luminosity function at this magnitude because of
the binning, except in the case of the blue points in Figure 6).
Could atmospheric effects cause it? We do not know, so it
remains an unsolved issue.
Similarly, we cannot unambiguously pinpoint the origin of
the magnitude discrepancy in the minimum of the luminosity
function. Both sets use Vega magnitudes, so it cannot be ain the MF slope at the assumed hydrogen-burning limit. We are
reluctant to take for good this result, because of the tentative
mass–luminosity relation we have adopted.
3.2. The M Dwarf to Brown Dwarf Transition in 47 Tucanae
As emphasized by Gerasimov et al. (2024), the terminationof the hydrogen-burning MS is expected to exhibit a gap in the
CMD between the faintest MS stars and the brightest browndwarfs. Thus, it is tantalizing to identify this gap (i.e., thetermination of the MS) with the evident gap at mF322W2 ∼ 24.2
seen in Figure 5. In their simulations Gerasimov et al. (2024)
do not ﬁnd a similar gap at this luminosity (their Figure 6),
whereas a gap may appear at a much fainter luminosity,
mF322W2 ∼ 26 (their Figure 10). This corresponds to theminimum in their synthetic luminosity function at this latter
magnitude, as illustrated in their Figure 9. Thus, Gerasimov
et al. (2024) associate this minimum with the hydrogen-burning
limit.
Our observed luminosity function shows also a sharp
minimum (see Figure 6), but it is found at mF322W2 ∼ 25.3,
i.e., around half a magnitude brighter than predicted by
Gerasimov et al. (2024). So, where is the hydrogen-burning
8

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

Figure 6. Left: reproduction of the mF322W2 vs. mF115W − mF322W2 CMD of Figure 4, where we used black colors to mark the probable MS and brown dwarfs of
47 Tucanae. We superimposed on the CMD the best-ﬁt isochrones for MS stars (Dotter et al. 2008; Milone et al. 2023a) and the isochrone of Phillips et al. (2020) for
low-mass stars. The horizontal lines mark the magnitudes corresponding to four stellar masses inferred from the red isochrone and from Gerasimov et al. (2024).
Right: F322W2 luminosity function of 47 Tucanae stars along the MS and the brown-dwarf sequence obtained by using 0.4 mag wide bins. The open circles arederived from stars where the correction for the incompleteness of the photometric sample is smaller than 50%. The top panel shows the luminosity functions of stars in
detectors A and B of the long-wavelength channel of NIRCam, separately, which have different radial distances ∼ 11 5 and 14 5 from the cluster center. The middle
panel shows the results from the entire ﬁeld of view. The gray points in the middle panel are derived by using 0.4 mag wide bins but by changing the starting
magnitude value of the luminosity function by 0.1, 0.2, and 0.3 mag with respect to the value used for the black points. The bottom panel shows the MF for stars in theentire ﬁeld of view. The vertical dotted line corresponds to 0.11 Me. See the text for details.

Panel (c) shows the histograms and the kernel-density
distributions of δ(mF115W − mF322W2) in six intervals of
F322W2 magnitude. The latter is obtained by using a Gaussian
kernel with dispersion σ = 0.02 mag.
When considering a MS portion with nearly constant
F322W2 magnitude, the majority of M dwarfs in thisluminosity interval exhibit blue colors. In particular, thehistograms of Figure 7(c) exhibit a peak aroundδ(mF115W − mF322W2) = 0.02 mag, which corresponds to thebulk of 1P stars. However, there is a tail of stars that extends
toward red colors.
The color broadening is mostly due to the star-to-star oxygen
variations, which are associated with multiple stellar populations.
The main culprit is the absorption of different chemical species
containing oxygen, mainly H2O (the primary absorber) and OH.
The absorption features of these molecules strongly affect the ﬂux
in the F322W2 band whereas the F115W ﬁlter is insensitive tooxygen variations. As a consequence, the 1P stars, which havehigher oxygen abundances than the 2Ps, exhibit fainter F322W2
magnitudes and bluer mF115W − mF322W2 colors than stars withsimilar atmospheric parameters (Milone et al. 2012b, 2023a;
Dotter et al. 2015; VandenBerg et al. 2022; Ziliotto et al. 2023).
To quantify the amount of oxygen that is needed to reproduce
the mF115W − mF322W2 MS broadening, we compared isochrones
with different oxygen abundances (Dotter et al. 2008; Milone et al.
2023a) with the observed CMD. Speciﬁcally, we considered theisochrones with [O/Fe] = 0.4 and [O/Fe] = −0.1. To compare thezero-point issue. Maybe it is due to a systematic effect on thestellar models, in a regime in which a small difference in mass
can result in a large difference in luminosity. On theobservational side, at these faint magnitudes the luminosity
function of galaxies is raising steeply, so we cannot exclude
that our luminosity function of 47 Tucanae is partly contaminated by faint galaxies. Still we have excluded morphologically nonstellar objects in constructing the luminosity
function of 47 Tucanae.
3.3. Multiple Stellar Populations
The mF322W2 versus mF115W − mF322W2 CMD in Figure 4
clearly shows that the MS color broadening, which iscomparable with that expected from observational errors forstars brighter than the MS knee, suddenly increases among M
dwarfs fainter than the MS knee, and approaches a value ofmore than 0.2 mag in the F322W2 interval between ∼20 and
23 mag.
Figure 7 provides additional information on the color
distribution of MS stars. We derived by eye the aquadashed–dotted line shown in panel (a) to deﬁne the blue
boundary of the MS. Subsequently, we utilized this line togenerate the verticalized mF322W2 versus δ(mF115W − mF322W2)
diagram represented in panel (b). To achieve this, we
subtracted the color of each star from the color of the aqua
ﬁducial associated with the identical F322W2 magnitude.
9The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

Figure 7. Reproduction of the mF322W2 vs. mF115W − mF322W2 CMD of Figure 4 zoomed in around the bottom of the MS of 47 Tucanae (panel (a)). The aqua line isthe blue boundary of the MS and is used to derive the mF322W2 vs. δ(mF115W − mF322W2) verticalized diagram shown in panel (b). The red and blue lines are the best-ﬁt
isochrones for MS stars with [O/Fe] = 0.4 and [O/Fe] = −0.1 dex, respectively (Dotter et al. 2008; Milone et al. 2023a). Panel (c) illustrates the δ(mF115W − mF322W2) histogram distributions for stars in six F322W2 magnitude intervals and the corresponding kernel-density distributions.

Figure 5 and the CMD simulated by Gerasimov et al. (2024, see
their Figure 10), which accounts for the chemical compositions of
the MPs of 47 Tucanae. The simulated CMD exhibits a wide
mF150W2 − mF322W2 color range ∼0.5 mag for stars with
mF322W2 ∼ 24 mag, which seems in contrast with the narrow
sequence of ultracool stars with similar luminosity observed in
Figure 5. This fact suggests a possible lack of very O-poor 2Pstars among ultracool stars.
Additional information on MPs in 47 Tucanae is provided by
the mF814W versus CF606W,F814W,F322W2 pseudo-CMD of the stars
in ﬁeld B shown in the top panel of Figure 8. In this diagram, thestars fainter than the MS knee, identiﬁed at approximately
mF814W = 20.0 mag, reveal hints of four discrete sequences.
According to the isochrones provided by Dotter et al. (2008) and
Milone et al. (2023a), which consider the chemical composition ofMPs, the reddest sequence is composed of 1P stars, while the 2Pstars with particularly extreme chemical compositions display
bluer CF606W,F814W,F322W2 pseudocolors.
To maximize the information on MPs, we combine themF814W versus CF606W,F814W,F322W2 pseudo-CMD and themF814W versus mF606W − mF814W CMD of M dwarfs with21.5 < mF814W < 24.0 mag. The inset of Figure 8 shows theresulting ΔCF606W,F814W,F322W2 versus ΔF606W,F814W chromosome map (ChM), where 1P stars are clustered in the ChM

relative colors of the 47 Tucanae stars with those of the isochrones
properly, we have calculated the color difference between the
O-poor isochrone and the O-rich one and shifted both verticalized
isochrones by δ(mF115W − mF322W2) = 0.02 mag, in such a waythat the red isochrone is superimposed on the bulk of 1P stars. The
δ(mF115W − mF322W2) interval spanned by stars brighter thanmF322W2 ∼ 22.5 mag, which have masses larger than ∼0.1 Me, isconsistent with an oxygen variation ∼0.5 dex, comparable with the[O/Fe] range inferred for RGB stars from high-resolution
spectroscopy (e.g., Carretta et al. 2009; Dobrovolskas et al.
2014). We conclude that the stars in the mass interval between∼0.1 and 0.9 Me span a similar range of [O/Fe].
Intriguingly, as noticed by Milone et al. (2023a) for the stars inﬁeld B, the MS color broadening of stars fainter thanmF322W2 ∼ 23 mag appears to be much smaller than that observed
for brighter M dwarfs in the F322W2 magnitude range∼20.0–22.5. Unfortunately, there are no available isochrones that
account for multiple stellar populations among ultracool stars.
Hence, we cannot provide a ﬁrm conclusion on whether or not
this fact is due to the lack of stellar populations with extreme
chemical composition among these very-low-mass stars. Additional qualitative information on MPs among stars less massive
than ∼0.1 Me is provided by the visual comparison between theobserved mF322W2 versus mF150W2 − mF322W2 CMD shown in
10

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

Figure 8. Top: mF814W vs. CF606W,F814W,F322W2 pseudo-CMD for stars in ﬁeld B. The inset shows the ΔCF606W,F814W,F322W2 vs. ΔF606W,F814W ChM for MS stars with21.5 < mF814W < 24.0 mag, and the corresponding ΔCF606W,F814W,F322W2 kernel-density distribution. Bottom: NIRSpec spectra for the two stars marked with large
dots in the pseudo-CMD and in the ChM (left). Two simulated spectra with [O/Fe] = +0.3 (red) and [O/Fe] = −0.1 (blue) are superimposed on the observed spectrain the bottom-right panel.

region with CF606W,F814W,F322W2  0.05 mag. We detect anextended 1P sequence that we associate with star-to-starmetallicity variations, as found by Marino et al. (2023) basedon high-resolution spectroscopy of RGB stars (see also Milone
et al. 2017b; Legnardi et al. 2022; Milone et al. 2023a). The 2P
sequence exhibits hints of four groups of stars clustered around
ΔF606W,F814W ∼ −0.10, −0.14, −0.17, and −0.20 mag.

comparable F814W magnitudes. However, these stars exhibit
distinct values for ΔCF606W,F814W,F322W2 and ΔF606W,F814W.
Based on their location in Figure 8, upper panel, the 1P and 2Pstars are labeled as red dots and blue dots, respectively.
The two targets have been spectroscopically observed as part of
GO2560, and their NIRSpec spectra are displayed in the bottomleft panel of Figure 8. The red spectrum represents the star with
R.A. = 0h22m22.14s and decl. = −72d04m10.s5 while the blue
spectrum correspond to the star with R.A. = 0h22m12.36s anddecl. = −72d04m30.s8. In presenting the spectra, we applied thesame shift to both, plotting them such that the average value of−2.5 log10(ﬂux) in the region with λ < 1.75 μm for the blue
spectrum equals zero.
Clearly, the spectrum of the 1P star exhibits much stronger
molecular bands than that of the 2P star. In principle, such large
differences in the spectra could be either due to different

4. The Chemical Composition of M Dwarfs from NIRSpec
SpectraIn this section we present NIRSpec spectra for two M dwarfs
photometrically associated with different stellar populations.
The location of these stars in the photometric diagrams isshown in Figure 8. The large dots superimposed on the pseudoCMD and the ChM of Figure 8 mark indeed two stars with
11

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

the 2P deﬁnes a sequence ranging from (ΔF606W,F110W,
ΔF110W,F160W,F115W,F322W2) ∼ (−0.2, 0.1) toward the top left
corner of the ChM.
To assess the proportion of stars within each population, we
expanded upon the approach outlined by Zennaro et al. (2019)
to derive the fraction of stars in the four stellar populations of
the GC NGC 2419 and apply it to the ChM of 47 Tucanae (seealso Milone et al. 2012a; Nardiello et al. 2018; Milone et al.
2020). In a nutshell, we estimated the mean values of
ΔF606W,F110W and ΔF110W,F160W,F115W,F322W2 for the stars
within each population (depicted as colored dots in Figure 9).
In the case of 1P stars, we used two values for the center toaccount for subpopulations with different metallicity (Legnardi
et al. 2022; Marino et al. 2023). These values served to deﬁne
ﬁve distinct regions denoted as R1, R2A, R2B, R2C, and R2D anddelimited by the colored solid lines. Owing to photometricerrors, each region may encompass stars from all populations.
As an example, the observed count of stars within region R1 is:

atmospheric parameters or to different chemical compositions.
The location of the two stars at similar mF814W magnitudes
suggests that they have similar atmospheric parameters. Indeed,by ﬁtting the mF814W versus mF606W − mF814W CMD with theisochrones of Dotter et al. (2008) and Milone et al. (2023a) we
obtain Teff/log g = 3700/5.0 and 3600/5.0 for the blue and thered star, respectively. Such a small difference in the atmospheric parameters alone cannot account for the differences ofthe observed spectra.
As outlined in the spectra of Figure 8, the main molecularfeatures appearing in the observed spectral range are due torotational–vibrational H2O band heads at 1.9 and 2.7 μm, andsome contribution from the CO second-overtone bands arerather masked by the H2O bands. The water vapor molecular
bands in the near-infrared are good indicators of oxygen
abundance.
While for a full detailed spectroscopic analysis of all theobserved NIRSpec spectra we refer the reader to an upcoming
paper, here we illustrate two synthetic spectra constructed withthe atmospheric parameters assumed for our two M dwarfs toinfer a ﬁrst estimate of the O range in the low-mass regime of
47 Tucanae. The bottom-right panel of Figure 8 shows acomparison of the observed spectra with simulated spectra withdifferent light-element abundances. The simulated spectra arederived as in our previous works (Milone et al. 2023a; Ziliotto
et al. 2023). In a nutshell, we calculated the model atmospheres
using the ATLAS 12 computer program, employing the
opacity-sampling technique and assuming local thermodynamic equilibrium (Kurucz 1970; Sbordone et al. 2004). We
incorporated molecular line lists for all diatomic molecules
listed on Kurucz’s website, in addition to including H2O
molecules from Partridge & Schwenke (1997). We used the
SYNTHE computer program (Kurucz & Avrett 1981;
Castelli 2005; Kurucz 2005; Sbordone et al. 2007) to computethe spectra in the region between 1.7 and 3.2 μm covered by
the available NIRSpec spectra. Speciﬁcally, the crimson
spectrum has oxygen content [O/Fe] = 0.3 dex, whereas theteal one has [O/Fe] = −0.1 dex. Noticeably, the blue spectrum
has higher effective temperature than the red one by 70 K. Thiseffective temperature difference is consistent with the fact thatthe 2P stars of 47 Tucanae are helium enhanced with respect tothe 1Ps (e.g., Milone et al. 2018b; Lagioia et al. 2018).
We conclude that the comparison of the spectra of the two
M dwarfs belonging to different stellar populations indicatesthat M dwarfs exhibit a range in O abundances very similar tothat reported in the literature for higher-mass stars along the
RGB. The result presented here is the ﬁrst detection of the O
range directly performed on stellar spectra of M dwarfs, andstrongly corroborates our previous ﬁndings about the similarityof the multiple stellar populations in stars with different masses(from the RGB to M dwarfs), already suggested by photometric
diagrams.

R1
R1
R1
NR1 = N1P f1P+ N2PA f2PA
+ N2PB f2PB
R1
R1
+ N2PC f2PC
+ N2PD f2PD,

(1 )where N1P, N2PA, N2PB, N2PC, and N2PD, are, respectively, thenumbers of 1P, 2PA, 2PB, 2PC, and 2PD stars in each
population in our sample.
The number of stars within the regions R2A, R2B, R2C, andR2D are linked to the fraction of stars of each population
through four comparable equations. The values of NR1, NR2A,
NR2B, NR2C, and NR2D used in these equations are derived bycounting the stars within the corresponding region.
The fractions of 1P, 2PA, 2PB, 2PC, and 2PD stars within the
R1
R1
R1
R1
region R1, f1PR1, f2PA
, f2PB, f2PC, and f2PD
are derived fromsimulated ChMs composed of ASs, and we did the same forinferring the fractions of stars of each populations in theregions of the ChM R2A, R2B, R2C, and R2D. To do this,we simulated 50,000 ASs for each population, disposed along
the centers of each population in the ChM.
The numbers of stars in the ﬁve populations of 47 Tucanae,
N1P, N2PA, N2PB, N2PC, N2PD, are calculated by solving forthese ﬁve equations. The results are provided in Table 2, where
we also provide the fractions of stars in the ﬁve stellar
populations that we obtain by extending the same procedure to
the ChM shown in Figure 8 for stars in ﬁeld A. In particular, we
ﬁnd that 1P stars comprise 38.1% ± 1.2% and 46.3% ± 4.2 %of the total number of ChM stars of ﬁeld A and ﬁeld C,
respectively.
The fraction of 2P stars in 47 Tucanae inferred from either
RGB or HB stars signiﬁcantly changes as a function of radial
distance. It ranges from about 80% near the cluster center toless than 60% for distances larger than ∼5′ (Milone et al.
2012b, 2017b; Cordero et al. 2014; Dondoglio et al. 2021;
Lee 2022). Figure 10 compares the fractions of 2P stars derived
by Dondoglio et al. (2021) from HB stars with different radial
distances (black points) and those obtained in this paper from
M dwarfs (red triangles). We conclude that there is no
signiﬁcant difference between the fractions of 2P stars derived
from stars with different masses. In the context of the multiple
generations scenario, where the 2P stars form in the densecluster center, this would imply that the initial mass function
does not depend on the density of the environment.

5. Population Ratios
To derive the fraction of stars in the main stellar populations of
47 Tucanae we started using the ΔF110W,F160W,F115W,F322W2
versus ΔF606W,F110W ChM of stars in ﬁeld C, which is derived
from the mF322W2 versus mF110W − mF160W + mF115W − mF322W2
pseudo-CMD plotted in Figure 2 and the mF322W2 versus
mF606W − mF110W CMD.
The results are plotted in the left panel of Figure 9. The
1P stars are clustered around the origin of the ChM, while
12

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

Figure 9. ΔF110W,F160W,F115W,F322W2 vs. ΔF606W,F110W ChM of M dwarfs in ﬁeld C (panel (a)). Panels (b)–(e) illustrate the procedure used to estimate the fractions of
1P and 2P stars. The observed ChM is further reproduced in panel (b), while panel (c) shows the simulated ChM where simulated 1P stars are colored in red. The
colored solid lines delimit the four regions used to infer the fraction of stars in each population, while the colored dots mark the average positions in the ChM of
the simulated stellar populations. Panels (d) and (e) show the ΔF110W,F160W,F115W,F322W2 and ΔF606W,F110W kernel distributions for the observed (black line) and thesimulated (gray line) stars.

Table 2Fractions of Stars in the Five Stellar Populations of 47 Tucanae in Fields A and C
Field

1P

2PA

2PB

2PC

2PD

A
C

0.381 ± 0.012
0.463 ± 0.042

0.352 ± 0.012
0.329 ± 0.035

0.148 ± 0.009
0.122 ± 0.028

0.065 ± 0.006
0.049 ± 0.021

0.054 ± 0.006
0.036 ± 0.014

devoted to the analysis of multiple stellar populations in GCs
conducted with the JWST. However, in addition to studies onMPs, we demonstrated that our data set has the potential to be
used to investigate the general properties of poorly exploredstellar populations, such as M dwarfs and brown dwarfs. In thiswork, we present early results which allowed us to study lowmass stars, in the domain of M dwarfs and beyond, down to thehydrogen-burning limit and the brown-dwarf sequence.
This overview sheds light on the MPs among the low-massstars of 47 Tucanae. The main results can be summarized asfollows.
1. The photometric diagrams constructed with the F322W2
ﬁlter of NIRCam, including the mF322W2 versus
mF115W − mF322W2 CMD and the mF322W2 versus
CF606W,F814W,F322W2 pseudo-CMD, reveal that, below
the knee, MS stars with masses larger than ∼0.1 Me andsimilar F322W2 luminosities span a wide color, or
pseudocolor, range. This result corroborates the evidenceof multiple stellar populations among M dwarfs (see, e.g.,
Dondoglio et al. 2022; Milone & Marino 2022; Milone
et al. 2023a; Cadelano et al. 2023, for previous
photometric studies of MPs among very-low-mass stars
in 47 Tucanae).
2. The comparison between isochrones with different
chemical compositions and the observed CMD reveals
that the MS color broadening observed among M dwarfs
more massive than ∼0.1 Me is consistent with stellar
populations with different oxygen abundances. The[O/Fe] range ∼0.5 dex that is needed to reproduce the

Figure 10. Fraction of 2P stars as a function of the radial distance from thecluster center. The dlack dots are derived by Dondoglio et al. (2021) by using
HB stars, whereas the results obtained in this paper from M dwarfs are
represented with red triangles. The horizontal lines mark the radial interval
associated with each point. The dotted and dashed vertical lines mark the core
and the half-light radius from the 2010 version of the Harris (1996) catalog.

6. Summary and Conclusions
We have presented our JWST project on multiple stellar
populations in the GC 47 Tucanae. Our main data set
comprised ∼13 hr of observing time with NIRSpec + NIRCamparallel observations, allowing us to infer crucial constraints
on the MP phenomenon both from spectroscopic andphotometric analysis. This is the ﬁrst project speciﬁcally
13

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

observations is comparable with that observed among
RGB stars by means of high-resolution spectroscopy. We
notice that the mF115W −mF322W2 color broadening of very
faint stars (mF322W2 > 23) appears to be narrower than thatobserved among the M dwarfs with brighter luminosity,
similar with what was observed by Milone et al. (2023a) forstars in ﬁeld B. An appropriate comparison with isochrones
that account for the chemical compositions of MPs isneeded to associate this fact with the possible lack of stellar
populations with extreme chemical compositions among
ultracool stars.
3. We present the NIRSpec spectra for two stars with similar
F814W magnitudes that occupy extreme positions in the
ChM and are associated with the populations 1P and 2PCstars. For a ﬁxed wavelength, the 1P star spectrum ismore absorbed than the spectrum of the 2PC star. The
comparison with synthetic spectra with different chemicalcompositions reveals that the two stars have different
oxygen abundances, with the 2PC M dwarf being
depleted by [O/Fe] = 0.4 dex with respect to the 1P
star. Hence, the ﬂux difference can be attributed tomolecules consisting of oxygen (predominantly H2O),
which exhibit strong absorption in the spectra of 1P stars,
known for their oxygen-rich composition. This outcome
represents the ﬁrst spectroscopic determination of thechemical composition of M dwarfs within a GC. It
validates the earlier prediction, derived from synthetic
spectra, that the observed color variation among
M dwarfs in GCs is attributable to multiple stellar
populations characterized by varying oxygen abundances(Milone et al. 2012b).
4. The 2P M dwarfs consist ∼62% and 54% of the total
number of M dwarfs, respectively at ∼5′ and ∼11′ fromthe cluster center. These values are similar to the fractionof 2P stars measured among HB and RGB stars withsimilar radial distances (Milone et al. 2012b; Dondoglio
et al. 2021), thus indicating that the fraction of 2P stars
does not depend on stellar mass, at least for stars more
massive than ∼0.1 Me. This result, together with theevidence that M dwarfs and giant stars span a similar
range of [O/Fe], provides a serious challenge to thescenarios for the formation of MPs that are based on
accretion.

We can conclude by emphasizing that the initial observationsof GCs carried out with the JWST have showcased thetelescope’s capability to separate MPs among M dwarfs
effectively and explore the properties of very-low-mass stars.
These results mark the opening of new horizons in the
exploration of the lower-mass regime. The heightened
sensitivity of JWST to the infrared domain will be indeed
instrumental in enabling an efﬁcient and systematic explorationof cool, low-mass stars across a substantial sample of GCs. A
current limitation is represented by the insufﬁcient availabilityof theoretical isochrones and mass–luminosity relations down
to the faint limits now reached with JWST.
Acknowledgments
We thank the anonymous referee for suggestions which
improved the manuscript. This work has been funded by theEuropean Union – NextGenerationEU RRF M4C2 1.1(PRIN 2022 2022MMEB9W: “Understanding the formation of
globular clusters with their multiple stellar generations”, CUP
C53D23001200006). T.Z. acknowledges support from theEuropean Union's Horizon 2020 research and innovation program
under the Marie Skodowska-Curie grant Agreement No.
101034319 and from the European Union NextGenerationEU.
Y.C. acknowledges support from grant RYC2021-032718-I,
ﬁnanced by MCIN/AEI/10.13039/501100011033 and theEuropean Union NextGenerationEU/PRTR. This work has been
partially supported by the Spanish MINECO grant PID2020117252GB-I00 and by the ABAUR/Generalitat de Catalunya
grant SGR-386/2021. S.J. acknowledges support from the NRF
of Korea (2022R1A2C3002992 and 2022R1A6A1A03053472).
Data Availability
The data underlying this article will be shared on reasonable
request to the corresponding author.
ORCID iDs
A. F. Marino https://orcid.org/0000-0002-1276-5487
A. P. Milone https://orcid.org/0000-0001-7506-930X
M. V. Legnardi https://orcid.org/0000-0003-3153-1499
A. Renzini https://orcid.org/0000-0002-7093-7355
E. Dondoglio https://orcid.org/0000-0001-8415-8531
Y. Cavecchi https://orcid.org/0000-0002-6447-3603
G. Cordoni https://orcid.org/0000-0002-7690-7683
A. Dotter https://orcid.org/0000-0002-4442-5700
E. P. Lagioia https://orcid.org/0000-0003-1713-0082
T. Ziliotto https://orcid.org/0000-0001-8538-2068
M. G. Carlos https://orcid.org/0000-0003-1757-6666
S. Jang https://orcid.org/0000-0002-1562-7557
A. Mohandasan https://orcid.org/0000-0001-5182-0330
M. Tailo https://orcid.org/0000-0002-1128-098X

The deep NIRCam photometry collected as part of GO2560
has allowed us to detect objects in the ﬁeld of 47 Tucanae down
to mF322W2 = 27.0. We have thus explored the faintest MS andbrown-dwarf regions, which are poorly investigated in GCs.
The corresponding results can be summarized as follows.
1. Based on both the mF322W2 versus mF115W − mF322W2 andthe mF322W2 versus mF150W2 − mF322W2 CMDs, we
detected a main discontinuity along the sequence ofvery-low-mass stars (masses smaller than ∼0.1 Me) at
mF322W2 ∼ 24.2 mag. We also notice clear changes of the
MS slope around mF322W2 = 23.3 and 24.5 mag.
2. The F322W2 luminosity function of MS stars exhibits adrop in the number of stars around mF322W2 = 25.3 mag,and the number of stars per magnitude interval rises up at
fainter luminosities. We tentatively associate this gap
with the hydrogen-burning limit. Hence, the deep CMD
obtained from GO2560 photometry unveils, for the ﬁrst
time in a GC, the brown-dwarf cooling sequence.

References
Anderson, J. 2022, AAS Meeting, 54, 206.01
Anderson, J., & King, I. R. 2000, PASP, 112, 1360
Anderson, J., & King, I. R. 2006, PSFs, Photometry, and Astronomy for the
ACS/WFC ACS 2006-01, Instrument Science ReportAnderson, J., Piotto, G., King, I. R., Bedin, L. R., & Guhathakurta, P. 2009,
ApJL, 697, L58
Anderson, J., Sarajedini, A., Bedin, L. R., et al. 2008, AJ, 135, 2055
Ballesteros-Paredes, J., Hartmann, L. W., Pérez-Goytia, N., & Kuznetsova, A.
2015, MNRAS, 452, 566

14

The Astrophysical Journal, 965:189 (15pp), 2024 April 20

Marino et al.

Bastian, N., Lamers, H. J. G. L. M., de Mink, S. E., et al. 2013, MNRAS,
436, 2398
Bastian, N., & Lardo, C. 2018, ARA&A, 56, 83
Bell, R. A., Hesser, J. E., & Cannon, R. D. 1983, ApJ, 269, 580
Bellini, A., Anderson, J., & Bedin, L. R. 2011, PASP, 123, 622
Bellini, A., Anderson, J., Bedin, L. R., et al. 2017, ApJ, 842, 6
Bianchini, P., Varri, A. L., Bertin, G., & Zocchi, A. 2013, ApJ, 772, 67
Bondi, H., & Hoyle, F. 1944, MNRAS, 104, 273
Briley, M. M., Hesser, J. E., & Bell, R. A. 1991, ApJ, 373, 482
Brown, J. A., Wallerstein, G., & Oke, J. B. 1990, AJ, 100, 1561
Bushouse, H., Eisenhamer, J., Dencheva, N., et al. 2023, JWST Calibration
Pipeline v1.10.0, Zenodo, 10.5281/zenodo.7795697
Cadelano, M., Pallanca, C., Dalessandro, E., et al. 2023, A&A, 679, L13
Calura, F., D’Ercole, A., Vesperini, E., Vanzella, E., & Sollima, A. 2019,
MNRAS, 489, 3269
Carretta, E., Bragaglia, A., Gratton, R., & Lucatello, S. 2009, A&A, 505, 139
Castelli, F. 2005, MSAIS, 8, 25
Cordero, M. J., Pilachowski, C. A., Johnson, C. I., et al. 2014, ApJ, 780, 94
Cordoni, G., Milone, A. P., Mastrobuono-Battisti, A., et al. 2020, ApJ, 889, 18
Cottrell, P. L., & Da Costa, G. S. 1981, ApJL, 245, L79
Dantona, F., Gratton, R., & Chiefﬁ, A. 1983, MmSAI, 54, 173
D’Antona, F., Vesperini, E., D’Ercole, A., et al. 2016, MNRAS, 458, 2122
de Mink, S. E., Pols, O. R., Langer, N., & Izzard, R. G. 2009, A&A, 507, L1
Decressin, T., Meynet, G., Charbonnel, C., Prantzos, N., & Ekström, S. 2007,
A&A, 464, 1029
Denissenkov, P. A., & Hartwick, F. D. A. 2014, MNRAS, 437, L21
Dickens, R. J., Bell, R. A., & Gustafsson, B. 1979, ApJ, 232, 428
Dobrovolskas, V., Kučinskas, A., Bonifacio, P., et al. 2014, A&A, 565, A121
Dondoglio, E., Milone, A. P., Lagioia, E. P., et al. 2021, ApJ, 906, 76
Dondoglio, E., Milone, A. P., Renzini, A., et al. 2022, ApJ, 927, 207
Dotter, A., Chaboyer, B., Jevremović, D., et al. 2008, ApJS, 178, 89
Dotter, A., Ferguson, J. W., Conroy, C., et al. 2015, MNRAS, 446, 1641
Gaia Collaboration, Brown, A. G. A., Vallenari, A., et al. 2021, A&A, 649, A1
Gerasimov, R., Burgasser, A. J., Caiazzo, I., et al. 2024, ApJ, 961, 139
Gieles, M., Charbonnel, C., Krause, M. G. H., et al. 2018, MNRAS, 478, 2461
Girardi, L., Groenewegen, M. A. T., Hatziminaoglou, E., & da Costa, L. 2005,
A&A, 436, 895
Gratton, R., Bragaglia, A., Carretta, E., et al. 2019, A&ARv, 27, 8Harris, W. E. 1996, AJ, 112, 1487
Jang, S., Milone, A. P., Legnardi, M. V., et al. 2022, MNRAS, 517, 5687
Kraft, R. P. 1994, PASP, 106, 553
Krause, M., Charbonnel, C., Decressin, T., Meynet, G., & Prantzos, N. 2013,
A&A, 552, A121
Kurucz, R. L. 1970, SAO Special Report 309
Kurucz, R. L. 2005, MSAIS, 8, 14

Kurucz, R. L., & Avrett, E. H. 1981, Solar Spectrum Synthesis I. A sample
atlas from 224 to 300 NM (Cambridge: Smithsonian AstrophysicalObservatory)
Lacchin, E., Mastrobuono-Battisti, A., Calura, F., et al. 2024, A&A, 681, A45
Lagioia, E. P., Milone, A. P., Marino, A. F., et al. 2018, MNRAS, 475, 4088
Lagioia, E. P., Milone, A. P., Marino, A. F., et al. 2021, ApJ, 910, 6
Lee, J.-W. 2022, ApJS, 263, 20
Legnardi, M. V., Milone, A. P., Armillotta, L., et al. 2022, MNRAS, 513, 735
Legnardi, M. V., Milone, A. P., Cordoni, G., et al. 2023, MNRAS, 522, 367
Marino, A. F., Milone, A. P., Casagrande, L., et al. 2016, MNRAS, 459, 610
Marino, A. F., Milone, A. P., Dondoglio, E., et al. 2023, ApJ, 958, 31
Marino, A. F., Milone, A. P., Renzini, A., et al. 2019, MNRAS, 487, 3815
Milone, A. P., Cordoni, G., Marino, A. F., et al. 2023b, A&A, 672, A161
Milone, A. P., & Marino, A. F. 2022, Univ, 8, 359
Milone, A. P., Marino, A. F., Bedin, L. R., et al. 2014, MNRAS, 439, 1588
Milone, A. P., Marino, A. F., Bedin, L. R., et al. 2019, MNRAS, 484, 4046
Milone, A. P., Marino, A. F., Cassisi, S., et al. 2012b, ApJL, 754, L34
Milone, A. P., Marino, A. F., Cassisi, S., et al. 2017a, MNRAS, 469, 800
Milone, A. P., Marino, A. F., Da Costa, G. S., et al. 2020, MNRAS, 491, 515
Milone, A. P., Marino, A. F., Dotter, A., et al. 2023a, MNRAS, 522, 2429
Milone, A. P., Marino, A. F., Mastrobuono-Battisti, A., & Lagioia, E. P. 2018a,
MNRAS, 479, 5005
Milone, A. P., Marino, A. F., Renzini, A., et al. 2018b, MNRAS, 481, 5098
Milone, A. P., Piotto, G., Bedin, L. R., et al. 2012a, A&A, 537, A77
Milone, A. P., Piotto, G., Renzini, A., et al. 2017b, MNRAS, 464, 3636
Nardiello, D., Griggio, M., & Bedin, L. R. 2023, MNRAS, 521, L39
Nardiello, D., Milone, A. P., Piotto, G., et al. 2018, MNRAS, 477, 2004
Partridge, H., & Schwenke, D. W. 1997, JChPh, 106, 4618
Phillips, M. W., Tremblin, P., Baraffe, I., et al. 2020, A&A, 637, A38
Renzini, A., D’Antona, F., Cassisi, S., et al. 2015, MNRAS, 454, 4197
Renzini, A., Marino, A. F., & Milone, A. P. 2022, MNRAS, 513, 2111
Richer, H. B., Heyl, J., Anderson, J., et al. 2013, ApJL, 771, L15
Sabbi, E., Lennon, D. J., Anderson, J., et al. 2016, ApJS, 222, 11
Sbordone, L., Bonifacio, P., Castelli, F., & Kurucz, R. L. 2004, MSAIS, 5, 93
Sbordone, L., Bonifacio, P., & Castelli, F. 2007, in IAU Symp. 239,
Convection in Astrophysics, ed. F. Kupka, I. Roxburgh, & K. L. Chan
(Cambridge: Cambridge Univ. Press), 71
Tailo, M., Milone, A. P., Lagioia, E. P., et al. 2020, MNRAS, 498, 5745
VandenBerg, D. A., Edvardsson, B., Casagrande, L., & Ferguson, J. W. 2022,
MNRAS, 509, 4189
Ventura, P., D’Antona, F., Mazzitelli, I., & Gratton, R. 2001, ApJL, 550, L65
Wang, L., Kroupa, P., Takahashi, K., & Jerabkova, T. 2020, MNRAS,
491, 440
Zennaro, M., Milone, A. P., Marino, A. F., et al. 2019, MNRAS, 487, 3239
Ziliotto, T., Milone, A., Marino, A. F., et al. 2023, ApJ, 953, 62

15