11institutetext:  INAF-Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5,
50125 Firenze, Italy
11email<EMAIL_ADDRESS>
Department of Physics “E. Fermi”, University of Pisa, Largo Bruno Pontecorvo
3, 56127 Pisa, Italy
INFN, Sezione di Pisa, Largo Bruno Pontecorvo 3, 56127 Pisa, Italy
SUPA, School of Science and Engineering, University of Dundee, Nethergate, DD1
4HN, Dundee, UK
Department of Physics, University of Warwick, Gibbet Hill Road, Coventry CV4
7AL, UK
Centre for Exoplanets and Habitability, University of Warwick, Gibbet Hill
Road, Coventry CV4 7AL, UK
European Southern Observatory, Karl-Schwarzschild-Strasse 2, 85748 Garching
bei München, Germany
Purple Mountain Observatory, Chinese Academy of Sciences, 10 Yuanhua Road,
Nanjing 210023, China
School of Science, University of New South Wales, Canberra, ACT 2600,
Australia
Institute of Theoretical Physics and Astrophysics, University of Kiel,
Leibnizstraße 15, 24118 Kiel, Germany
Max Planck Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany

# Protoplanetary and debris disks in the $\eta$ Chamaeleontis Association:

A sub-millimeter survey obtained with APEX/LABOCA††thanks: This publication is
based on data acquired with the Atacama Pathfinder Experiment (APEX). APEX is
a collaboration between the Max-Planck-Institut für Radioastronomie, the
European Southern Observatory, and the Onsala Space Observatory. Based on
086.C-0174A-2010, PI: A. Sicilia-Aguilar; 082.F-9304-2008, PI: R.Liseau.
observations

V. Roccatagliata 112233 A. Sicilia-Aguilar 44 M. Kim 5566 J. Campbell-White
7744 M. Fang 8899 S. J. Murphy 1010 S. Wolf 1111 W. A. Lawson 1010 Th. Henning
1212 J. Bouwman 12121122334455667788101011111212

(Received April 14, 2023; Accepted XX, 2023)

###### Abstract

Context. Nearby associations are ideal regions to study coeval samples of
protoplanetary and debris disks down to late M-type stars. Those aged
5-10,Myrs, where most of the disk should have already dissipated forming
planets, are of particular interest.

Aims. We present the first complete study of both protoplanetary and debris
disks in a young region, using the $\eta$ Chamaeleontis ($\eta$ Cha)
association as a test bench to study the cold disk content. We obtained sub-
millimeter data for the entire core population down to late M-type stars, plus
a few halo members.

Methods. We performed a continuum sub-millimeter survey with APEX/LABOCA of
all the core populations of $\eta$ Cha association. These data are combined
with archival, multi-wavelength photometry to compile a complete spectral
energy distribution. Disk properties have been derived by modeling
protoplanetary and debris disks using RADMC 2D and DMS, respectively. We
compute a lower limit of the disk millimeter fraction, which is then compared
to the corresponding disk fraction in the infrared for $\eta$ Cha. We also
revisit and refine the age estimate for the region, using the Gaia eDR3
astrometry and photometry for the core sources.

Results. We find that protoplanetary disks in $\eta$ Cha typically have holes
with radii of the order of 0.01 to 0.03 AU, while ring-like emission from the
debris disks is located between 20 au and 650 au from the central star. The
parallaxes and Gaia eDR3 photometry, in combination with the PARSEC and
COLIBRI isochrones, enable us to confirm an age of $\eta$ Cha between 7 and 9
Myrs. In general, the disk mass seems insufficient to support accretion over a
long time, even for the lowest mass accretors, a clear difference compared
with other regions and also a sign that the mass budget is further
underestimated. We do not find a correlation between the stellar masses,
accretion rates, and disk masses, although this could be due to sample issues
(the objects are few and mostly low-mass). We confirm that the presence of
inner holes is not enough to stop accretion unless accompanied by dramatic
changes to the total disk mass content. Comparing $\eta$ Cha with other
regions at different ages, we find that the physical processes responsible for
debris disks (e.g., dust growth, dust trapping) efficiently act in less than 5
Myrs.

###### Key Words.:

individual objects: eta Cha – stars: pre-main-sequence – stars: circumstellar
matter

##  1 Introduction

Protoplanetary disks are the site of planet formation and the source of the
building blocks of forming planets (e.g., Drażkowska et al. 2023). The last
phase of the evolution of protoplanetary disks, the so-called debris disk
phase, is characterized by remnant, second-generation dust without gas (e.g.,
Wyatt 2005). Recently, an increasing number of debris disks have been found to
host gas (e.g., Hughes et al. 2018), which can be a leftover from primordial
or a second-generation gas (e.g., Kóspál et al. 2013; Kral et al. 2017). In
particular, in the debris disks, only CO isotopologues have been detected
(Smirnov-Pinchukov et al. 2022) most probably due to the self-shielding of CO
itself.

The dissipation of dust (e.g., Hernández et al. 2007; Sicilia-Aguilar et al.
2013a) and gas (Fedele et al. 2010) in protoplanetary disks takes place within
a few million years (e.g. Manara et al. 2023) after the collapse of the
molecular cloud. The major mechanisms responsible for the dissipation of the
disk material are mass accretion (e.g., Hartmann et al. 2005; Manara et al.
2023), disk photoevaporation by the radiation of the central star or by
external radiation (Alexander et al. 2006; Owen et al. 2010; Picogna et al.
2019), planet formation (e.g., Drażkowska et al. 2023), binarity (Bouwman et
al. 2006), and dynamical interactions (e.g., Cuello et al. 2019). The
fundamental implication of these studies is that, in about 5 Myr, almost all
the material required to form planets is gone, so planets must form within the
first million years of disk evolution (e.g., Sicilia-Aguilar et al. 2006;
Manara et al. 2018). Inside-out disk dissipation, which was long proposed as
the expected mechanism for disk dispersal (Strom et al. 1989), was also
thoroughly confirmed via observations with the Spitzer Space Telescope
(Sicilia-Aguilar et al. 2006; Hernández et al. 2007; Teixeira et al. 2012;
Sicilia-Aguilar et al. 2013a), typically using a set of diagnostic diagrams in
the infrared color-color plane (e.g., Koepferl et al. 2013). The impact of
high-mass stars in the dissipation timescale of protoplanetary disks is still
controversial (Richert et al. 2018, 2015; Guarcello et al. 2016; Gaczkowski et
al. 2013; Roccatagliata et al. 2011).

The study of disk lifetimes, however, was not systematically extended to far-
infrared and longer wavelengths, as highlighted, for instance, by Ercolano &
Pascucci (2017). This lack of completeness was caused by both the sensitivity
of the instruments and the fact that not all members are identified in nearby
associations.

At millimeter wavelengths, Williams & Cieza (2011) and Williams (2011) found a
rapid evolution of disks within a few mega-years, interpreted as an efficient
formation of millimeter-size grains rather than a decline in the disk mass.
Recent ALMA large surveys, aiming mainly at Class II sources, led to a
statistically significant analysis of the dust masses in clusters at different
ages (e.g., Williams et al. 2019), confirming a decrease of dust mass over
time, as well as a positive correlation between disk and stellar mass (e.g.,
Ansdell et al. 2016, 2018) and which is even steeper in older regions
(Pascucci et al. 2016). ALMA observations of Class III sources in Ophiuchus
(Cieza et al. 2019) and Lupus (Lovell et al. 2020) suggested a rapid dispersal
of millimeter-sized dust from protoplanetary disks and planetesimal belt
formation already by 2 Myr. Protoplanetary disks resolved by ALMA, e.g., in
the Disk Substructures at High Angular Resolution Project (DSHARP) disks
(Stammler et al. 2019), triggered a new discussion on different evolutionary
paths followed by resolved structured and not-structured Herbig (e.g. Garufi
et al. 2017) vs Tauri disks (e.g. van der Marel et al. 2018). Debris disks
have been suggested to derive from the efficient dust traps in the
planetesimal belts of structured disks (e.g., Cieza et al. 2021; Jiang & Ormel
2021; Michel et al. 2021) while, for other disks, the decrease in dust mass
would be mainly caused by radial drift.

The evolution of protoplanetary and debris disks can be also strongly
influenced by flybys (e.g. Cuello et al. 2019, 2020). In particular, Bertini
et al. (2023) found a high incidence of close encounters for debris disks.
This might have fundamental implications in the evolution of both the
perturber source and the perturbed debris system since an exchange of material
can take place during a close encounter (e.g., as predicted by Picogna &
Marzari 2014) and change in the inclination in the orbit of inner planets.

Nearby associations are ideal regions to draw complete samples of pre-main
sequence stars and faint protoplanetary disks around stars with spectral types
extending to late M-type stars. At about 100 pc and with an age between 5 and
10 Myr (e.g., Luhman & Steeghs 2004), the $\eta$ Chamaeleontis ($\eta$ Cha)
association is unique in terms of coeval stars with spectral types between A
and M6, different disk morphologies, accretion properties, and evolutionary
stages. Its members have been studied via near-infrared ground-based data,
Spitzer mid-infrared surveys (Megeath et al. 2005; Sicilia-Aguilar et al.
2009) and Herschel far-IR data (Riviere-Marichalar et al. 2015). Initially
discovered via X-ray (Mamajek et al. 1999), a deficit of low-mass stars in the
core population suggested either a very compact formation scenario (Moraux et
al. 2007) or the presence of a halo or dynamically ejected cluster population.
This halo population has been identified via accretion and stellar youth
tracers as well as kinematic properties (Lawson et al. 2002; Luhman & Steeghs
2004; Lyo et al. 2004; Song et al. 2004; Murphy et al. 2010). The total number
of $\eta$ Cha cluster members has been recently revised to be at least 23
using the eROSITA all-sky survey and Gaia eDR3 (Robrade et al. 2021).

In this paper, we present a new sub-millimeter survey obtained with APEX to
trace the emission of the outer part of disks and thus complement previous
studies at shorter wavelengths. We also perform the first detailed analysis of
disk properties for both protoplanetary and debris disks. In Sect. 2 we
summarize the Atacama Pathfinder Experiment (APEX)/Large APEX BOlometer CAmera
(LABOCA) observations, data reduction, and the results, as well as the revised
age of the association using the Gaia EDR3 photometric data and parallaxes. In
Sect. 3 we analyze the spectral energy distributions (SED) of the
protoplanetary and debris disks. Sect. 4 discusses the observed dust masses in
the context of the mass accretion rates available in the literature, as well
as the fraction of different classes of disks compiled from star-forming
regions at different ages. A summary of our results and conclusions is given
in Sect. 5.

##  2 Observations and data reduction

The $\eta$ Cha association was observed with APEX/LABOCA at 870 $\mu$m and
with an angular resolution of 18.6 arcsec (Siringo et al. 2009). The
observations were carried out in On-Off mode (PHOT) mainly between the end of
October and December 2010 and completed in November 2013. The data consists of
a series of scans (and a variable number of sub-scans) and Tables 4 and 5
summarize the corresponding total integration time. The opacity was determined
with sky dips performed during the night, varying between $\tau$ = 0.09 and
0.41. The flux calibration was estimated by observing several calibrators per
night, including Carina, CW-LEO, N2071IR, B13134, and PKS1057-79.

The on-off observations were calibrated and reduced using the LABOCA pipeline
within the BoA package111www.apex-
telescope.org/bolometer/laboca/reduction/BoA-woo.html, developed for APEX
bolometer data. BoA allows us to derive the opacity using the sky dips and to
reduce the calibrators during each night. The reduction of the On-Off data is
obtained using the command doOO (included in the version of the BoA package of
July 2010). This command reads as input a list of scans of the science target
and the appropriate opacity of each scan. We used the weak option which allows
the optimization of the reduction for faint sources. The reduction proceeds to
compute the flux at each nodding phase. The final flux is the cumulative flux
weighted with the flux error of each nodding. We also used the clip option to
specify a sigma clipping as a threshold to remove the single noddings that
fall outside it. The clip value adopted during the data reduction is 3. The
final fluxes are reported in Table 1. The associated error, $\sigma$, is the
statistical uncertainty measured on the data, computed as the standard
deviation of all measurements, divided by the square root of the number of
measurements. For non-detected sources, we report the 3 $\sigma$ upper limit.

To this dataset, we add archival222Program number: 082.F-9304A-2008, PI:
R.Liseau LABOCA maps of RECX 16, RECX 17 and RECX 18333Additional observations
available in the archive with corrupted scans were not recoverable.. Details
of these observations are listed in Table 6. BoA was also used to reduce the
maps, in particular, the program optimizes for faint point sources (i.e.,
reduce-map-weaksource.boa). After the first iteration the signal to noise
(S/N) over the reduced map was computed. To optimize the S/N, we performed
three iterations of the data reduction. None of the three sources were
detected. In Table 1, we report the 3 $\sigma$ upper limit, where $\sigma$,
for this observation mode, represents the standard deviation over the field
covered by the map, without considering the edge and the bad pixels remaining
in the map after the iterations. Table 1 lists the LABOCA fluxes and the
astrometric measurements of parallax and proper motions compiled from GAIA
eDR3 (Gaia Collaboration et al. 2021).

---
Figure 1: APEX flux versus spectral type for the $\eta$ Cha members. For
comparison, the objects in the list from Manara et al. (2023) are shown in
grey, scaled to the distance of $\eta$ Cha. Note the lack of significant
correlations between these quantities. Filled dots represent flux detections,
while inverted triangles show upper limits. A further color ring is added to
specify the type of disk according to the classical SED classification adopted
by Sicilia-Aguilar et al. (2009): transitional disks (TD), Class II, and Class
III (including the debris disks). RS Cha and $\eta$ Cha are considered as
upper limits due to cloud contamination.  Table 1: Summary of source
properties and LABOCA measurements. Source names adopted in the paper and the
corresponding SIMBAD database IDs are given in the first two columns, the
spectral types in the third column, while parallaxes and proper motions from
Gaia eDR3 are in columns four to six. The last column contains the LABOCA-
measured fluxes and upper limits. The error in the fluxes is the root-mean-
square (rms) of the observation. The upper limits are computed as $3\,\sigma$.

Source | SIMBAD | SpT | $\pi$ | $\mu_{\alpha}$ | $\mu_{\delta}$ |  F870μm
---|---|---|---|---|---|---
| name |  | $[mas]$ | $[mas/yr]$ | $[mas/yr]$ | [mJy]
Detections |  |  |  |  |  |
J0844 | V* EO Cha | M0 |  10.1867 $\pm$ 0.0106 |  -30.287 $\pm$ 0.015 |  26.927 $\pm$ 0.013 |  4.96 $\pm$ 0.83
J0843 | V* ET Cha | M2 |  9.6756 $\pm$ 0.2249 |  -27.932 $\pm$ 0.28 |  29.149 $\pm$ 0.27 |  5.42 $\pm$ 0.83
RECX-5 | V* EK Cha | M5 |  10.1522 $\pm$ 0.0138 |  -30.242 $\pm$ 0.018 |  26.862 $\pm$ 0.019 |  20.01 $\pm$ 2.45
RECX-6 | V* EL Cha | M2 |  10.1933 $\pm$ 0.0156 |  -29.106 $\pm$ 0.02 |  27.014 $\pm$ 0.021 |  8.38 $\pm$ 0.54
RECX-8 1 | RS Cha | A7V |  10.137 $\pm$ 0.0213 |  -27.263 $\pm$ 0.032 |  28.179 $\pm$ 0.029 |  18.62 $\pm$ 2.63
RECX-9 | V* EN Cha | M4 | – | – | – |  5.71 $\pm$ 1.29
RECX-11 | V* EP Cha | K5 |  10.1198 $\pm$ 0.0111 |  -30.145 $\pm$ 0.015 |  26.801 $\pm$ 0.011 |  12.74 $\pm$ 1.68
J0801 2 | 2MASS J08014860-8058052 | M4.4 |  8.9432 $\pm$ 0.0233 |  -19.006 $\pm$ 0.031 |  28.879 $\pm$ 0.03 |  7.58 $\pm$ 0.86
J0820 2 | 2MASS J08202975-8003259 | M4.3 |  10.2685 $\pm$ 0.0229 |  -27.872 $\pm$ 0.028 |  29.282 $\pm$ 0.032 |  8.12 $\pm$ 1.1
RX_J1005 2 | RX J1005.3-7749 | M1e |  10.0257 $\pm$ 0.0129 |  -37.326 $\pm$ 0.016 |  14.656 $\pm$ 0.015 |  7.75 $\pm$ 0.56
Marginal | Detections |  |  |  |  |
$\eta$ Cha | * eta Cha | B8V |  10.1679 $\pm$ 0.0665 |  -29.43 $\pm$ 0.18 |  26.831 $\pm$ 0.119 |  3.37 $\pm$ 1.86
J0841 | V* ES Cha | M4 |  10.1205 $\pm$ 0.0247 |  -29.957 $\pm$ 0.034 |  28.036 $\pm$ 0.036 |  2.88 $\pm$ 0.97
Upper | limits |  |  |  |  |
RECX-3 | V* EH Cha | M3 |  10.1168 $\pm$ 0.0107 |  -28.442 $\pm$ 0.014 |  27.001 $\pm$ 0.016 | ¡ 1.65
RECX-4 | V* EI Cha | K7 |  10.1366 $\pm$ 0.0121 |  -30.826 $\pm$ 0.016 |  26.006 $\pm$ 0.018 | ¡ 3.03
RECX-7 | V* EM Cha | K3 |  10.1278 $\pm$ 0.0118 |  -30.184 $\pm$ 0.014 |  27.576 $\pm$ 0.014 | ¡ 4.02
HD 75505 | HD 75505 | A1V |  10.2039 $\pm$ 0.0213 |  -30.646 $\pm$ 0.028 |  27.134 $\pm$ 0.031 | ¡ 7.86
RECX-12 | V* EQ Cha | M3 |  10.2837 $\pm$ 0.0354 |  -30.397 $\pm$ 0.045 |  26.256 $\pm$ 0.041 | ¡ 2.01
RECX-16 | 2MASS J08440915-7833457 | M5.5 |  10.0441 $\pm$ 0.0393 |  -29.844 $\pm$ 0.054 |  26.262 $\pm$ 0.049 | ¡ 40.4
RECX-17 | ECHA J0838.9-7916 | M5 |  10.0827 $\pm$ 0.0363 |  -28.868 $\pm$ 0.044 |  27.035 $\pm$ 0.041 | ¡ 9.10
RECX-18 | ECHA J0836.2-7908 | M5.5 |  9.9086 $\pm$ 0.05 |  -29.092 $\pm$ 0.059 |  27.871 $\pm$ 0.053 | ¡ 29.8
Herschel |  |  |  |  |  |
RECX1 | V* EG Cha | K4Ve |  10.1099 $\pm$ 0.0917 |  -29.62 $\pm$ 0.118 |  26.852 $\pm$ 0.103 |

1: The detection of RECX-8 is contaminated by a local extended nebulosity
highlighted by Planck between 30 an 217 GHz. For this reason, it is more
appropriate to consider the emission from RECX-8 as an upper limit.
2: Halo members from the work of Murphy et al. (2010).

##  3 Analysis

In this section, we provide the details of the data analysis. After computing
the dust disk mass from the sub-millimeter flux, the spectral energy
distributions (SEDs) of the cluster members were compiled, including our new
submillimeter photometric point from our APEX/LABOCA survey as well as
archival optical magnitudes, 2MASS JHKs data (Sicilia-Aguilar et al. 2009),
near-infrared (WISE and Akari), mid-infrared observations from Spitzer
(Sicilia-Aguilar et al. 2009) and far-infrared data from Herschel (Riviere-
Marichalar et al. 2015).

So far, only the dust composition in the inner part of the disk, observed by
Spitzer/IRS, has been analyzed for the $\eta$ Cha targets (Bouwman et al.
2006; Sicilia-Aguilar et al. 2009), together with a first interpretation of
the SEDs including up to mid-infrared wavelengths. The SEDs including the
Herschel/PACS photometry (Riviere-Marichalar et al. 2015) were not modeled,
which we now do including the submillimeter data. Our LABOCA survey did not
cover RECX 1, but we add its Herschel data to our analysis.

We then analyze the structure of protoplanetary disks and debris disks in the
$\eta$ Cha association fitting the complete SEDs. Protoplanetary disks are
fitted using the Monte Carlo radiative transfer code RADMC 2D (Dullemond &
Dominik 2004; Dullemond 2011), while for debris disks we use the DMS (Debris
disks around Main sequence Stars; Kim et al. 2018) software, optimized for
optically thin emission. For 2MASS J08014860, since its nature is not clear,
we proceed with the modeling using both codes, RADMC 2D and DMS.

###  3.1 Dust disk masses

In Fig. 1 we show the fluxes as a function of the spectral type of the star,
noting that we detect sources down to late M-type stars. For comparison, we
also show the sources with known spectral types in the list from Manara et al.
(2023), scaled to the distance of $\eta$ Cha. We highlight that there is no
correlation between the fluxes detected and the spectral type of the host
star. A Spearman rank test returns a false-alarm probability of 12% if we
consider only the detections, or 65% if upper limits are included, which is
fully consistent with the two quantities being uncorrelated. We highlight that
our work is homogeneously tracing different spectral types in various regimes
only barely covered in other works.

Under the assumption that the emission at millimeter wavelengths is optically
thin, the dust disk mass ($M_{\rm dust}$) is directly proportional to the
millimeter flux ($S_{\nu}$), so

$M_{\rm dust}=\frac{S_{\nu}D^{2}}{k_{\nu}B_{\nu}(T_{\rm dust})},$ (1)

where $k_{\nu}=k_{0}(\nu/\nu_{0})^{\,\beta}$ is the mass absorption
coefficient, $\beta$ parametrizes the frequency dependence of $k_{\nu}$,
$S_{\nu}$ is the observed flux, $D$ is the distance to the source, ${\it
T}_{\rm dust}$ is the dust temperature, and $B_{\nu}(T_{\rm dust})$ is the
Planck function. For consistency, we adopt the dust properties from Ansdell et
al. (2016), using typical assumptions of a single dust grain opacity $k_{0}$ =
3.37 cm2/g at 890 $\mu$m and a single dust temperature $T_{\rm dust}$ = 20 K,
which corresponds to the median temperature of the Taurus disks (Andrews &
Williams 2005). This approach allows us to directly compare our results with
those of Manara et al. (2016). In this calculation, we have always highlighted
that any trend in the dust mass can reflect the variation in opacity or the
differences in temperature among the variety of disks in the cluster.
Moreover, as already highlighted by Hartmann et al. (2006), disk masses are
systematically underestimated by a factor that could be up to hundreds (e.g.,
Liu et al. 2022; Rilinger et al. 2023), due to the uncertainties in the
factual dust properties, disk size, and inclination.

The fact that most disks from Manara et al. (2023) are substantially more
massive/more luminous than the $\eta$ Cha ones is a signature of a bias and
suggests that the understanding of protoplanetary disks and their dissipation
timescales and mechanisms requires including the faint, harder-to-detect, more
evolved targets as well. This is not the only bias present, since until a few
years ago, debris disks were more commonly detected around early-type stars
(e.g., Matthews et al. 2014; Sibthorpe et al. 2018) basically reflecting a
skew in the observed sample. As van der Marel & Mulders (2021) started to
highlight, the observed debris and structured disks show instead similar
detection rates as a function of their host stellar mass.

###  3.2 Modelling the protoplanetary disks with RADMC-2D

To study the structure of the protoplanetary disks in the $\eta$ Cha
association, we fitted the full SED using the Monte Carlo radiative transfer
code RADMC 2D (Dullemond & Dominik 2004; Dullemond 2011) for different
underlying disk structures to explore the various possibilities for each
system. RADMC 2D assumes a fully irradiated disk where the temperatures of all
the dust species for each given location are the same, which imposes the main
limitation on the models. The models can be set to have well-mixed gas and
dust (iterating on the solution and using the local temperature to establish
the flaring), or decoupled gas and dust (setting a given power law for the
vertical scale height $H$ vs. disk radius $R$) to simulate a more settled
disk, even though all species would be settled in the same way, which is a
simplification. The stellar photospheres rely on the same MARCS models
(Gustafsson et al. 2008) used in Sicilia-Aguilar et al. (2009), with a
temperature and radius that reproduce the stellar SED for each object.
Interstellar extinction is considered to be negligible, which is well
supported by an independent analysis carried out by Rugel et al. (2018) of
X-Shooter spectra of the $\eta$ Cha core members. The stellar properties are
summarized in Table 1.

We follow the same strategy as in Sicilia-Aguilar et al. (2015) where, rather
than trying to derive the best model, we try to understand whether each disk
can be reproduced by a standard, continuous model, or if it requires radial
(e.g., variations in composition, inner holes or gaps) or vertical changes
(e.g., variations in the disk scale height in the inner and outer disk,
settling and deviations from a well-mixed dust and gas disk). We explore
different disk hole sizes and vertical scale heights to reproduce the observed
disk emission. Taking into account that the parameter space is non-continuous
and highly degenerated even for relatively complete SEDs with Herschel and
millimeter data (Sicilia-Aguilar et al. 2015, 2016), and the variety of
processes that can affect disks, the best fit is rarely unique (and often not
perfect), but even in this case it is possible to assess whether a disk
presents signs of evolution. Since most of the SEDs show clear evidence of
holes/gaps and a radial variation in the dust properties between the inner and
the outer disk, we concentrate on exploring the different possibilities for
inner vs. outer disk structure to reproduce the observed SEDs. This includes
changes in composition, flaring, vertical scale height, and imposing a given
H/R law that simulates dust setting vs. flared gas in hydrostatic equilibrium.
Because of the complexity of the systems and the lack of resolved
observations, it is often impossible to clearly distinguish between scenarios,
but deviations with respect to standard and less evolved disks are evident.

|
---|---
|
|
|
Figure 2: SED models for the protoplanetary disks. For all cases, the
photometry data is marked by blue circles (empty symbols for marginal
detections). Spitzer/IRS spectra are plotted in black. The MARCS model stellar
photospheres are marked by dotted grey lines, and the various disk models are
represented by colored lines and labeled according to their main
characteristics (see Table 2). RECX-5: A large-scale gap is not needed to
reproduce the long wavelengths, which also means that SED alone cannot
constrain among many diverse but equally plausible SED structures. RECX-9: A
change in vertical scale height at 10-15 au is needed, which could be caused
by a gap, wall, warp, or any other structure affecting the density and the
scale height probably created by the existing companion at 20 au. RECX-11:
Best fit with relatively massive and flattened disks. A more puffed innermost
disk (either a puffed-up rim or a more extended $\sim$0.6 au region) is
required, with the disk becoming increasingly flattened and settled at larger
radii. J0801 and J0820 appear to be examples of relatively massive transition
disks with large, strongly mass-depleted inner holes. ECHA J0841: Very
flattened SED. ECHA J0843: Small gap or hole required. For ECHA J0844: Gap
and/or change in the vertical scale height needed to explain the far-IR flux.

For the dust model, we used a collisional distribution of amorphous magnesium
and iron silicate grains with sizes between 0.1 and 1000 $\mu$m and a 25%
content of amorphous carbon. The dust opacities were taken from the Jena
database (Jäger et al. 2003). In the innermost disk (inside the hole, which
does not necessarily need to be completely clean but may have lower
densities), we can keep the same dust distribution as in the outer disk, or
consider an optically thin population composed of 0.1 $\mu$m grains with a
mixture of amorphous silicates and crystalline forsterite in a 1:3 proportion.
Small and crystalline grains in the inner disk are included depending on
whether the object has a strong silicate feature. Large grains in the inner
disk tend to make the silicate feature fade, although technically it is very
hard to distinguish a radial vs. vertical (e.g., due to differential settling)
stratification of large and small grains. For disks with weak silicate
emission, keeping the same composition for the inner and outer disk can
typically reproduce the observed SED, so we chose to use a uniform grain
distribution (with variable density and vertical scale height, if needed) for
simplicity. The interface between the inner and outer disk composition is set
using the Schmidt number (Sc), for which larger values imply a more abrupt
composition change. In general, we take Sc=2/3 to ensure a smooth transition,
but this may be increased to simulate more abrupt gaps. The gas-to-dust ratio
is set to 100.

By default, the outer disk (or the whole disk, if no distinct radial
substructures are included) is considered to be flared with $H/R$ $\propto$
$R^{\rm\,1/7}$, where $H$ is the vertical scale height of the disk and $R$ is
the disk radius. The value of H/R in the outer part of the outer
($H/R_{\rm\,out}$) and inner ($H/R_{\rm\,in}$) disk, and the minimum vertical
scale height can be changed to reproduce the observed SED. Without resolved
observations, there is a strong degeneracy between the disk mass and the outer
disk radius, even in the presence of submillimeter data. We thus take the disk
radius to be Rout 100 au. The inclination angle is taken to be 45 degrees in
all cases, noting that the most dramatic differences occur only at extreme
inclinations. For the inner disk, when present, we set up an inner and outer
region, and a surface density that changes with a certain power of the radius,
typically taken to be -1. The inner disk flaring is set by default to $H/R$
$\propto$ $R^{\rm\,1/7}$, except in very flattened models. To avoid
overloading the models with extra parameters, we do not include extra puffed
inner rims, although if the height of the inner and outer disks do not match,
this effectively behaves as a puffed rim.

Since the parameter space is degenerated, we investigated the minimal model
needed to account for the observed SED, rather than attempting to derive the
(very uncertain) best fit. We thus start with continuous disks with a single
power law to define the flaring and density profile and depart from that model
by adding a differentiated inner disk if the SED cannot be reproduced with a
simpler model. For each SED, several models were run, aiming at reproducing
the observations with very different disk parameters in order to probe all
possible scenarios. In most cases, the result is that each given SED may be
equally (or nearly equally) well reproduced by very different models within a
non-continuous parameter space. Very often, none of the models produces a good
fit at all wavelengths, which is an indication that the disk structure is
likely more complex than what our simple models consider. Table 2 contains a
summary of the models, and below we provide an object-by-object discussion.
The resulting models are presented in Fig. 2, together with the observed SED
and the IRS/Spitzer spectra. Specific objects are discussed in the following
sections.

####  3.2.1 RECX-5

The stellar parameters considered for RECX-5 are a stellar mass $M_{*}$ = 0.3
M⊙ and a stellar radius $R_{*}$ = 0.7 R⊙. The SED of RECX-5 is characterized
by a very strong silicate feature (Bouwman et al. 2006; Sicilia-Aguilar et al.
2009). It is now found to have also a very high submillimeter flux, which
requires a total disk mass of about 1.8 $\times$ 10${}^{\,-3}$ M⊙ (or 6
$\times$ 10${}^{\,-3}$ $M_{*}$) to reproduce the LABOCA emission with the
given grain sizes, and a maximum scale height of 0.06 $\times$ $R_{\rm\,out}$.
RECX-5 has been modeled in detail by Bouwman et al. (2010), who suggested
that, in addition to a small inner hole 0.6 au in size, a second gap centered
at 24 au was also necessary to reproduce the flux. Details of this model,
which was also created using RADMC-2D, can be found in the original paper.

The results of our modeling are that, while an essentially empty hole (filled
only with small, optically thin dust grains including crystals) is required to
reproduce the SED, the SED alone does not require including additional gaps at
larger distances to reproduce the Herschel and LABOCA fluxes (see Fig. 2). The
required size of the inner hole is of the order of 0.5-1 AU, with a
temperature for the inner disk rim of about 200 K, but the fact that the disk
SED can be well-reproduced using very different structures at moderate to
large radii is a sign of the degeneracy resulting from unresolved
observations.

Given the number of disks that are observed to have asymmetries and large-
scale gaps (e.g., van der Marel et al. 2013; Garufi et al. 2014; Avenhaus et
al. 2014; ALMA Partnership et al. 2015; Isella et al. 2016; Benisty et al.
2018, among many others) and that the relation between accretion and Herschel
fluxes observed in transition disks hints toward the presence of further gaps
(hard to detect from the SED alone; Sicilia-Aguilar et al. 2015), we cannot
exclude the presence of further radial and azimuthal asymmetries in the disk
of RECX-5, but spatially resolved data will be required for final
confirmation.

####  3.2.2 RECX-9

For this disk, the mismatch (or discontinuity) between the mid-IR spectrum and
the far-IR data makes it very hard to produce a good fit without some drastic
radial changes at intermediate distances (12-25 au). Having a gap and/or a
change in flaring/settling at these distances can reproduce the observed
profile. The location of this asymmetry is roughly consistent with the
position of the known companion at $\sim$ 0.2 arcsecs ($\sim$ 20 au; Bouwman
et al. 2006). There is still substantial submillimeter emission, suggesting
that the disk is not fully truncated so there is still some circumbinary
material present, which may account for the extra disk mass inferred from the
submillimeter flux. We checked if the submillimeter emission could be
contaminated by another source. The Planck maps between 30 and 70 GHz do not
reveal any nearby source that could affect the submillimeter flux of
RECX-9444The nearest sources are quasars at $\sim$6∘ and 8∘.. The lack of a
strong silicate feature is a sign of strong settling and/or a strong
population of large grains in the inner disk, both characteristic of an
evolved protoplanetary disk. In fact, the disk can be fitted with the same
combination of small and large grains extending throughout the whole disk.
Some care has to be taken regarding the fit of the far-IR data, given that the
mismatch between the Spitzer 70 $\mu$m and the Herschel/PACS data points may
suggest either contamination by cloud emission or some episode of accretion
variability, although there is no evidence that RECX-9 suffers accretion
variations beyond a factor of few (see Table 3).

####  3.2.3 RECX-11

Besides being relatively settled and rich in large grains (which causes the
observed low-contrast silicate feature), the SED of RECX-11 does not display
any clear evidence of a distinct radial substructure, except for maybe a
slightly puffed-up inner disk and a potential change in flaring/gap at
intermediate radii (see Fig. 2). The strong submillimeter emission is
characteristic of a relatively settled (or flattened), but still massive disk.
Note that the models here assume that grains of all sizes are well-mixed,
which may not be the case, especially regarding the larger grains that
dominate the far-IR and submillimeter points. We also assume that there is no
difference in composition between the inner and the outer disk.

####  3.2.4 ECHA J0820

One of the members of the $\eta$ Cha extended population discovered by (Murphy
et al. 2010), together with ECHA J0801, ECHA J0820 also shares with the former
a very low near-IR emission. The disk has substantial emission at 22 $\mu$m as
well as at 870 $\mu$m setting strong constraints on the presence of gaps and
holes in the inner disk. In this case, a gap or a partly evacuated hole 2-3 au
in size offers a very good fit to the SED. Because of the very low near-IR
emission, the inner disk is required to be both depleted and flattened,
compared to the outer disk.

####  3.2.5 ECHA J0843

Also known as RECX-15 and ET Cha, fitting its strong silicate feature and the
high emission from 18 $\mu$m on requires a change of flaring angle,
composition, or slope in the innermost disk, which could be a sign of a gap in
the disk at small radii ($\leq$ 0.5 au). The disk is consistent with a strong
and optically thick inner part, followed by a gap or a change in vertical
scale height. The precise parameters of the gap depend strongly on the dust
properties and on the shape of the inner rims of the gap, which we are keeping
as a simple vertical slope in this case. Since the silicate feature is strong,
we use a small-grain (amorphous silicate and forsterite, 0.1 $\mu$m in size)
dominated inner disk and an outer disk with the usual larger grain
distribution, but a similar result can be attained by including the small
grains on the outermost layers of the disk or within the gap.

####  3.2.6 ECHA J0841

This disk presents an anomalously high submillimeter datapoint, compared to
the rest of the SED, but since it is a 2.9 $\sigma$ (marginal) detection, it
may be considered as a possible upper limit. The innermost disk SED, including
the Herschel data, is consistent with a strongly settled system with very low
mass, while the mass predicted from the submillimeter flux, despite its
uncertainty, is higher than what one would expect for a standard, continuous
disk with such a flat inner disk. One possibility would be the presence of
large gaps at large radii (to explore this possibility, we would need data in
the 200-700 $\mu$m range), or a nearly geometrically flat disk (which could be
caused by settling in a large-grain, gas-poor disk. Another possibility would
be that the millimeter flux is overestimated and thus an upper limit.
Examination of the IR images does not reveal any evidence of a potential
contaminant submillimeter source at distances relevant to the APEX beam. In
any case, the existing near- to far-IR data show a very evolved, settled
system.

####  3.2.7 ECHA J0844

The structure of this disk is very similar to ECHA J0843, although the
gap/wall is located at larger distances and the disk appears more settled and
evolved. We again encounter the same difficulties, that the flux in the
mid/far-IR depends strongly on the properties of the gap and the gap walls,
leading to large uncertainty in a broad and non-continuous parameter space.
None of the simple models used here offers a perfect fit. The comparison with
the well-mixed disk (where dust follows the hydrostatic equilibrium
distribution of the gas) suggests that the disk is artificially puffed up in
the innermost part or at least, that it has a puffed-up inner rim. Varying the
dust distributions and surface densities in the inner disk could be done to
improve the fit, but the uncertainty cannot be resolved without further (if
possible, spatially-resolved) information.

####  3.2.8 $\eta$ Cha

Figure 3: SED fit for $\eta$ Cha. Note that the LABOCA 870 $\mu$m point is not
fitted, as it is most likely contaminated as explained in the text.

The intermediate-mass star $\eta$ Cha is a special case, compared to the rest
of the objects, since it has a disk that seems to be intermediate between a
protoplanetary disk and a debris disk, but it does not seem to have a large
inner hole or gap and the disk may start straight away at the dust destruction
radius. The disk had been also proposed to arise from decretion, as in a
Classical Be star, but optical spectroscopy does not reveal the required gas
features (Sicilia-Aguilar et al. 2009). We have further examined in detail the
existing spectroscopy of $\eta$ Cha in the public ESO database, and none of
the observations reveal the gas lines expected in a Classical Be star, so that
the disk around $\eta$ Cha appears to be either an evolved protoplanetary disk
or a young debris disk extended over a relatively large range of radii.

When fitting the SED for $\eta$ Cha, we notice that the LABOCA marginal
detection must be either a non-detection, contaminated or the error
underestimated, because it is not possible to reproduce it with any sensible
model of disk associated with $\eta$ Cha. The emission would need to arise
from very cold dust ($<$10 K) and therefore, be located at a very large
distance from the relatively luminous $\eta$ Cha. The very low temperature
would require a very large dust mass (several Jupiter masses in dust alone) so
that it likely corresponds to molecular cloud emission (or it is rather a non-
detection than a marginal detection).

The inner disk requires optically thin dust, with a minimum size of the order
of 20 $\mu$m to avoid silicate emission. As mentioned, an inner hole is not
needed, but the low flux at 70um suggests that the disk is relatively small
($<$100 au, with the best fit being around 50 au). The flaring of the disk
cannot be constrained because it is optically thin, but the fact that it is
very hard to reproduce the shorter and longer wavelengths with a uniform disk
suggests that the disk matter may be distributed asymmetrically in the inner
vs. the outer disk, or that the disk has a slightly faster-falling power law
of the surface density with radius (the best model has an exponent of -1.2,
instead of -1 which is used in the rest of cases). The results of the fitting
procedure are shown in Fig. 3.

|
---|---
|
|
|
Figure 4: Left: SED models for the debris disks. The photometry data are
marked by red circles and triangles (the upper limits). Stellar photospheric,
disk, and total emission (stellar photospheric + disk models) are represented
by grey dotted lines, and red and blue continuous lines, respectively (see
Table 2 for details on individual models). The blue dots in the SED of RX
J1005.3-7749 show the detection from Murphy et al. (2010). Right: Results of
the SED modeling of the debris disks showing the mass of the disk as a
function of the location of the narrow dust ring.

###  3.3 Modelling the debris disks with DMS

We define Class III sources with ${f_{\rm\,dust}\,(\equiv
L_{\rm\,dust}/L_{\rm\,*})\,<\,10^{\,-2}}$ as debris disks. These disks are
modeled assuming an optically thin emission from a dust distribution around a
stellar source. As the first approximation, the debris dust in the disk is
distributed in a very narrow ring with a fixed radial width of 20 au (Fig. 8).
The re-emission of the stellar radiation from the disks is reproduced by
varying the location of the inner and outer ring radii between 10 and 700 AU,
respectively, as well as the total disk mass. This approximation of the ring
emission is well supported by recent observations at different wavelengths of
resolved debris disks distributed as a narrow ring (e.g., Kalas et al. 2013;
Milli et al. 2017; Schneider et al. 2018; Marino et al. 2018; Booth et al.
2023). Debris dust material is assumed to be astrosilicates with a bulk
density of 3.5 g cm-3 (Draine 2003). We note that most of the debris disks or
Class III (without debris material around the central star) objects are stars
with very late spectral type. For these cases, the radiation pressure exerted
by these stars is insufficient to achieve a blowout size, particularly for
stars with temperatures below 5250 K (Kirchschlager & Wolf 2013). Thus, we
consider a range of dust grain sizes from 0.1 $\mu$m to 1000 $\mu$m, where the
grain size distribution $n\,(a)$ follows the power-law distribution $n\,(a)$
$\propto$ $a^{\rm-3.5}$ (Dohnanyi 1969).

We use a newly developed software tool DMS (Debris disks around Main sequence
Stars; Kim et al. 2018) which is optimized for the simulation of debris disks,
i.e., optically thin systems. In particular, it allows us to simulate
scattered light and thermal dust reemission images, the continuum spectral
energy distributions, and scattered light polarization images. The optical
properties of the dust grains are computed using the tool miex (Wolf &
Voshchinnikov 2004). The characteristics of the central stars are listed in
Table 1.

The results of the modeling (Fig. 4) of the debris disks show that in two
cases the ring the close to the central star. The first case is represented by
2MASS J08014860, discussed in detail in Sect. 3.4, while the second case is
represented by RECX 3, an M3 star, where the infrared photometry already
suggested the presence of cold material around the star. Our modeling allows
us to define a range in distance between 20 and 160 au for a ring with a mass
of 1-9 $\times$ 10-7M⊙. In the case of RECX 1, a K7 star, where two detections
in the far-infrared lie above the stellar photosphere, the missing LABOCA
observation shows the poor fitting of the SED which leads to a poor constraint
of the location of a 0.2-1.2 $\times$ 10-7M⊙ disk between 100 and 700 au.
Finally, the M1 object, RX J1005.3-7749555Gaia DR2 5202670052321415040, was
first classified as a pre-main-sequence object by Covino et al. (1997) and
then Murphy et al. (2010) as a possible dynamical $halo$ member. It harbors
the most massive ring (1.0-1.5 $\times$ 10-5 M⊙) at a location between 590 and
680 au, more distant than expected, meaning that the dust temperature is
extremely low. This, in turn, results in a remarkably high derived dust mass.
We note that there is a possibility that the millimeter emission attributed to
the debris disk could be contaminated by background galaxies or asymptotic
giant branch stars. Similar issues have been addressed in the previous
observations, in particular, for the Class III young stellar objects, showing
similar infrared excesses that lead to similar expected signatures in
observations (e.g., Oliveira et al. 2009; Evans et al. 2009). However, with
the recent release of the second set of data from the Gaia mission Gaia
Collaboration et al. 2018 and the availability of high-resolution optical
spectra in the present study, allowing for more accurate identification and
classification, we are better equipped to differentiate and exclude these
sources from consideration in our analysis.

Tentative SED models have been also carried out for the probable debris disks,
RECX 6 and RECX8. We explore different disk models, represented by blue lines
in Fig. 9 in the Appendix, but none of them were able to reproduce all of the
observed data, as is the case with other systems. We note that RECX 8 (RS Cha)
is an eclipsing binary system with two components of similar mass and
luminosity, which is confirmed using the TESS lightcurve by (Steindl et al.
2021). Furthermore, the third low-mass component has been recently proposed by
Woollands et al. (2013). Due to the complexity of this system, it cannot be
approximated by our assumed single thin-ring model.

Table 2: Parameters used in the disk models. Note that Rgap and Tgap refer to the location of a radial discontinuity, which may involve not only a thin gap but a change in radial properties (e.g., dust grain properties, vertical scale height, density power law). Note that for some models, we set the inner disk temperature in RADMC, and the distance is calculated accordingly, although the opposite is also feasible. In this table, we list the distances as well because they are easier to compare with resolved observations and other models in the literature, and the derived quantities are marked with ∗. The parameter pH/R corresponds to the exponent in ($H/R$) $\propto$ R${}^{\,\rm{p}}$, and p${}_{\Sigma\,\rm{dust,in}}$ is the power law vs radius for the surface density in the inner disk. The notes include any further deviation from the standard model (see text) not reflected in the previous parameters. Source | Model | Teff | R∗ | Rin | Tin | Rgap | Tgap | H/Rout | H/Rin | $\Sigma_{\rm dust,in}$ | pH/R | Mdisk | Notes
---|---|---|---|---|---|---|---|---|---|---|---|---|---
| # | (K) | R⊙ | (au) | (K) | (au) | (K) | (out/min) | (out/min) | (g cm-2) | (in, out) | (M⊙) |
RECX-5 | #1 | 3350 | 0.76 | 0.018 | 1480∗ | 1∗ | 200 | 0.06/0.005 | 0.04/0.01 | 1e-5 | 1/7, 1/7 | 1.8e-3 | Schmidt Nr=10, pΣdust,in=-1.4,
|  |  |  |  |  |  |  |  |  |  |  |  | Small grains
RECX-9 | #1 | 3305 | 0.95 | 0.02∗ | 1500 | – | – | 0.07/0.005 | – | – | 1/7 | 1.4e-4 |
| #2 |  |  | 0.07 | 830∗ | 12 | 63 | 0.1/0.005 | 0.3/0.05 | 1e-5 | 1/7, 1/7 | 1.2e-4 |
| #3 |  |  | 0.07 | 830∗ | 20∗ | 50 | 0.2/0.005 | 0.2/0.05 | 1e-5 | 1/7, 1/7 | 1.4e-4 |
RECX-11 | #1 |  |  | 0.05 | 1500∗ | 0.64 | 430∗ | 0.08 | 0.12 | 2e-5 | 1/7, 1/7 | 2.5e-4 |
| #2 |  |  | 0.05 | 1500∗ | 0.64 | 430∗ | 0.05 | 0.14 | 5e-5 | 1/7, 1/7 | 3.4e-4 |
| #3 |  |  | 0.05 | 1500∗ | 0.64 | 430∗ | 0.05 | 0.14 | 5e-5 | 1/7, 1/7 | 3.4e-4 | Small grains in inner disk
| #4 |  |  | 0.05 | 1500∗ | 0.64 | 430∗ | – | – | 5e-5 | – | 3.4e-4 | Well-mixed gas and dust
J0801 | #1 | 3280 | 0.62 | 0.02 | 1250∗ | 12.5 | 50∗ | 0.03 | 0.03 | 1e-5 | 1/7, 1/7 | 1e-4 |
| #2 |  |  | 0.02 | 1250∗ | 12.5 | 50∗ | 0.03 | 0.03 | 1e-6 | 1/7, 1/7 | 8e-4 |
| #3 |  |  | 0.02 | 1250∗ | 19.5 | 40∗ | 0.02 | 0.03 | 1e-6 | 1/7, 1/7 | 8e-4 |
| #4 |  |  | 0.02 | 1250∗ | 19.5 | 40∗ | 0.01 | 0.01 | 1e-6 | 1/7, 1/7 | 8e-4 |
J0820 | #1 | 3280 | 0.60 | 0.03 | 1000∗ | 3 | 100∗ | 0.17 | 0.05 | 1e-4 | 1/7, 1/7 | 4e-4 |
| #2 |  |  | 0.03 | 1000∗ | 2.1 | 120∗ | 0.17 | 0.05 | 1e-4 | 1/7, 1/7 | 4e-4 |
| #3 |  |  | 0.03 | 1000∗ | 3 | 100∗ | 0.05 | 0.20 | 1e-4 | 1/7, 1/7 | 4e-4 |
J0841 | #1 |  |  | 0.03 | 900∗ | 0.1∗ | 495 | 0.04/0.005 | 0.1/0.05 | 1e-5 | 1/7, 1/7 | 1.2e-4 |
| #2 |  |  | 0.03 | 900∗ | 0.1∗ | 495 | 0.03/0.005 | 0.1/0.05 | 1e-5 | 1/7, 1/7 | 4e-5 |
| #3 |  |  | 0.03 | 900∗ | 0.1∗ | 495 | 0.03/0.005 | 0.1/0.05 | 1e-5 | 1/7, 0 | 8e-5 |
| #4 |  |  | 0.03 | 900∗ | 1.5∗ | 130 | 0.02/0.005 | 0.1/0.01 | 3e-5 | 0, 0 | 5e-4 |
J0843 | #1 | 3430 | 0.79 | 0.03 | 1200∗ | 0.18∗ | 495 | 0.28 | 0.17 | 5e-3 | 1/7, 1/7 | 3e-4 | Small grains in inner disk
| #2 |  |  | 0.03 | 1200∗ | 0.9∗ | 220 | 0.25 | 0.17 | 5e-3 | 1/7, 1/7 | 2.4e-4 | Small grains in inner disk
J0844 | #1 | 3145 | 0.37 | 0.03 | 750∗ | 3∗ | 78 | – | – | 1e-4 | – | 3.9e-4 | Well-mixed gas and dust
| #2 |  |  | 0.03 | 750∗ | 2∗ | 92 | 0.15 | 0.2 | 1e-4 | 1/7, 0 | 3.9e-4 |
| #3 |  |  | 0.03 | 750∗ | 3∗ | 78 | 0.17 | 0.1 | 1e-4 | 1/7, 0 | 3.4e-4 |
| #4 |  |  | 0.03 | 750∗ | 3∗ | 78 | 0.15 | 0.1 | 1e-4 | 1/7, 0 | 3.9e-4 |
$\eta$ Cha | #1 | 10500 | 2 | 1.6∗ | 900 | – | – | 0.01/0.002 | – | – | 0 | 1.5e-6 | Large grains, not flared, 100 au
| #2 |  |  | 0.6∗ | 1500 | – | – | 0.06/0.002 | – | – | 1/7 | 3e-7 | Large grains
| #3 |  |  | 0.9∗ | 1200 | – | – | 0.03/0.002 | – | – | 1/7 | 3e-7 | Large grains
| #4 |  |  | 0.6∗ | 1500 | – | – | 0.06/0.002 | – | – | 1/7 | 3e-7 | Large grains, 50 au. Surface density
|  |  |  |  |  |  |  |  |  |  |  |  | power law pΣdust=-1.2

###  3.4 ECHA J0801

ECHA J0801 (2MASS J08014860) is one of the M5 sources classified by Murphy et
al. (2010) as possible halo members of $\eta$ Cha, and it may be a binary
system due to its position in the color-magnitude diagram and the different
high-velocity measurements (Murphy et al. 2010). Dynamical simulations
suggested the system was ejected from the cluster.
Although ECHA J0801 has a negligible excess only in the mid-IR, we find a
significant submillimeter flux, which might suggest a massive disk. The
presence of substantial circumstellar material can also explain the observed
variable accretion (Murphy et al. 2011). We highlight, however, that this high
submillimeter flux can be contaminated by the local diffuse nebulosity
detected by Planck between 30 and 217 GHz. Taking into account this probable
contamination, the submillimeter detection will be considered as an upper
limit for disk emission. However, since the nature of such a system is still
under debate, we model the SED using both RADMC 2D and DMS.

Since the stellar photosphere had not been modeled previously, we use the
stellar photosphere of ECHA J0841 (with a slightly different radius to account
for the differences in luminosity) to simulate the stellar emission, given
that the two stars are very similar. Using RADMC 2D, the lack of far-IR
observations does not allow us to pinpoint the disk structure at intermediate
radial distances, but the object is consistent with a nearly depleted inner
disk together with a rather massive outer disk, starting some point after 12
au. Given its very low mass, it belongs to the class of transitional objects
with very large gaps, which is rare around late-type stars.
Using instead DMS, ECHA J0841 is assumed to be a debris disk with an optically
thin dust emission. With the same approximation as in the other debris
systems, we find that ECHA J0841 is the only source with a massive narrow ring
of 1.6-1.7 $\times$ 10-6M⊙ at between 17 and 20 au from the central star.
It is very interesting to highlight that both approaches led to a large gap of
about 15 au and a very massive disk compared to other systems (the absolute
disk mass is not comparable since the grain distribution and disk size are
different).

###  3.5 The age of the $\eta$ Cha association

---
Figure 5: Top: Colour-Magnitude diagram of the $\eta$ Cha association members
from Gaia eDR3 photometry. Isochrones are overplotted and determined
$\chi^{2}$ values are shown in the legend, each with $p$ values ¿ 9e6.
Confirmed binaries, as indicated in Tab. 3 are not included in the $\chi^{2}$
calculations. Bottom: Proper motion plot from Gaia eDR3 astrometry. Markers
are colored by the star’s parallax, as shown by the color bar. RX J1005
(-38,15) and J0801 (-19,29) are not shown.

Recent work on Gaia observations of the nearby associations (Gagné et al.
2018) reported an age of 11 $\pm$ 3 Myr for the $\eta$ Cha association. This
value was not based on a new analysis of the available Gaia DR2 data presented
in that paper, but instead on the previous work of Bell et al. (2015). They
compiled the MV, V-J color-magnitude diagram (CMD) of 18 cluster members,
applying a weighted average of a few stars within the cluster with
trigonometric parallaxes from the revised Hipparcos reduction of van Leeuwen
(2007), giving a distance of 94.27 $\pm$ 1.18 pc. The final given age is the
average from the fitting of four different sets of isochrones (Dotter et al.
2008; Tognelli et al. 2011; Bressan et al. 2012; Baraffe et al. 2015).

We have used the Gaia eDR3 photometric fluxes and parallaxes (Gaia
Collaboration et al. 2021), finding the best age for the association to be 8
$\pm$ 1 Myr (Fig. 5). The mean parallax for the association from eDR3 is 10.11
$\pm$ 0.13 mas, resulting in a distance of $\sim$ 99.0 $\pm$ 1.3 pc. Absolute
G band magnitudes were calculated using the individual parallaxes for 17 of
the cluster members and the three halo members. The only star that does not
have a parallax measurement in eDR3 is RECX-9, for which we use the average
parallax. PARSEC (Bressan et al. 2012) and updated COLIBRI (Pastorelli et al.
2020) isochrones were over-plotted to the CMD, with a standard $\chi^{2}$
calculation giving the best fitting age of 8 Myr. All $\chi^{2}$ calculations
had a $p$-value ¿ 9e6. However, due to the intrinsic variability of PMS stars
and the corresponding expected to scatter, 7 and 9 Myr are also likely with
their respective low $\chi^{2}$ scores. The five confirmed binaries in the
association (as indicated in Tab. 3) were not included in the $\chi^{2}$
calculations. Each of these binary members is shown to lie above the youngest
isochrone in Fig. 5. Alecian et al. (2007) calculate the age of the eclipsing
binary RS Cha to be 9.13 $\pm$ 0.12 Myr, which is in good agreement with our
estimate for the rest of the association. J0838 and J0836 are also slightly
above the youngest isochrone.

Two members are positioned below the oldest isochrones, J0843 and J0844. As
discussed, these two stars have extended disks with similar properties.
Intrinsic variability of the stars or occultations from disk material may be
the reason for their G band flux being lower than expected for the age of the
association. Another possible explanation could be accretion, which strongly
affects the spectra of both sources (Fig. 7 in Rugel et al. 2018), or
scattered light from the large disks, which can make objects appear bluer than
expected. The three halo members, which are further away from the cluster
center, have similar ages to the rest of the association. J0820 is located in
the same proper motion space as the association, however, RX J1005 and J0801
are outliers in the proper motion plot.

##  4 Discussion

###  4.1 Disk properties

Models of the SEDs of protoplanetary and debris disk systems in $\eta$ Cha
allow us to characterize the disk geometry and dust distribution. In
particular, none of the observed disks can be well-fitted using a non-evolved
disk model, with well-mixed dust and gas (no settling) and a smooth
distribution of dust across the disk. The SEDs of all of the disks in $\eta$
Cha show clear deviations from a standard protoplanetary disk model, such as
radial variations of the dust properties, settling, mass depletion, and the
presence of radial asymmetries or substructures(holes, gaps, and/or changes in
the vertical scale height) in the disk.

As expected, individual disk models are highly degenerated and the parameter
space is often non-continuous. We thus prefer to keep the models simple and
explore the deviations from a standard, continuous disk rather than to obtain
good fits. Note that in most cases, we expect the structure of the disk to be
much more complicated (e.g., containing radial and azimuthal asymmetries,
structures within gaps and holes and at the disk edges, etc.) than what is
included in our simple models.

Debris disks in the $\eta$ Cha association are found around late spectral
types, from K7 (RECX 1) and M-type stars. At first approximation, the modeling
of the debris disks (see Fig. 4) assumes a ring distribution of the debris
dust. The SED models allow us to constrain a range of possible locations and
ring masses. With masses between 0.1 and 1.8 $\times$ 10-6 M⊙ the ring
location varies between 15 and 100 au from the central star. The most massive
ring, about $\sim 10^{-5}$ M⊙, is found at more than 600 au around RX J1005.3
- 7749. A less massive ring around RECX 1 is also found at very large
distances, between 200 and 600 au from the central star. These results are
fundamental to further exploring whether the origin of these structures may be
related to the presence of an inner solar system.

###  4.2 Dust disk mass versus mass accretion

---
Figure 6: Top: Accretion rate vs spectral type. Note the lack of significant
correlations between these quantities. Bottom: Circumstellar dust mass derived
according to Eqn. 1 vs. accretion rate. For comparison, the objects in the
list from Manara et al. (2023) are shown in grey. Note that the total disk
mass would depend on the gas-to-dust ratio. APEX flux and mass upper limits
are marked by inverted triangles. The red line corresponds to the model fitted
by Manara et al. (2016). The accretion rates are taken from the literature
(see Table 3). For each object, we show the individual measurements (small
symbols) together with the average rate (large symbols). A further color ring
is added to specify the type of disk according to the classical SED
classification.

We first discuss the results on $\eta$ Cha alone, then in the context of other
young sources (compiled by Manara et al. 2023) to cover a larger spread in
disk dust masses and mass accretion rate.

The relation between submillimeter emission (or disk mass) and accretion has
been suggested by several authors, as a means to explore and test viscous
accretion and photoevaporation (e.g., Hartmann 1998; Sicilia-Aguilar et al.
2013a; Manara et al. 2016, among others). The disks in $\eta$ Cha have the
advantage of being older than most, which allows us to explore whether this
relation holds for slow accretors with evolved disks. Table 3 lists the mass
accretion rates from the literature. To avoid introducing further
uncertainties from the various types of disk models that can fit a given SED,
and to make our results comparable to those of Manara et al. (2016), we
derived the disk mass following Equation 1 with the same values they use
($\beta$ = 1, kν0 = 10 cm2g-1 at $\nu_{0}$ = 1000 GHz, and Tdust = 20 K).
Because of the differences in the underlying dust model assumed to obtain the
dust absorption coefficient, these masses are different from the masses
derived from the RADMC models, which assume a larger maximum grain size and
calculate the dust temperature consistently from radiative transfer depending
on the disk model and structure. Note that the available data do not constrain
the maximum grain sizes, so the differences reflect the typical uncertainties
deriving disk masses.

Spurious correlations may appear because accretion rates and disk masses tend
to be correlated with the stellar mass (e.g., Clarke & Pringle 2006; Mulders
et al. 2017). We thus checked whether this could be the main cause of the
marginal correlation, but we found no correlation between the spectral type
and the APEX flux (see Fig. 1) or the accretion rate (see Fig. 6). The
relation between the accretion rate and the disk mass is shown in the bottom
panel of Fig. 6. The $\eta$ Cha data are consistent with the Mdust vs Ṁ
correlation proposed by Manara et al. (2016), although our data points
themselves are not significantly or only marginally correlated (when upper
limits are included, the false-alarm probability for a Spearman rank test is
7%). This is likely caused by the large scatter of the mass-accretion relation
and the fact that we have very few data points concentrated towards the lower
end of mass and accretion rate ranges.

We also note that there is no evidence in the literature of accretion onto any
of the sources that do not have a significant IR excess and that sources with
disks classified as Class III (which also correspond to anemic or dust-
depleted disks; Lada et al. 2006; Currie et al. 2009; Sicilia-Aguilar et al.
2013b), have (if any) very low accretion rates. This is the same trend
observed in other clusters, where the lack of detectable accretion is strongly
correlated with clear signs of disk evolution and substantial mass depletion
(e.g., Sicilia-Aguilar et al. 2013b). The presence of very low accretion rates
in objects with disks that, albeit evolved and with radial asymmetries, are
not cleared from gas and dust in their innermost regions, also poses strong
constraints to the photoevaporation rate. The fact that all disks with
significant disk mass are accreting is further confirmation that, in order to
stop accretion, a dramatic change in the disk has to occur, and simple holes
and gaps are normally not enough if not accompanied by a clear decrease in the
disk mass.

Putting together all observations of accretion, one interesting fact is that
accretion towards $\eta$ Cha members seems to be highly variable, despite its
age. Although there are only 2 or 3 measurements per object, 7 out of 12 stars
in $\eta$ Cha have accretion variations of one order of magnitude or more. In
contrast, only 10% of stars in Tr 37 have variable accretion (beyond a factor
of few and at least over 3 $\sigma$) when comparing the existing 2
measurements of a larger collection of members (Sicilia-Aguilar et al. 2013a).
Examining the results in more detail, several of the objects show variable
line profiles (e.g., Ingleby et al. 2013; Lawson et al. 2004). Nevertheless,
variability in the line profile does not necessarily indicate variations in
the accretion rate and could be instead due to rotational modulation of non-
axisymmetric accretion channels (Campbell-White et al. 2021). Considering
this, together with the fact that detecting very low accretion rates (for
which the accretion luminosity is very low compared to the stellar luminosity)
is intrinsically very hard, we prefer to be cautious in this respect.

The $\eta$ Cha members also suffer the known problem that the observed disk
mass (as extrapolated from dust mass, assuming a typical gas-to-dust ratio) is
not enough to support accretion for a very long time. In fact, assuming the
typical gas-to-dust ratio of 100 and considering the observed accretion rates,
all the disks in $\eta$ Cha would need to entirely disperse within hundreds to
thousands of years to support the current accretion rates. Although these
relatively old and evolved disks are not expected to live very long,
considering the typical evolutionary timescales (Haisch et al. 2001; Sicilia-
Aguilar et al. 2006; Hernández et al. 2007; Fedele et al. 2010), these
estimates are too short. Moreover, all disks are consistent with having their
bulk mass at a relatively large distance from the star, which makes it
difficult the transport within this short timescale, predicting even shorter
lifetimes for disks and accretion. This is statistically implausible because
it would mean that all disks would be about to dissipate, in contradiction
with the usual smooth disk fraction vs. time curve. Even though the mismatch
between accretion and disk masses is a known problem and it tends to affect
less the lower-mass end (e.g., Sicilia-Aguilar et al. 2016), the data on the
$\eta$ Cha members is a further indication that something is missing in the
global picture of accretion and mass transport in disks.

Looking at the young objects from Manara et al. (2023), Fig. 6 reveals that
our sources are complementary to the sources from the literature. Taking into
account the mass accretion variability, mass accretion and disk mass
distribution of $\eta$ Cha does not allow a clear distinction of disk
evolutionary path between viscosity, MHD winds, or external photoevaporation
as suggested in Manara et al. (2023). We decided on purpose to not specify the
age of the sources from Manara et al. (2023) since those are not directly
comparable in any case with an older region as $\eta$ Cha where there is an
intrinsic decreasing of the disk fraction already took place: the surviving
accreting sources at such an advanced age would need to be compared only with
the same percentage of disks in younger clusters and, as pointed out by
Sicilia-Aguilar et al. (2010), their younger counterparts would probably
correspond to a small subset of those with initially higher disk masses.

###  4.3 Disk evolution at longer wavelengths

The millimeter observations and the age of the cluster allow us to discuss the
dissipation of the disks within the first 5 Myr, which represents a key age of
the planet formation process since after that age the amount of inventory
material is more than halved. Our survey includes all the core members of the
$\eta$ Cha associations, not only the Class II sources, plus three halo
members from the works of Murphy et al. (2010, 2011). The disk fraction for
$\eta$ Cha in the infrared is based on the Spitzer work of Sicilia-Aguilar et
al. (2009) and Fang et al. (2013), including the three halo members.

Since, to a first approximation, the millimeter emission from a protoplanetary
disk is tracing its outer region, while the mid-infrared emission traces the
inner part, the comparison between the disk fraction in these two wavelength
ranges is directly related to the mechanism responsible for the disk
dissipation.

We find that all the sources with a near-infrared excess do show a millimeter
excess over the stellar photosphere. The only exception is $\eta$ Cha itself,
which is not a reliable detection as previously discussed (see Section 3.2.8).
This suggests that the dissipation processes affect the disk at all radii, as
suggested by Currie & Sicilia-Aguilar (2011).

A further discussion of this hypothesis can be done using a recent work of
Michel et al. (2021), where they summarized the updated membership (including
recent results from Gaia) and redefined disk classes in a consistent way.
Using the definition of disk classes from Greene et al. (1994) based on the
slope in the SED in the 2.2-10 $\mu$m, Michel et al. (2021) derived also the
total number of Class I, Class F, and Class II, and the Class III in several
young star-forming regions: Corona Australis (hereafter labeled Cr A),
Ophiuchus, Taurus, Chamaeleon I (Cha I), Chamaeleon II (Cha II), IC 348,
Lupus, $\epsilon$ Chamaeleontis ($\epsilon$ Cha), TW Hydra (TW Hya), $\eta$
Cha, and Upper Sco.
While Michel et al. (2021) analyzed the possible different mechanisms
responsible for dust dissipation in disks with and without structures, we
compiled instead the fraction of sources classified as Class I, Class F, and
Class II and compared them to the fraction of Class III and debris disks. For
consistency, we mostly adopt the compiled values from Michel et al. (2021).
However, when comparing disk masses to debris disks Michel et al. (2021)
refers to a more complete survey on debris disks carried out by Holland et al.
(2017), where, apart from the TW Hya debris disks, all the others are older
than 20 Myrs and most of them isolated field objects. For this reason, we use
some additional criteria to compile the final disk fractions which differ from
Michel et al. (2021). The age of $\eta$ Cha is taken from our work, $\epsilon$
Cha from Murphy et al. (2010), TW Hya from Luhman (2023), and Cr A from
Sicilia-Aguilar et al. (2011) with a spectroscopic survey of the in-cloud
population. The fraction of Class I, F, II, and III for Cr A arises from
Peterson et al. (2011), which includes only the in-cloud core population of
the cluster members spectroscopically characterized. The debris disk fraction
for $\eta$ Cha results from this work, and the one of TW Hya, from the SONS
survey (Holland et al. 2017). Following the definition of debris disks used in
our work, we also compile the confirmed debris disks from Michel et al. (2021)
selecting Class III disks with fractional dust luminosities
${f_{\rm\,dust}\,(\equiv L_{\rm\,dust}/L_{\rm\,*})\,<\,10^{\,-2}}$.

Fig. 7 shows the fraction of Class I + Class F + Class II disks and Class III
disks of young star-forming regions (upper panel), and the lower limit of
debris disks in these regions (bottom panel). We find that the fraction of
Class III is lower than Class I, F, and II in CrA, Ophiuchus, Cha I, Cha II,
and Lupus, while the inverse behavior is seen in Taurus, IC 348, $\epsilon$
Cha, $\eta$ Cha, TW Hya, and Upper Sco, where the fraction of Class III is
higher. We also find an exponential decrease in the fraction of Class I, F,
and II with a slope of -0.32 $\pm$ 0.11, with the fraction of such disks
decreasing together up to ages 2 to 3 Myrs. After this timescale, the drastic
decrease in the disk fractions of Class I, F, and II disks corresponds to an
increase in Class III. While this is consistent with previous studies on dust
and gas dissipation in low-mass star-forming regions at different ages (e.g.,
Hernández et al. 2007; Fedele et al. 2010; Williams & Cieza 2011), our work
also explores whether the percentage of debris disks among Class III objects
allow us to further constrain formation and evolution timescales.

For the younger regions, we find that most of the clusters show only a few
percent of debris disks, without allowing proper treatment of the timescale
for debris disk formation and evolution. However, our survey on $\eta$ Cha
opens a new discussion on this class of objects showing a percentage of debris
disks ten times higher than in other regions. A result in this direction is
drawn also including the results of the debris disks in TW Hya from the
youngest sources in the SONS survey (explicitly dedicated to debris disks;
Holland et al. 2017). Furthermore, we find a constant lower limit for debris
disks until 8 Myrs, where a much higher percentage of debris disks is found in
$\eta$ Cha. We note that this can be, however, an effect of a bias in the
observations of the younger regions that might also host a higher number of
debris disk systems, not yet detected.

Nonetheless, this general trend might have some fundamental implications for
the timescale of grain growth, dust trapping, and giant planet formation. By 5
Myr, a significant amount of mass will be already as a building block of
planets and few sources will have evolved dust up to millimeter sizes. If,
instead, a higher fraction of debris disks were to be found in younger
regions, it would mean that all these processes can efficiently create stable
systems already in a few Myrs.

---
Figure 7: Relation between disk fraction in several young star-forming regions
and their age. The fraction of Class I + Class F + Class II disks is marked by
filled dots, while the fraction of Class III disks is shown as crosses. Each
star-forming region is color-coded according to the legend. The filled
triangles in the lower panel represent the lower limits of the debris disk
fraction. The blue line represents the fitting of the fraction of Class I +
Class F + Class II disks, with a slope of -0.32 $\pm$ 0.11.

There are some aspects that need to be taken into consideration, as they may
affect the disk fractions. Notably, advancements in observational
capabilities, particularly through Gaia (e.g., Galli et al. 2018;
Roccatagliata et al. 2018; Herczeg et al. 2019; Roccatagliata et al. 2020;
Luhman 2023), have revealed multiple populations within star-forming regions.
To date, the treatment of age and the incorporation of updated disk fractions
are only partially addressed. Recent reanalysis of disk fractions with revised
membership from Gaia (e.g., Mendigutía et al. 2022) recover consistent disk
fractions (with discrepancies up to 10%), despite considering a more extensive
area of star-forming regions. Combining different techniques to complete the
cluster membership in different spectral type ranges, as done by Pelayo-
Baldárrago et al. (2023), also found that the disk fraction in IC 1396
decreased from 39% down to 28$\%$ at about 4 Mys.

##  5 Conclusions and Summary

The main results and conclusions of our study can be summarized as follows.

  * •

Thanks to its proximity, our pilot submillimeter survey carried out with the
APEX/LABOCA instrument detected 12 (two of them marginally) out of the 20
members of the $\eta$ Cha cluster, providing valuable upper limits for the
rest. Our work also confirms the distance of the $\eta$ Cha association to be
at 99.0 $\pm$ 1.3 pc, with an age between 7 and 9 Myr.

  * •

The modeling of the SEDs, including the new submillimeter data, suggests that
protoplanetary disks in $\eta$ Cha have holes with radii between 0.01 and 0.03
AU, while ring-like emission from the debris disks arises from between 20 au
and 650 au from the central star. In the case of ECHA J0801, both
protoplanetary disk and debris disk models led to a large gap of about 15 au
and a very massive disk compared with other systems. We confirm the trend that
low disk masses and low accretion rates are correlated, as well as the fact
that having gaps and holes is not enough to stop disk accretion, if not
accompanied by a substantial decrease in the disk mass.

  * •

All disks in the cluster depart from the typical structure of primordial
protoplanetary disks, showing evidence of radial changes in their structures.
Although most disks observed at high resolution have signatures of radial and
azimuthal asymmetries (so that their presence may not be representative of a
rapid evolutionary stage, but may also correspond to long-lived features), for
$\eta$ Cha disks these asymmetries are already evident in the SED, which
indicates a further degree of evolution. Disk dispersal is found to be
dominated by inside-out processes.

  * •

Comparing accretion rates and disk masses, the classical problem of lack of
disk matter to support accretion is also present in $\eta$ Cha, extending well
into the low-mass regime where the problem tends to be less serious than for
intermediate-mass stars.

  * •

Our work on the $\eta$ Cha association helps put in context the evolution of
disks by addressing all the different types of disks and their changes with
age. An important point is that we include the debris disks, for which so far
there is only a lower limit. Our study shows consistency with previous results
concerning the fractions of Class I, F, Class II, and Class III objects, while
also bringing up the discussion on the formation and evolution of debris
disks. Their higher fraction in older clusters suggests that there are
physical mechanisms capable of efficiently forming these systems in less than
5 Myr.

###### Acknowledgements.

VR acknowledges the support of the Italian National Institute of Astrophysics
(INAF) through the INAF GTO Grant “ERIS $\&$ SHARK GTO data exploitation”. ASA
and JCW were in part supported by the STFC grant number ST/S000399/1 (“The
Planet-Disk Connection: Accretion, Disk Structure, and Planet Formation”). JCW
is funded by the European Union under the European Union’s Horizon Europe
Research $\&$ Innovation Programme 101039452 (WANDA). MK and SW acknowledge
the DFG for financial support under contract WO 857/15-1 and WO 857/15-2
within the context of the Research UnitFOR 2285 “Debris Disks in Planetary
Systems”. MK also acknowledges the funding from the Royal Society. This
publication is based on data acquired with the Atacama Pathfinder Experiment
(APEX). APEX is a collaboration between the Max-Planck-Institut für
Radioastronomie, the European Southern Observatory, and the Onsala Space
Observatory. This research has made use of the SIMBAD database, operated at
CDS, Strasbourg, France.

##  Appendix A Additional Material

In this Appendix, we present the extended Tables of all the data used during
the discussion in the main paper.

Table 3: Mass accretion rates compiled from the literature. The accretion rates labeled as Ṁmodel are derived from magnetospheric accretion model fits, while rates marked as ṀV10 are obtained from the relation between accretion rate and the H$\alpha$ line width at 10% of the maximum (V10), from Natta et al. (2004). The uncertainties in the disk mass include detection and calibration uncertainties. ECHA 0841 has TO/flat SED, EW(Ha) = -12 AA. ECHA 0844 has flat SED and is likely accreting with EW(Ha) = -58 AA References: J06: Jayawardhana et al. (2006); L04: Lawson et al. (2004); M11: Murphy et al. (2013); I13: Ingleby et al. (2013); R18: Rugel et al. (2018); * = V10 below 200 km s-1, which is the minimum for the Natta et al. (2004) relation. Source |  log(Ṁmodel) | Refs. |  log(ṀV10) | Refs.
---|---|---|---|---
|  (M⊙/yr) |  |  (M⊙/yr) |
RECX-1 | -11.9 | J06 |  |
RECX-3 | -12.1* | L04 | -11.8* | J06
RECX-4 | -11.9* | L04 | -11.5* | J06
RECX-5 | -10.3 | L04 | -9.7,-11.0*, -9.89 | L04, J06, R18
RECX-6 | -11.4* | L04 | -11.5* | J06
RECX-7 | -11.7* | L04 | -10.1 | J06
RECX-9 | -10.4 | L04 | -10.0,-9.1, -9.92 | L04, J06, R18
RECX-10 |  |  | -11.8*, -11.9* | L04, J06
RECX-11 | -10.4,-9.77 | L04,I13 | -9.7, -9.5, -9.49 | L04, J06, R18
RECX-12 |  |  | -11.3*, -11.4* | L04,J06
ECHA J0843 | -9.0,-9.10 | L04, I13 | -7.7,-8.7,-9.12 | L04 , J06
ECHA J0844 |  |  | -10.18 | R18
2MASS J0801 | -11.3*, -9.7 | M11 |  |
2MASS J0820 | -10.6, -9.7 | M11 | -8.7 | M11
Table 4: Log of the observations of the $\eta$ Cha members observed in on-off mode. Source | Obs. Date | Scan # | Exp. time [s]
---|---|---|---
$\eta$ Cha | 2010-10-30 | 81846 | 600
|  | 81847 | 600
|  | 81854 | 600
J0801 | 2010-12-11 | 91881 | 600
|  | 91882 | 600
|  | 91883 | 600
|  | 91884 | 600
|  | 91885 | 600
|  | 91886 | 600
|  | 91887 | 600
J0841 | 2010-12-10 | 91567 | 600
|  | 91568 |
|  | 91569 |
|  | 91570 |
|  | 91571 |
|  | 91572 |
J0843 | 2010-12-10 | 81872 | 600
|  | 81873 |
|  | 96114 |
|  | 96115 |
|  | 96116 |
|  | 96117 |
|  | 96118 |
J0844 | 2010-12-10 | 91697 | 600
|  | 91698 |
|  | 91699 |
| 2010-12-10 | 91700 | 600
| 2010-12-27 | 96127 | 600
|  | 96128 |
|  | 96129 |
|  | 96130 |
|  | 96131 |
RECX-3 | 2010-12-26 | 95914 | 600
|  | 95915 |
|  | 95916 |
|  | 95917 |
|  | 95918 |
|  | 95919 |
|  | 95920 |
RECX-4 | 2010-12-26 | 95928 | 600
|  | 95929 |
|  | 95930 |
|  | 95931 |
|  | 95932 |
|  | 95933 |
RECX-5 | 2010-10-30 | 81865 | 600
|  | 81867 |
|  | 81868 |
Table 5: Journal of the APEX/LABOCA observations of $\eta$ Cha members obtained in on-off mode. Source | Obs. Date | Scan # | Exp. time [s]
---|---|---|---
RECX-6 | 2013-11-18 | 82368 | 360
|  | 82369 | 600
|  | 82370 | 600
|  | 82371 | 600
|  | 82372 | 600
|  | 82373 | 60
| 2013-11-21 | 83354 | 600
|  | 83355 | 600
|  | 83356 | 600
|  | 83357 | 600
RECX-7 | 2013-11-19 | 82822 | 600
|  | 82823 | 600
| 2013-11-20 | 83105 | 600
|  | 83106 | 600
RECX-9 | 2010-10-30 | 81869 | 600
|  | 81870 |
|  | 81871 |
| 2013-11-19 | 82824 | 600
|  | 82825 | 600
| 2013-11-20 | 83103 | 600
|  | 83104 (no) | 600
RECX-11 | 2010-10-30 | 81858 | 600
|  | 81859 |
|  | 81860 |
RECX-12 | 2013-11-18 | 82347 | 600
|  | 82348 | 600
|  | 82349 | 570
|  | 82350 | 600
|  | 82351 | 600
|  | 82352 | 330
| 2013-11-21 | 83361 | 600
|  | 83362 | 600
RS Cha | 2013-11-18 | 82375 | 600
|  | 82376 | 600
HD75505 | 2013-11-19 | 82820 | 600
|  | 82821 | 600
Table 6: Journal of the observations of the $\eta$ Cha members in mapping mode (we report only the scans that have been analyzed with more than one sub-scan.) Source | Obs. Date | Scan # | Exp. time [s]
---|---|---|---
RECX-16 | 2008-10-21 | 55435 | 35
|  | 55436 | 35
|  | 55437 | 35
|  | 55438 | 35
|  | 55439 | 35
|  | 55440 | 35
|  | 55441 | 35
|  | 55442 | 35
RECX-17 | 2008-12-26 | 73721 | 420
|  | 73722 | 420
|  | 73723 | 420
|  | 73724 | 420
|  | 73725 | 420
|  | 73726 | 420
|  | 73727 | 420
|  | 73728 | 420
RECX-18 | 2008-12-26 | 73697 | 420
|  | 73698 | 420
|  | 73699 | 420
|  | 73700 | 420
|  | 73701 | 420
|  | 73702 | 420
|  | 73703 | 420
|  | 73704 | 420

Table 7: Multi-wavelength photometry including the near- and mid-infrared data from Sicilia-Aguilar et al. (2009) and far-infrared from Riviere-Marichalar et al. (2015). Name | 2MASS ID |  3.6 $\mu$m |  4.5 $\mu$m |  5.8 $\mu$m |  8.0 $\mu$m |  24 $\mu$m |  70 $\mu$m |  70 $\mu$m |  100 $\mu$m |  160 $\mu$m |  870 $\mu$m
---|---|---|---|---|---|---|---|---|---|---|---
|  |  Spitzer/IRAC | IRAC | IRAC | IRAC |  Spitzer/MIPS | MIPS |  Herschel/PACS | PACS | PACS | LABOCA
$\eta$ Cha | 08411947-7857481 | a | a | a | a |  4.33 $\pm$ 0.02 |  3.87 $\pm$ 0.13 |  $<$6 |  |  $<$16 |  3.37$\pm$1.86
RS Cha | 08431222-7904123 | a | a |  5.40 $\pm$ 0.01 |  5.40 $\pm$ 0.01 |  5.34 $\pm$ 0.03 |  |  |  |  |  18.62$\pm$2.63
HD 75505 | 08414471-7902531 |  6.97 $\pm$ 0.01 |  6.97 $\pm$ 0.01 |  6.97 $\pm$ 0.01 |  6.97 $\pm$ 0.01 |  6.97 $\pm$ 0.06 |  |  $<$7 |  |  $<$19 |  ¡7.86
J0836 | 08361072-7908183 |  10.57 $\pm$ 0.03 |  10.45 $\pm$ 0.01 |  10.40 $\pm$ 0.01 |  10.37 $\pm$ 0.01 | b | b |  $<$10 |  |  $<$24 |
J0838 | 08385150-7916136 |  10.07 $\pm$ 0.02 |  10.00 $\pm$ 0.01 |  9.92 $\pm$ 0.01 |  9.90 $\pm$ 0.01 | b | b |  $<$8 |  |  $<$15 |
J0841 | 08413030-7853064 |  10.53 $\pm$ 0.01 |  10.26 $\pm$ 0.05 |  10.04 $\pm$ 0.01 |  9.48 $\pm$ 0.01 |  7.07 $\pm$ 0.06 |  |  $<$15 |  $<$12 |  $<$25 |  2.88$\pm$0.97
J0843 | 08431857-7905181 |  8.38 $\pm$0.03 |  7.91 $\pm$ 0.01 |  7.42 $\pm$ 0.03 |  6.51 $\pm$ 0.01 |  3.52 $\pm$ 0.01 |  1.89 $\pm$ 0.02 |  204$\pm$6 |  142$\pm$5 |  71$\pm$7 |  5.42$\pm$0.83
J0844 | 08440914-7833457 |  11.02 $\pm$0.01 |  10.75 $\pm$ 0.01 |  10.42 $\pm$ 0.01 |  9.76 $\pm$ 0.01 |  7.11 $\pm$ 0.07 | b |  29$\pm$4 |  |  59$\pm$8 |  4.96$\pm$0.83
RECX-1 | 08365623-7856454 |  7.17 $\pm$ 0.02 |  7.20 $\pm$ 0.02 |  7.16 $\pm$ 0.01 |  7.11 $\pm$ 0.02 |  6.96 $\pm$ 0.06 |  |  |  |  |
RECX-3 | 08413703-7903304 |  9.27 $\pm$ 0.02 |  9.21 $\pm$ 0.02 |  9.09 $\pm$ 0.05 |  9.15 $\pm$ 0.01 |  8.49 $\pm$ 0.13 |  |  6$\pm$1⋆⋆ |  $<$10 |  $<$13 |  $<$1.65
RECX-4 | 08422372-7904030 |  8.45 $\pm$ 0.02 |  8.45 $\pm$ 0.02 |  8.37 $\pm$ 0.01 |  8.32 $\pm$ 0.02 |  7.73 $\pm$ 0.09 |  |  9$\pm$1⋆⋆ |  $<$9 |  $<$14 |  $<$3.03
RECX-5 | 08422710-7857479 |  9.59 $\pm$ 0.03 |  9.50 $\pm$ 0.03 |  9.37 $\pm$ 0.01 |  8.89 $\pm$ 0.01 |  5.05 $\pm$ 0.02 |  2.38 $\pm$ 0.04 |  113$\pm$24 |  203$\pm$13 |  138$\pm$28 |  20.01$\pm$2.45
RECX-6 | 08423879-7854427 |  9.15$\pm$ 0.04 |  9.07 $\pm$ 0.07 |  9.04$\pm$ 0.01 |  9.04 $\pm$ 0.01 |  8.88 $\pm$ 0.16 |  |  $<$5 |  |  $<$13 |  8.38$\pm$0.54
RECX-7 | 08430723-7904524 |  7.52$\pm$ 0.01 |  7.54 $\pm$ 0.02 |  7.49$\pm$ 0.01 |  7.46 $\pm$ 0.01 |  7.37 $\pm$ 0.08 |  |  $<$15 |  $<$11 |  $<$21 |  $<$4.02
RECX-9 | 08441637-7859080 |  8.99$\pm$ 0.02 |  8.80 $\pm$ 0.01 |  8.57$\pm$ 0.01 |  7.97 $\pm$ 0.02 |  5.8 $\pm$ 0.03 |  3.16$\pm$0.07 |  |  77$\pm$3 |  53$\pm$7 |  $<$5.46
RECX-10 | 08443188-7846311 |  8.58$\pm$ 0.01 |  8.61 $\pm$ 0.01 |  8.55$\pm$ 0.02 |  8.49 $\pm$ 0.01 |  8.45 $\pm$ 0.13 |  |  $<$7 |  |  $<$14 |
RECX-11 | 08470165-7859345 |  7.09$\pm$ 0.02 |  6.86 $\pm$ 0.01 |  6.57$\pm$0.01 |  5.97$\pm$0.01 |  3.68$\pm$0.01 |  1.80$\pm$0.02 |  212$\pm$21 |  205$\pm$13 |  174$\pm$26 |
RECX-12 | 08475676-7854532 |  8.19$\pm$0.01 |  8.15 $\pm$ 0.01 |  8.10$\pm$ 0.02 |  8.07$\pm$0.01 |  7.95$\pm$0.10 |  |  $<$5 |  |  $<$13 |  $<$2.01
666⋆⋆:Cieza et al. (2013). |  |  |  |  |  |  |  |  |  |  |
Figure 8: Sketch of the basic idea of the model used to interpret the SED of
the debris disks.

We show here the results of the tentative interpretation of the SEDs of RECX 6
and RECX 8. Both sources are well detected with LABOCA at 870 $\mu$m above 3
$\sigma$.

|
---|---
Figure 9: Tentative SED models for the debris disks. The photometry data are
marked by red circles and red triangles (the upper limits). Stellar
photospheric, disk, and total emission (stellar photospheric + disk models)
are represented by grey dotted lines, and red and blue continuous lines,
respectively, but none of those are able to reproduce all the observed data
well.

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