# EXORCISM: a spectroscopic survey of young eruptive variables (EXor and
candidates)

T. Giannini11affiliation: INAF - Osservatorio Astronomico di Roma, via
Frascati 33, 00078, Monte Porzio Catone, Italy , A. Giunta22affiliation: ASI -
Agenzia Spaziale Italiana - Via del Politecnico, 00133 Roma, Italy , M.
Gangi11affiliation: INAF - Osservatorio Astronomico di Roma, via Frascati 33,
00078, Monte Porzio Catone, Italy , R. Carini11affiliation: INAF -
Osservatorio Astronomico di Roma, via Frascati 33, 00078, Monte Porzio Catone,
Italy , D. Lorenzetti11affiliation: INAF - Osservatorio Astronomico di Roma,
via Frascati 33, 00078, Monte Porzio Catone, Italy , S.
Antoniucci11affiliation: INAF - Osservatorio Astronomico di Roma, via Frascati
33, 00078, Monte Porzio Catone, Italy , A. Caratti o Garatti33affiliation:
INAF - Osservatorio Astronomico di Capodimonte, Via Moiariello 16, 80131
Napoli, Italy 44affiliation: Dublin Institute for Advanced Studies, School of
Cosmic Physics, Astronomy & Astrophysics Section, 31 Fitzwilliam Place, Dublin
2, Ireland , L. Cassará55affiliation: INAF-IASF Milano, via Alfonso Corti 12,
20129 Milano, Italy , B. Nisini11affiliation: INAF - Osservatorio Astronomico
di Roma, via Frascati 33, 00078, Monte Porzio Catone, Italy , A.
Rossi66affiliation: INAF-Osservatorio di Astrofisica e Scienza dello Spazio,
Via Piero Gobetti, 93/3, I-40129 Bologna, Italy , V. Testa11affiliation: INAF
- Osservatorio Astronomico di Roma, via Frascati 33, 00078, Monte Porzio
Catone, Italy , F. Vitali11affiliation: INAF - Osservatorio Astronomico di
Roma, via Frascati 33, 00078, Monte Porzio Catone, Italy

(Received June 1, 2019; Revised January 10, 2019)

###### Abstract

We present an optical/near-IR survey of 11 variable young stars (EXors and
EXor candidates) aimed at deriving and monitoring their accretion properties.
About 30 optical and near-infrared spectra ($\Re$ $\sim$ 1500-2000) were
collected between 2014-2019 with the Large Binocular Telescope (LBT). From the
spectral analysis we have derived the accretion luminosity ($L_{\mathrm{acc}}$
) and mass accretion rate ($\dot{M}_{\mathrm{acc}}$ ), the visual extinction
(AV) , the temperature and density of the permitted line formation region (T,
$n_{\mathrm{H}}$), and the signature of the outflowing matter.

Two sources (ASASSN-13db and iPTF15afq) have been observed in outburst and
quiescence, three during a high-level of brightness (XZ Tau, PV Cep, and NY
Ori), and the others in quiescence. These latter have $L_{\mathrm{acc}}$ and
$\dot{M}_{\mathrm{acc}}$ in line with the values measured in classical T Tauri
stars of similar mass. All sources observed more than once present
$L_{\mathrm{acc}}$ and $\dot{M}_{\mathrm{acc}}$ variability. The most extreme
case is ASASSN-13db, for which $\dot{M}_{\mathrm{acc}}$ decreases by two
orders of magnitude from the outburst peak in 2015 to quiescence in 2017.
Also, in NY Ori $L_{\mathrm{acc}}$ decreases by a factor 25 in one year.
In 80% of the sample we detect the [O I] 6300 Å line, a tracer of mass loss.
From the variability of the H$\alpha$/[O I] 6300 Å ratio, we conclude that
mass accretion variations are larger than mass loss variations.
From the analysis of the H I recombination lines a correlation is suggested
between the density of the line formation region, and the level of accretion
activity of the source.

Star formation (1565); Pre-main sequence objects (1290); Circumstellar disks
(235); Stellar mass loss (1613); Eruptive phenomena (475)

##  1 Introduction

EXor objects (named after the prototype EX Lupi, Herbig 1989) are Pre-Main
Sequence (PMS) stars showing episodes of eruptive accretion caused by
magnetospheric accretion events (Shu et al. 1994). Since the unsteady mass
accretion is a central theme in the star formation process (e.g. Cieza et al.
2018), these objects have been the subject of many investigations (Kuffmeier
et al. 2018; Contreras-Peña et al. 2019; Meng et al. 2019; MacFarlane et al.
2019a, 2019b; see also Audard et al. 2014 for a comprehensive view of the EXor
phenomenon). EXors are characterized by outbursts of short duration (typically
months) occurring at different time-scales (months, years) and showing
amplitudes of several magnitudes at optical and near-IR wavelengths. An
evolutionary scheme has been proposed in which EXors represent a less
energetic and (possibly) later stage than the powerful FU Orionis eruptions
(FUors, Hartmann & Kenyon 1985, Hartmann et al. 1993), although EXors present
substantial differences, such as shorter and more frequent outbursts, spectra
dominated by emission lines instead of absorption lines, and smaller values of
the mass accretion rate during outbursts (10${}^{-7}-10^{-6}$ M⊙ yr-1 vs.
10${}^{-5}-10^{-4}$ M⊙ yr-1). Yet, at present, a detailed model for the disk
structure and its evolution does not exist for EXors and different hypotheses
have been proposed for the possible trigger of their outbursts: thermal
instability in the disk (e.g. Audard et al. 2014 and references therein),
gravitational perturbations induced by a binary companion (Lodato $\&$ Clarke
2004) or by the migration of a giant planet (Bonnell & Bastien 1992), or
sudden changes in the stellar magnetic activity (D’Angelo & Spruit 2012,
Armitage 2016).

Observational constraints to theoretical models are nowadays provided by
photometric monitoring programs (e.g. Grankin et al. 2008; Morales-Calderón et
al. 2011; Guo et al. 2020) and by ongoing all-sky surveys (e.g.
Gaia111https://sci.esa.int/web/gaia, All-Sky Automated Survey for Supernovae
(ASAS-SN)222http://www.astronomy.ohio-state.edu/asassn/index.shtml,
Intermediate Palomar Transient Factory
(iPTF)333https://www.ptf.caltech.edu/iptf, Panoramic Survey Telescope & Rapid
Response System (Pan-STARSS)444https://panstarrs.stsci.edu/, and VISTA
Variables in the Via Lactea (VVV)555https://vvvsurvey.org/), which
sistematically follow (although often at a poor level of sensitivity) the
photometric variations of many EXors and EXor candidates, thus allowing us to
track the changes of their integrated properties (light-curves, Spectral
Energy Distributions -SEDs- and colors). A dramatic improvement in this
context will be provided in the next future by the the Vera Rubin
Observatory666https://www.lsst.org/, which will explore the southern sky with
unprecedented sensitivity and cadence.

Conversely, the optical and IR spectroscopic monitoring of the EXors is
currently not at a comparable level of sky coverage and cadence, despite its
crucial importance for understanding how the gas properties (e.g. excitation,
ionization, dynamics) vary during different activity phases. Indeed, so far
only few EXors have been spectroscopically investigated (e.g. EX Lupi: Sipos
et al. 2009; Sicilia-Aguilar et al. 2012; V2492 Cyg: Hillenbrand et al. 2013;
V1118 Ori: Audard et al. 2010; Lorenzetti et al. 2015; Giannini et al. 2016,
2017, 2020; HBC722: Kóspál et al. 2016). In addition, Guo et al. (2020)
recently presented a near-IR spectroscopic monitoring of a few VISTA variables
(supposedly eruptive candidates).

Waiting for instruments like
SoXS777https://www.eso.org/sci/facilities/develop/instruments/SoXS.html at
ESO-NTT (Schipani et al. 2018), which will be specifically dedicated to the
spectroscopic monitoring of large samples of objects, we have started in 2014
the EXOR OptiCal and Infrared Systematic Monitoring program EXORCISM
(Antoniucci et al. 2013). Over the years, new EXor candidates have been
discovered and included in our database. As a result, we provide here the
first flux-limited deep spectroscopic atlas of EXors (and candidates)
collected through state-of-the-art instrumentation (Large Binocular Telescope
- LBT) at optical and near-IR wavelengths. As outburst phases are expected to
be short-lived and rare events, most of our spectra have been taken during
quiescence phases. These spectra allow us to infer the baseline values of the
physical parameters, such as intrinsic line fluxes, accretion luminosity
($L_{\mathrm{acc}}$), and mass accretion rate ($\dot{M}_{\mathrm{acc}}$),
which can be used as reference in case of outburst. This is what we did for
the few sources that we observed both in quiescence and outburst, for which we
were able to study the variation of the parameters between the two phases.

The paper is organized as follows: Section 2 gives general information on the
investigated sample while Section 3 provides details on the observations and
data reduction procedures. In Section 4 we comment on the light-curves of the
individual sources. Sections 5 and 6 are dedicated to the description of the
spectra and to the discussion of the results obtained from the spectroscopic
analysis. Final remarks are given in Section 7.

##  2 The investigated sample

Table 1 lists our targets with their general properties. It provides
coordinates, distance, location, binarity, and jet or outflow detection.
Together with the name given in the Simbad Astronomical
Database888http://simbad.u-strasbg.fr/simbad/, we report other designation(s)
commonly used in the literature.

Our sample is divided into two groups. The first is composed of 5 known EXors
(XZ Tau, UZ Tau E, VY Tau, NY Ori, V1143 Ori) listed in the compilation by
Audard et al. (2014) and confirmed by many works in literature as members of
the class. In addition, we have also obtained multi-epoch spectroscopy of the
classical EXor V1118 Ori at both optical and IR wavelengths. These data are
not discussed here since they have been separately presented in a series of
dedicated papers (Giannini et al. 2020 and references therein).

The second group, generically named as ’EXor candidates’ in Table 1, consists
of 6 sources that have been reported through different alerts (e.g.
astronomers’ telegrams and public surveys alerts) for being young and highly
variables objects ($\Delta$V $\stackrel{{\scriptstyle>}}{{{}_{\sim}}}$2 mag).
Here we briefly comment on each of them. Two objects, namely V1647 Ori and PV
Cep, are present in the Audard et al. (2014) compilation, but their nature is
still controversial. V1647 Ori has been previously classified as a candidate
FUor, but it is currently considered as a peculiar object, since its spectrum
now differs significantly from that of a bona fide FUor (Connelley & Reipurth
2018), showing both FUor and EXor features. A few more objects with similar
characteristics have been identified during the VVV spectroscopic survey and
dubbed ’MNors’ (Contreras Peña et al. 2017). PV Cep and DR Tau were originally
included in the first list of EXor variables by Herbig (1989). PV Cep
repeatedly undergoes intense outbursts (up to 5 mag) of short duration and
brightness between that of FUors and EXors (Andreasyan et al. 2021). Its high
luminosity ($\approx$ 100 L⊙) is likely attributable to a young Herbig Ae star
(see e.g. Lorenzetti et al. 2011). DR Tau has progressively increased its
brightness by about three magnitudes between 1960 and 1980. Since then its
average brightness has been at the same level but still with remarkable
photometric fluctuations with amplitude up to $\sim$ 1 mag in the optical
(Banzatti et al. 2014 and references therein). Similarly, the $B$ magnitude of
V350 Cep changed from $>$ 21 to $\sim$ 17 between 1954 and 1978. Afterward,
the source has remained at roughly the same level of brightness, with the
exception of at least two fading events of about 2 mag in $B$ in 1979 and
2004. For these reasons Herbig (2008) did not classify V350 Cep as an EXor.
Considering its high and stable brightness level and its spectrum rich in
emission lines, it is likely that also V350 Cep is an object with properties
in between those of FUors and EXors (Andreasyan et al. 2021). iPTF15afq is a
new candidate EXor, suggested by Miller et al. (2015) on the basis of the
amplitude of its outburst of about 2.5 mag (in $R$-band) and its spectrum
dominated by lines in emission. ASASSN-13db (discovered by the ASAS-SN survey)
has been originally classified as an EXor by Holoien et al. (2014). Later
observations of blue-shifted absorption features in several hydrogen and
metallic lines suggest an intermediate behavior between EXors and FUors
(Sicilia-Aguilar et al. 2017).

In Table 2 we summarize the photometric properties of the sources, with the
range of variation observed in the indicated band and a short (non-exhaustive)
summary of the variability history as retrieved from the literature.

In Table 3 we list the relevant stellar parameters as given in the literature,
namely the visual extinction (AV), bolometric ($L_{\mathrm{bol}}$), stellar
($L_{\mathrm{*}}$) and accretion ($L_{\mathrm{acc}}$) luminosity, spectral
type (SpT) and effective temperature ($T_{\mathrm{eff}}$), stellar mass
($M_{\mathrm{*}}$), mass accretion rate ($\dot{M}_{\mathrm{acc}}$), and
evolutionary class (typically derived from the spectral slope between 2 and 24
$\mu$m). All sources but PV Cep, have mass
$\stackrel{{\scriptstyle<}}{{{}_{\sim}}}$1 M⊙ and spectral types M-K. V1647
Ori and PV Cep present a high and variable AV while all other sources are only
slightly extincted. Mass accretion rates have been measured for about 70% of
the sample. During quiescence, they are typically of the order of
10${}^{-8}-10^{-7}$ M⊙ yr-1, but show variations up to several orders of
magnitude in bursts.

In summary, our sample is composed of sources whose variability may differ in
amplitude, cadence, and duration. In all cases, however, the increase of
brightness has been attributed in the literature to accretion events,
sometimes accompanied by significant variations of the local extinction.

##  3 Observations and data reduction

The observations were carried out between 2014 and 2019 with the 8.4m Large
Binocular Telescope (LBT) located at Mount Graham (Arizona, USA). Optical and
near-IR spectra of the targets were obtained with the Multi-Object Double
Spectrograph (MODS, Pogge et al. 2010) and the LBT Utility Camera in the
Infrared (LUCI, Seifert et al. 2003), respectively.

Given the remarkable brightness of the sources, our project was executed as a
filler program during unfavorable atmospheric conditions that usually make
observations of weaker sources unfeasible. As a consequence, we were not able
to complete the survey with all the needed spectra. However, the poor
photometric quality of the sky (typical seeing
$\stackrel{{\scriptstyle>}}{{{}_{\sim}}}$1$\farcs$5) did not represent a
significant problem for the validity of our results, since the flux
calibration of the spectra was based on photometric observations taken close
in time to the spectra.

In total, we acquired 19 spectra with MODS (11 targets) and 10 spectra with
LUCI (8 targets), see the journal of observations in Table 4. Optical and
near-IR spectra have been taken during the same night only for V1647 Ori, and
for DR Tau at a temporal distance of 4 days. For all the other targets,
optical and near-IR spectra are spaced by several months or years, and, in
some cases, the spectrum of one of the two segments is missing.

MODS observations were performed with the dual grating mode (Blue + Red
channels, spectral range 350$-$950 nm) by using a 0$\farcs$80 slit ($\Re\sim$
1500 and 1800 in the Blue and Red channels, respectively). The slit angle
matched the parallactic angle to minimize the wavelength dependence on the
slit transmission. For each source, the adopted integration time is reported
in Table 4.

LUCI observations were carried out with the G200 low-resolution grating
coupled with the 0$\farcs$75 slit. Two datasets were acquired with the
standard ABB’A’ technique using the $zJ$ and $HK$ grisms. This provides a
final spectrum covering the 1.0$-$2.4 $\mu$m wavelength range at
$\Re\sim$1500\. The integration time for the individual sources is provided in
Table 4.

Data reduction was performed at the Italian LBT Spectroscopic Reduction
Center999http://www.iasf-milano.inaf.it/Research/lbt_rg.html, by means of
scripts optimized for LBT data. Data reduction steps of each MODS spectral-
image are: correction for dark and bias, bad-pixel mapping, flat-fielding, sky
background subtraction, and extraction of the one-dimensional spectrum by
integrating the stellar trace along the spatial direction. In the optical
range few telluric features are present, being the most prominent that around
7600 Å , around which, however, there are no lines used for the subsequent
accretion and ejection analysis. Therefore, we have not applied any telluric
correction to the optical spectra. Wavelength calibration was obtained from
arc lamps. Intercalibration between Blue and Red spectral segments was
verified superposing the spectral range between 5300 and 5900 Å in common
between the two channels. In all cases, the Blue and Red spectra resulted
optimally aligned without the need for further corrections.

The raw LUCI spectral images were flat-fielded, sky-subtracted, and corrected
for optical distortions in both the spatial and spectral directions. Telluric
absorptions were removed using the normalized spectrum of a telluric standard
star, after the removal of its intrinsic spectral features. Wavelength
calibration was obtained from arc lamps spectra. In a few cases, we applied an
inter-calibration correction to align the $zJ$ and $HK$ segments, which did
not exceeded 20% of the flux density.

The bad atmospheric conditions prevented us from flux calibrating the spectra
by using spectrophotometric standards. Therefore, we used as absolute
calibrators the acquisition images obtained immediately before the spectral
images in the filter specified in Table 4. Photometry was obtained by taking
as reference all the visible stars in the field, for which either a PAN-
STARRS101010panstarss.stsci.edu (in the optical) or a
2MASS111111ipach.caltech.edu/project/2mass (in the near-IR) photometry is
available. In few cases, namely when not enough reference sources were present
in the acquisition image or our target appeared saturated, we used the
photometry retrieved from public databases and obtained a few days before or
after the LBT observation (as specified in the note of Table 4).

##  4 Light-curves

To establish the level of accretion activity of each target at the time of the
LBT observation we collected data from public surveys (ZTF, ASAS-SN, and
AAVSO121212https://www.aavso.org/) to derive the optical light-curves of the
sources over several years, which are displayed in Figure 1. For NY Ori and
V1647 Ori, only the light-curves of the ASAS-SN survey are available.
Unfortunately, the ASAS-SN images of these sources are strongly contaminated
by the diffuse emission of the ONC and Mc Neil’s nebulae, where NY Ori and
V1647 Ori are respectively located (Aspin et al. 2009, Herbig 2008).
Therefore, these two light-curves are not shown.

In the following, we briefly comment on the light-curves of the remaining 9
objects.

  * -

XZ Tau. The $V$ band magnitude varied between $\sim$ 14.5 and 12 mag within a
6 yr period. Our MODS spectra cover both a high- and low-level brightness
state.

  * -

UZ Tau E. This source shows a very small amplitude level of variability within
5 yr. It has maintained an average brightness level around $V$ = 12.7 mag,
with fluctuations of $\sim$ 0.3 mag. The spectroscopic observations were
obtained at photometric levels with 0.6 mag of difference.

  * -

VY Tau. A remarkable brightening occurred at the beginning of 2014,
unfortunately not covered by our observations. This was followed by a second
low-amplitude burst about a year later. The first MODS spectrum refers to the
declining phase of this burst.

  * -

V1143 Ori. Between 2013 and 2018 the source has remained at $V$ $\sim$ 17 mag
with typical fluctuations of tenths of magnitude. The MODS spectrum was
acquired at V$\sim$ 16.5 mag.

  * -

DR Tau. In the last seven years the DR Tau average photometry has been $V\sim$
12 mag, with large fluctuations up to 1 mag. All our observations have been
performed close to the average level.

  * -

ASASSN-13db. This source underwent a first, short-lived outburst in 2013,
followed by a second more intense outburst from 2014 to late 2016. We have
observed ASASSN-13db four times, the first three (two with MODS and one with
LUCI) during the second burst, and the last with MODS in 2017, when the source
was back to quiescence.

  * -

iPTF15afq. The source underwent an outburst at the beginning of 2018 (not
shown in the $r$ band light-curve, see Hillenbrand 2019), followed by a
quiescence period of about 2 years. A second outburst occurred in 2019,
covered by our LUCI observation during the rising phase and close to the peak.
A slow decay started at the beginninig of 2020 and is still ongoing. The
remaining three observations (between 2016 and 2018), have been taken during
quiescence ($R\sim$ 18.5-19 mag), according to the light-curve presented by
Miller et al. (2015).

  * -

PV Cep. Between 2013 and 2020, PV Cep has experienced a slow increase in
brightness, with a peak in 2016 ($V\sim$ 14.5 mag), followed by a decrease
down to $V>$ 18 mag at a roughly similar speed. Both our observations were
conducted close to the brightness peak.

  * -

V350 Cep. Our first LUCI observation was taken in November 2015, at $R$=15.4
mag (Semkov et al. 2017). The source experienced a deep dimming in spring 2016
followed by a new brightening. Since 2018, V350 Cep is again progressively and
slowly fading (from $r\sim$ 15.3 to $r\sim$ 15.8). Our second LUCI observation
(2018) was obtained when V350 Cep was still relatively bright ($r\sim$ 15.5
mag), while the next three MODS spectra were taken when $r$ had decreased down
to 15.6-15.8 mag.

Figure 1: Optical ($V-$ or $r-$band) light-curves retrieved from data of
public surveys (ZTF, ASAS-SN, and AAVSO). The x-axis represents the Modified
Julian Date (MJD), while some calendar dates are given on top as reference. A
blue (red) vertical line indicates the date of a MODS (LUCI) observation. The
black line corresponds to quasi-simultaneous MODS+LUCI observation of DR Tau.
Photometry taken closest in time to the spectroscopic observation is marked
with a green dot.

Summarizing, our targets have shown a variety of behaviors during the last
decade. Some of them have remained fairly constant, others have undergone
strong outbursts or long-term and remarkable brightness variations. As
expected, in most cases our observations were performed during phases of
quiescence. However, some exceptions exist: indeed we have observed XZ Tau and
PV Cep during a high level of their activity and even caught the outbursts of
ASASSN-13db and iPTF15afq. In addition, our acquisition images of NY Ori
indicate a variation of about two magnitudes in the $r$-band between 2014
($r\sim$ 12) and 2016 ($r\sim$ 14).

##  5 Description of the spectra

Figure 2: MODS spectra. Spectra obtained in different dates are shown with
different colors. Brightest lines and telluric features are labelled. Figure
3: As in Figure 2.. Figure 4: LUCI spectra. Spectra obtained in different
dates are shown with different colors. Brightest lines are labelled.  Figure
5: As in Figure 4.

Figures 2-3 and Figures 4-5 show the MODS and LUCI spectra of our sample,
respectively. Line fluxes have been derived by fitting the profiles with a
Gaussian function, and the relative uncertainties have been estimated by
multiplying the rms noise of the continuum close to the line times the FWHM.
Appendix A (Tables 8 \- 18) report the fluxes of lines used in the analysis.

Lines tracing accretion have been detected in all sources. In particular, the
Balmer recombination lines (typically from H$\alpha$ to H$\delta$, but for
many sources up to H15) are detected in all MODS spectra. In several objects
(NY Ori, DR Tau, ASASSN-13db, V1647 Ori, PV Cep) some of these lines are
partially or totally seen in absorption. Note that, given the late spectral
type of the sources (Table 3), the photospheric contribution to the lines seen
in absorption is negligible, with the possible exceptions of NY Ori and PV
Cep. Paschen and Brackett lines (both in the optical and in near-IR range) are
detected in 10 and 5 objects, respectively, and appear all in emission. In all
objects observed with LUCI, we detect the He I 1.08 $\mu$m line, while other
He I optical lines are seen in the spectra of the brightest objects. Metallic
lines of many species (Fe I, Fe II, Ti I, Ti II, Ca II, Cr I, Cr II) are also
detected.
Signatures of ejection activity, such as the [O I] 6300 Å line, are present in
all sources but VY Tau and V1143 Ori. In two objects (iPTF15afq and V350 Cep)
we also detected H2 emission at 2.12 $\mu$m, which is a typical tracer of
molecular outflows. In XZ Tau, UZ Tau E, DR Tau, and PV Cep, we detect some [S
II] lines, indicative of gas with a low ionization degree. Lines of [S II],
[Fe II], [Ni II], and [N II] have been detected in the spectrum of iPTF15afq,
which could signal the occurrence of a high-velocity jet (Giannini et al.
2019).

In some of the LUCI spectra we have identified at least one of the two CO ro-
vibrational bandheads (v=2-0 and v=3-1) in emission (UZ Tau E, iPTF15afq, and
PV Cep), or in absorption (VY Tau), as expected for EXor sources (e.g. Biscaya
et al. 1997, Hillenbrand et al. 2013). In the remaining objects the CO
bandheads are undetected, probably because the edge of the HK spectra is
typically very noisy.

In general, we remark that: 1) the spectra well resemble the optical/near-IR
spectra of accreting T Tauri stars (e.g. Alcalá et al. 2017) and no evident
differences exist among the spectra of known EXors and EXor candidates; 2) the
intensity of the accretion tracers (e.g. H I , Ca II) follows the variation of
the source brightness (e.g. XZ Tau, Table 8) and may even disappear or change
from emission to absorption in case of an outburst event (ASASSN-13db and
iPTF15afq).

##  6 Analysis and discussion

###  6.1 Extinction

The primary goal of the present work is to determine the accretion parameters
($L_{\mathrm{acc}}$ and $\dot{M}_{\mathrm{acc}}$ ) of the sample. To get a
meaningful estimate of these two quantities we first derived the visual
extinction (AV) toward the sources. As listed in Table 3, one or more AV
values are reported in the literature for some of them. Although these values
represent an important reference, several studies have shown that a remarkable
reduction (e.g. Hillenbrand et al. 2013), or less often, an increase of the
extinction (in edge-on disks, Stock et al. 2020) often accompany accretion
burst episodes. Therefore, we cannot assume AV as a constant, but we need to
estimate its value for each observation date. In the following, we separately
discuss the methods used to derive AV from both the optical and near-IR
spectra.

####  6.1.1 Extinction derived from optical spectra

The method used to derive the extinction from the optical spectra is described
in detail in Giannini et al. (2018). This is essentially based on the
empirical relationships found by Alcalá et al. (2014, 2017), between the
accretion luminosity, $L_{\mathrm{acc}}$, and the luminosities
$L_{\mathrm{i}}$ of selected emission lines (so-called ’line method’). In the
optical range these relationships exist for more than 20 lines, namely the H I
recombination lines of the Balmer and Paschen series from H$\alpha$ to H15,
and from Pa8 to Pa10, along with He I , O I, and Ca II lines. Therefore, the
relationships provide up to around 20 independent estimates of
$L_{\mathrm{acc}}$ , namely one for each observed line ($L_{\mathrm{acc(i)}}$
).

In our procedure, we fit simultaneously $L_{\mathrm{acc}}$ and AV , this
latter allowed to vary between 0 and 15 mag (in steps of 0.20 mag). The
extinction law by Cardelli et al. (1989) and total-to-selective extinction
ratio RV = 3.1 are assumed. We iteratively fit $L_{\mathrm{acc}}$ to minimize
the dispersion among the individual $L_{\mathrm{acc(i)}}$ , computed by de-
reddening the observed fluxes for the current AV and assuming the distance
listed in Table 1. The best estimate of AV is the value for which the
dispersion between the $L_{\mathrm{acc(i)}}$ is minimized and we assume as
$L_{\mathrm{acc}}$ the average of the $L_{\mathrm{acc(i)}}$ corresponding to
that AV.

Figure 6: Accretion luminosity derived from the flux of the indicated tracers
by using the relationships by Alcalá et al. (2017). The horizontal line shows
the average $L_{\mathrm{acc}}$ . The source name, the observation date, and
the fitted values of $L_{\mathrm{acc}}$ , and AV are indicated.

Figure 7: As in Fig. 6.

Figure 8: As in Fig. 6.

We included in the fit only the lines that do not present any absorption
component. We have verified that, rather than on the individual flux errors,
the main limitation of this method comes from the uncertainties on the
empirical relationships (Alcalá et al. 2017), which give an error on AV of
0.5$-$1.5 mag, depending on how many lines are considered in the fit. Further
causes of uncertainty come from the assumptions on the extinction law and RV.
For example, we have checked that if the extinction law by Weingartner &
Draine (2001) is applied, the fitted AV decreases by $\approx$ 0.4 mag.
Conversely, an increase in AV by $\sim$ 0.2-0.3 mag is produced for RV = 5.5.
In Figures 6-8 we show the results of our procedure, while in Table 5 we
report the fitted extinction values and the number of lines used.

An independent way to compute AV has been described by Alcalá et al. (2021,
so-called ’continuum’ method). The observed spectrum is iteratively divided by
a template of the same spectral type (SpT), artificially reddened by varying
the value of AV. The best estimate of AV is that for which the ratio has a
flat slope. Applying this method to strong accretors, one has to correct the
spectrum for the excess continuum (veiling) induced by the accretion hot spot
(UV) and disk (IR) emission. This has been estimated for the majority of our
sources by applying the procedure described in Appendix B.
Following the prescriptions by Alcalá et al. (2014) and Fischer et al. (2011),
we applied the continuum method in the spectral range between 550 and 800 nm,
where the veiling excess is minimized and expected not to change more than the
typical error of 0.3 associated with its determination (Appendix B). We
adopted as template spectra a grid of non-accreting (Class III) YSOs of the
same SpT of the observed sources (Manara et al. 2013, 2017), reddened for AV
between 0 and 15 mag in steps of 0.20 mag. The results of the fitting
procedure are provided in Figure 9 and the fitted AV are listed in Table 5.
Noticeably, in 6 out of the 9 optical spectra examined, the two AV
determinations agree within 0.5 mag.

Figure 9: Ratios between observed spectra and artificially reddened and
veiling-corrected templates. Ratios with minimum slopes (for which the
extinction AV is estimated) are indicated in blue. For each panel, object
name, date of observation, spectral type, and veiling are also reported.

Finally, the case of V1647 Ori deserves a short discussion. This source is a
well-known Class I/flat young source, still deeply embedded in the parental
dusty envelope. As shown in several works (e.g. Nisini et al. 2016, Fiorellino
et al. 2021), in embedded sources a relevant fraction of the optical emission
originating close to the star is scattered in the cavity excavated inside the
envelope. The effect is to enhance the emission at shorter wavelengths, that
in turn simulates a lower extinction. For this reason, we consider not
reliable the determination of AV (9.2 mag) derived in the optical range and
will assume the extinction derived from the LUCI spectrum on the same date (13
mag, Sect. 6.1.2).

####  6.1.2 Extinction derived from near-IR spectra

Figure 10: [$J-H$] vs. [$H-K$] color-color diagram. The arrow shows the
extinction vector for AV =5 mag, according to the reddening law by Cardelli et
al. (1989). The dashed line marks the locus of the unreddened T Tauri stars
(Meyer et al. 1997), while the continuous line represents Main Sequence stars
(Luhman et al. 2010). Data points of the objects discussed in the text are
labelled.

The AV estimate from the near-IR spectra was also obtained by applying two
independent methods, briefly described in the following paragraph.

The first method is based on the [$J$-$H$] and [$H$-$K$] colors at the date of
the observation. One photometric point is directly provided by the acquisition
image (Table 4), while the other two are computed from the flux-calibrated
spectrum convolved with the Johnson filter profiles131313retrived at
http://svo2.cab.inta-csic.es/theory/fps/, Spanish Virtual Observatory (SVO)
Filter Profile Service (Rodrigo & Solano 2020). We estimate that this
procedure leads to an uncertainty of about $\pm$0.1 mag in each band. From the
magnitudes we have constructed the [$J$-$H$] vs. [$H$-$K$] color-color diagram
shown in Figure 10. Here, we indicate with an arrow the direction of the
extinction vector, whose length corresponds to an extinction of 5 mag. We
assume the reddening law by Cardelli et al. (1989) and RV = 3.1. We estimate
AV from the position of the observed colors and the unreddened Classical T
Tauri locus (CTTs, Meyer et al. 1997) and list the values in Table 5.
Considering all the uncertainties involved in this procedure (absolute
spectral calibration, propagation of magnitude errors, assumptions on
extinction law and RV), we conservatively evaluate a final error on AV of 1
mag.

All sources have low or negligible AV, with the exception of V1647 Ori (AV
$\sim$ 13 mag). In the case of VY Tau, which is located at the left edge of
Figure 10, we are not able to estimate a reliable value of AV, while for two
observations (iPTF15afq and V350 Cep) we cannot apply the method because the
available spectra do not cover the three bands.

The second method relies on the detection of accretion line tracers, as done
for the optical spectra, see Figures 6-8. However, empirical relationships in
the near-IR are available only for the Paschen lines (from Pa8 to P$\beta$)
and for the Br$\gamma$ (Alcalá et al. 2017), which are moreover less sensitive
to extinction variations than the optical lines. For DR Tau, for which we have
quasi-simultaneous MODS and LUCI spectra, the optical and near-IR lines have
been considered all together in the fit. Conversely, as explained in the
previous section, the lines of V1647 Ori observed with LUCI have been fitted
separately from those observed with MODS, although the two spectra were
obtained on the same date. The fitted AV of 13 mag is in perfect agreement
with that derived from the color-color diagram.

We were able to apply both methods to five spectra. The two AV determinations
are in excellent agreement in UZ Tau E and V1647 Ori (Table 5), but differ
between 0.9 and 3.3 mag in the other three sources (DR Tau, PV Cep and V350
Cep). Since we have no reason to prefer one method over the other (also
considering that the lines used in the fit of $L_{\mathrm{acc}}$ are few), we
have assumed the entire range of AV determinations to compute the accretion
parameters $L_{\mathrm{acc}}$ and $\dot{M}_{\mathrm{acc}}$ (see Sect. 6.2).

###  6.2 Accretion luminosity and mass accretion rate

As explained in Sect. 6.1.1, the accretion luminosity, $L_{\mathrm{acc}}$ ,
has been derived together with AV by fitting the individual
$L_{\mathrm{acc(i)}}$ . With reference to the fits shown in Figures 6-8 we
note that, in general, the $L_{\mathrm{acc(i)}}$ derived from the individual
lines agree within the errors. There are however lines whose accretion
luminosity may be up to an order of magnitude higher than the average
$L_{\mathrm{acc}}$ . In particular, we signal Ca II K and He I at 492 and 502
nm. However, thanks to the high number of lines used, $L_{\mathrm{acc}}$ does
not change by more than 10% whether or not these three lines are included in
the fit.
For many sources, a further estimate of $L_{\mathrm{acc}}$ has been obtained
by applying the relationships between $L_{\mathrm{acc}}$ and $L_{\mathrm{i}}$
by fixing the AV to the value derived by the continuum method (in the optical)
or by the color-color diagram (in the infrared), see Table 6. This way, we
have checked that $L_{\mathrm{acc}}$ does not change by more than 30%.

The mass accretion rate, $\dot{M}_{\mathrm{acc}}$ , has been estimated using
the relationship by Gullbring et al. (1998):

$\dot{M}_{\mathrm{acc}}\approx
1.25{\frac{L_{\mathrm{acc}}R_{\mathrm{*}}}{GM_{\mathrm{*}}}}$ (1)

where $R_{\mathrm{*}}$ and $M_{\mathrm{*}}$ are the stellar mass and radius,
and G is the gravitational constant. We have adopted for $M_{\mathrm{*}}$ the
values given in the literature (Table 3) and derived $R_{\mathrm{*}}$ from the
relationship:

$R_{\mathrm{*}}={\frac{1}{2T_{\mathrm{eff}}^{2}}}\sqrt{{\frac{L_{\mathrm{*}}}{\pi\sigma}}}$
(2)

where $\sigma$ is the Stefan-Boltzmann constant, and $T_{\mathrm{eff}}$ and
$L_{\mathrm{*}}$ are taken from Table 3.

The derived $\dot{M}_{\mathrm{acc}}$ values are listed in Table 6.

It is worth noting that no remarkable differences exist between EXors and EXor
candidates, either in $L_{\mathrm{acc}}$ or in $\dot{M}_{\mathrm{acc}}$ .
Typical values are $-$2 $\stackrel{{\scriptstyle<}}{{{}_{\sim}}}$
Log($L_{\mathrm{acc}}$/L⊙) $\stackrel{{\scriptstyle<}}{{{}_{\sim}}}$ 1, and
$-$9 $\stackrel{{\scriptstyle<}}{{{}_{\sim}}}$ Log($\dot{M}_{\mathrm{acc}}$/M⊙
yr-1) $\stackrel{{\scriptstyle<}}{{{}_{\sim}}}$ $-$7\. V1647 Ori is the object
with significantly higher $L_{\mathrm{acc}}$ ($\sim$ 50 L⊙ ) and
$\dot{M}_{\mathrm{acc}}$ ($\sim$ 2 10-5 M⊙ yr-1 ), in line with its very young
age and its classification as a Class I/flat source. This source underwent two
outbursts during its recent history, one in 2003-2004 (Aspin et al. 2006) and
another one in 2008-2011 (Aspin 2011). During the second outburst, its
$r$-band photometry was $\sim$ 17.77 and AV = 8 mag. In 2016, we measured
$r$=18.68, namely about 1 mag fainter. Despite that, our determination of
$\dot{M}_{\mathrm{acc}}$ is about one order of magnitude higher than the
estimate of (4$\pm$2) 10-6M⊙ yr-1 measured by Aspin (2011). This disagreement
is due to the higher AV (13 mag) that we have estimated from the LUCI
spectrum. A considerable variability in AV is indeed exhibited by V1647 Ori.
During quiescence ($r\sim$ 23), AV was around 19 mag (Aspin 2011), therefore
it is plausible to assume that in an intermediate stage, as that we have
probed, also the extinction was in between the two extreme values of 8 and 19
mag. It would therefore be relevant to measure $\dot{M}_{\mathrm{acc}}$ during
the quiescent state (namely in the conditions of maximum AV ) to disentangle
the contributions of accretion and extinction in the large-amplitude
variability exhibited by this source.

The sources for which we have more than one observation are XZ Tau, UZ Tau E,
NY Ori, DR Tau, ASASSN-13db, iPTF15afq, PV Cep, and V350 Cep. Of these, all
but NY Ori and ASASSN-13db present an amplitude variability
$\stackrel{{\scriptstyle<}}{{{}_{\sim}}}$ 1 mag (Figure 1), in line with that
observed in Classical T Tauri stars (Costigan et al. 2014). These magnitude
fluctuations correspond to variations in $L_{\mathrm{acc}}$ and
$\dot{M}_{\mathrm{acc}}$ of a factor $\approx$ 1.5$-$3\. In general, these
results confirm the accepted scenario in which accretion is a non-stationary
phenomenon, although they are insufficient to explain the spread of more than
two orders of magnitude found in the $\dot{M}_{\mathrm{acc}}$ determinations
among sources of the same mass (e.g. Antoniucci et al. 2014).

However, larger magnitude fluctuations imply significant variations in
$L_{\mathrm{acc}}$ and $\dot{M}_{\mathrm{acc}}$. In ASASSN-13b, in the 2 years
between the outburst peak and the quiescence, $L_{\mathrm{acc}}$ changed from
0.2 to 0.03 L⊙ and $\dot{M}_{\mathrm{acc}}$ from 3 10-8 to 5 10-10 M⊙ yr-1 .
Also, in NY Ori $L_{\mathrm{acc}}$ decreased by a factor 25 in about one year
(2014 December - 2016 January). Although no photometric measurements are
available for the indicated period, this change in $L_{\mathrm{acc}}$ is
compatible with the variation of around 3 mag observed several times in this
source (see Table 2).

###  6.3 Physical conditions of the H I lines emitting region

Figure 11: Balmer decrement (Hi/H$\beta$ flux ratios vs. upper level of the
transition) of XZ Tau in the two dates of observation (black: 2014/12/21, red:
2018/03/21). The hydrogen density fitted for the two sets of data is reported
in logarithmic scale. Different line symbols represent different temperatures
as indicated in the top-right corner of the figure.

High-resolution spectroscopic observations indicate that both the high-
velocity stellar winds and the gas channeled from the disk onto the stellar
surface contribute to the hydrogen line profiles (e.g. Folha & Emerson 2001,
Antoniucci et al. 2017a, Moura et al. 2020). Although our low-resolution
observations do not allow to kinematically separate these two gas components,
we can evaluate the average physical conditions of the hydrogen lines emitting
regions using the observed flux ratios as temperature and density diagnostics.

Using their local line-excitation model, Kwan & Fischer (2011) provided
predictions of the Balmer decrements (ratios from H15 to H$\alpha$ with
respect to H$\beta$), in the range of hydrogen density 8 $\leq$
log[$n_{\mathrm{H}}$(cm-3)] $\leq\,$12, and temperature between 3750 K and
15000 K. Similarly, Paschen decrements (ratios from Pa$\gamma$ to Pa12 with
respect to Pa$\beta$), together with the ratio Br10/Br$\gamma$ are given by
Edwards et al. (2013) for 8 $\leq$ log[$n_{\mathrm{H}}$(cm-3)] $\leq\,$12.4
and T = 5000$-$20000 K.

Fits through the intrinsic hydrogen line fluxes have been done by least
squares method for sources where at least 4 lines are detected. As an example,
we show in Figure 11 the Balmer decrement measured in XZ Tau for the two
observation dates (2014 and 2018). We plotted the Hi/H$\beta$ ratios up to
H14/H$\beta$ and H12/H$\beta$ for the 2014 (black) and 2018 (red)
observations, respectively. First, we note the different shape of the Balmer
decrement in the two dates. Following the nomenclature of Antoniucci et al.
2017b, in 2014 the decrement shape was of ’type 4’, typical of strong
accretors, and changed in 2018 to the ’type 2’ shape, which is the one most
frequently seen in T Tauri stars. This shape can be explained by variation of
a factor $\sim$ 2.5 in $n_{\mathrm{H}}$, with ratios Hi/H$\beta$ increasing
with nup in models with higher densities. Conversely, no clear correlation is
found between the temperature and the shape of the Balmer decrement, and in
particular, all the temperatures considered in the Kwan & Fischer model are
compatible for the observation in 2018.

Table 7 summarizes the results for the objects of our sample, also reporting
the brightness level of each source as derived from the photometric data and
the number of lines involved in the fit. As pointed out for the case of XZ
Tau, the temperature is not well constrained, apart from the case of UZ Tau E.
It is therefore more interesting to focus on the results regarding the
hydrogen density. In our previous study on the 2015 outburst of the classical
EXor V1118 Ori, we highlighted the sudden increase in the hydrogen density of
about 2 orders of magnitude, from log($n_{\mathrm{H}}$)=9.4 to
log($n_{\mathrm{H}}$)=11.4, passing from the quiescent phase to the peak of
the outburst (Giannini et al. 2017). In the objects of the sample presented
here, we fit the hydrogen density in the range log($n_{\mathrm{H}}$) from 9.0
to 11.0, considering the uncertainties. Indeed, a certain level of correlation
may be found between $n_{\mathrm{H}}$ and the activity level of the source. In
Sect. 4 and 6.2, we identified XZ Tau, PV Cep, iPTF15afq, ASASSN-13db, and NY
Ori as sources that showed considerable photometric variability during the
last years, often accompained by a significant $L_{\mathrm{acc}}$ and
$\dot{M}_{\mathrm{acc}}$ variations. Specifically, in XZ Tau the $V$ band
magnitude increased from $V\sim$ 12.8 in 2014 to $V\sim$ 14.2 in 2018. Our
analysis shows that this resulted in a decrease of veiling from 0.8 to 0.0, of
$\dot{M}_{\mathrm{acc}}$ by a factor of 5 and of $n_{\mathrm{H}}$ by a factor
around 2.5-3.0. PV Cep has been observed only during a period of high
brightness. Close to the peak of the light-curve in 2015 we fit
log($n_{\mathrm{H}}$)=11.0. In ASASSN-13db and in iPTF15afq, we are able to
measure $n_{\mathrm{H}}$ only during the quiescent phase. It is interesting to
note, however, that during periods of high brightness level (2014-2015 for
ASASSN-13db and 2019 for iPTF15afq), the H I lines are always seen in
absorption, a signature typical of FUors or in general of sources with a high
accretion rate (Connelley & Reipurth 2018). In NY Ori the situation is more
controversial. From our acquisition images we find differences of about 2
magnitudes in the $r$-band between 2014 ($r\sim$ 12) and 2016 ($r\sim$ 14).
All H I lines (from the Balmer and Paschen series) are seen in absorption in
2016, while the Paschen lines are in emission in 2014. Given the spectral type
of NY Ori (K4), a possibility is that during quiescence the observed H I lines
are in absorption because they originate in the photosphere, while during
outburst the accretion component in emission dominates.

In the remaining sources, which are all observed during quiescence,
log($n_{\mathrm{H}}$) is typically between 9.0 and 10.0. Only in V350 Cep we
estimate log ($n_{\mathrm{H}}$) between 10.6 and 10.8 in 2015. Unfortunatly,
the photometric level at that date is unknown.

###  6.4 Mass loss variability and forbidden line emission

Matter accretion from a disk necessarily imposes the removal of angular
momentum through winds and jets. This makes the phenomena of mass accretion
and ejection intimately correlated, so it is interesting to investigate
whether they are also connected in terms of variability.

Mass loss rate determinations suffer from multiple assumptions on the geometry
and physical conditions of the outflowing gas (e.g. Sperling et al. 2020).
Therefore, rather than searching for correlations between the rates of
accretion and ejection, we have preferred to look for flux variations of lines
representative of the two phenomena.

The [O I] 6300 Å is the most prominent outflow emission feature. Indeed, it
was detected in all our objects but VY Tau and V1143 Ori, namely the two
sources with the lowest values of $\dot{M}_{\mathrm{acc}}$ .

As tracer of accretion we have selected the H$\alpha$ line, both because of
its brightness and, more importantly, because the ratio H$\alpha$/[O I] 6300 Å
, being independent of extinction, is directly related to the intrinsic ratio
between mass accretion and mass ejection rates.

Among the 5 sources showing a significant $\dot{M}_{\mathrm{acc}}$ variation,
we have multi-epoch determinations of the [O I] 6300 Å for XZ Tau, NY Ori and
iPTF15afq. With reference to the line fluxes reported in Appendix A, we find
that in XZ Tau F([O I] 6300 Å ) has remained fairly constant while
F(H$\alpha$) decreased by a factor $\sim$ 2.4 from 2014 to 2018. This
indicates that between 2014 and 2018 the variation of the mass accretion has
been significantly larger than that of the mass loss. Similarly, in NY Ori,
F(H$\alpha$) and F([O I] 6300 Å ) decreased by a factor 7 and 3 in two years
(2014$-$2016). In iPTF15afq we observe the same trend although the variation
is low for both lines, due to the short temporal distance of just one year.

More striking is the trend in the classical EXor V1118 Ori. Indeed, comparing
the H$\alpha$ and [O I] 6300 Å fluxes during the 2015 outburst (Giannini et
al. 2016, 2017), with those of quiescence (Lorenzetti et al. 2015) we find a
variation of F(H$\alpha$) about ten times larger than the variation of F([O I]
6300 Å ).

##  7 Conclusions

We analysed optical and near-IR low resolution spectra taken at LBT between
2014 and 2019 of a sample of 11 variable young sources, composed of 5 well
known eruptive variables (EXors) and 6 Pre-Main Sequence objects showing
episodic variability likely attributable to intermittent accretion events
(EXor candidates). Ten sources have been observed more than once, therefore
allowing us to investigate for correlations between photometric and
spectroscopic variability. At this stage, the collected observations represent
the first flux-limited spectroscopic survey of candidate eruptive variables
able to accurately determining fundamental quantities such as visual
extinction; accretion parameters; physical parameters of the H I lines
emitting region, and signatures of mass outflowing. The main results of this
survey can be summarized as follows:

  * -

The analysis of the light-curves indicates that six sources (UZ Tau E, VY Tau,
V1143 Ori, DR Tau, V1647 Ori, and V350 Cep) have been observed while in
quiescence and other two, XZ Tau, and PV Cep, during a high level of
brightness. Remarkably, ASASSN-13db and ipTF15afq have been caught during
outburst.

  * -

All targets present lines that are tracers of accretion, like H I
recombination lines, both in the optical and IR range. All IR spectra show
prominent Ca II and He I 1.08 $\mu$m lines. Metallic lines of many species are
also detected. Signatures of ejection activity ([O I] 6300 Å , [S II] , H2
2.12$\mu$m) characterize about 80% of the sources. No evident difference
exists between the spectra of known and candidate EXors, all of them
resembling spectra of accreting young T Tauri stars.

  * -

For many sources the veiling excess was computed. We found that a significant
variation of the veiling is associated with accretion events. During the 2015
outburst of ASASSN-13db the veiling at 710 nm increased from 0.1 to $>$ 3.

  * -

Since the visual extinction AV closely depends on the source activity level at
the time of the observation, a major effort has been made to accurately
evaluate this parameter. For both optical and near-IR spectra two independent
methods (based on lines, continuum, and colors) were used to derive AV values.

  * -

Empirical relationships between the accretion luminosity ($L_{\mathrm{acc}}$ )
and the luminosity of selected lines allowed us to determine the average
accretion luminosity for each source and date of observation. The mass
accretion rate was also evaluated. We did not find any remarkable difference
between the $L_{\mathrm{acc}}$ and $\dot{M}_{\mathrm{acc}}$ values of known
EXors and EXor candidates of our sample, which are in the range $-$2
$\stackrel{{\scriptstyle<}}{{{}_{\sim}}}$ Log($L_{\mathrm{acc}}$/L⊙)
$\stackrel{{\scriptstyle<}}{{{}_{\sim}}}$ 1 and $-$9
$\stackrel{{\scriptstyle<}}{{{}_{\sim}}}$ Log($\dot{M}_{\mathrm{acc}}$/M⊙
yr-1) $\stackrel{{\scriptstyle<}}{{{}_{\sim}}}$ $-$7, respectively.

  * -

All sources observed more than once present significant
$\dot{M}_{\mathrm{acc}}$ variability even if they have been observed during
quiescence. In ASASSN-13db a decrease of $\dot{M}_{\mathrm{acc}}$ of two
orders of magnitude is observed from the outburst peak to the quiescent phase.
Also, the accretion luminosity of NY Ori decreased by a factor 25 in one year.

  * -

Physical parameters ($n_{\mathrm{H}}$ and T) of the H I emitting region have
been evaluated from the observed hydrogen line ratios by using predictions
provided in the literature. These ratios are much more effective in
constraining the density ($n_{\mathrm{H}}$ ) than the temperature, which
remains poorly defined. Generally, a direct correlation is recognizable
between density and accretion activity of the source, with significant density
increases (up to two orders of magnitudes) associated with large brightness
fluctuations.

  * -

Tracers of mass outflows, such as the [O I] 6300 Å line, have been detected in
most of the investigated sources. By comparing the [O I] variations with those
of a mass accretion tracer (H$\alpha$), we conclude that mass accretion
variations are larger than mass loss variations.

In the next future, other facilities appropriate for this research will be
available in the southern hemisphere, which are expected both to discover
large numbers of eruptive variables (e.g. VRO-LSST), and to spectroscopically
follow-up them in the optical/near-IR band (e.g. SoXS at ESO-NTT). Their
observations will probably deeply modify our knowledge of the eruptive
accretion phenomenon. In such framework we hope that our coherent database
will represent a valuable reference study.

##  8 Acknowledgements

This work is based on observations made with the Large Binocular Telescope
(LBT). The LBT is an international collaboration among institutions in the
United States, Italy and Germany. LBT Corporation partners are: The University
of Arizona on behalf of the Arizona university system; Istituto Nazionale di
Astrofisica, Italy; LBT Beteiligungsgesellschaft, Germany, representing the
Max-Planck Society, the Astrophysical Institute Potsdam, and Heidelberg
University; The Ohio State University, and The Research Corporation, on behalf
of The University of Notre Dame, University of Minnesota and University of
Virginia. This work has been supported by the project PRIN-INAF-MAIN-STREAM
2017 “Protoplanetary disks seen through the eyes of new-generation
instruments” and by the European Union’s Horizon 2020 research and innovation
programme un-der the Marie Sklodowska-Curie grant agreement No 823823 (DUST-
BUSTERS). We also acknowledge the support by INAF/Frontiera through the
‘Progetti Premiali’ funding scheme of the Italian Ministry of Education,
University, and Research and by the Deutsche Forschungs-Gemeinschaft (DFG,
German Research Foundation) - Refno. FOR 26341/1 TE 1024/1-1. This research
also received financial support from the project PRIN-INAF 2019
”Spectroscopically Tracing the Disk Dispersal Evolution”. ACG receveid
fundings from the European Research Council (ERC) under the European Union’s
Horizon 2020 research and innovation programme (grant agreement No. 743029).
A. Rossi acknowledges support from the INAF project Premiale Supporto Arizona
& Italia.

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Table 1:  The investigated sample

Simbad | RA (J2000) | DEC (J2000) | D1 | Location | Binarity | Jet/Outflow | Other | Ref
---|---|---|---|---|---|---|---|---
Name | (h:m:s) | (∘:′:′′) | (pc) |  |  |  | designation |
EXors
XZ Tau | 04:31:40.07 |  $+$18:13:57.2 | 150a | L1551 | Y(0$\farcs$30) | HH152/Y | Haro6-15$-$SVS169 | 1
UZ Tau E | 04:32:43.07 |  $+$25:52:31.0 | 130.4 | B19 | Y(0$\farcs$34) |  $\cdots$ / $\cdots$ | UZ Tau A | 2
VY Tau | 04:39:17.41 |  $+$22:47:53.4 | 153.4 | L1536 | Y(0$\farcs$66) | Y(?)/ $\cdots$ | SVS38 | 3,4
NY Ori | 05:35:36.01 |  $-$05:12:25.3 | 403.5 | ONC |  $\cdots$ |  $\cdots$ / $\cdots$ | Parenago 2119 | 5
V1143 Ori | 05:38:03.90 |  $-$04:16:42.8 | 395.3 | L1640 |  $\cdots$ |  $\cdots$ / $\cdots$ | Sugano’s obj | 5
EXor candidates
DR Tau | 04:47:06.21 |  $+$16:58:42.8 | 193.0 | L1558 | N | Y/ $\cdots$ |  $\cdots$ | 6
ASASSN-13db | 05:10:11.08 |  $-$03:28:26.3 | 387.1 | near L1615/16 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
V1647 Ori | 05:46:13.14 |  $-$00:06:04.9 | 412.9 | Mc Neil’s nebula | N(?) | HH23(?)/Y | McNeil’s Nebula | 7,8,9
iPTF 15afq | 07:09:21.40 |  $-$10:29:34.4 | 1315.7 | CMa |  $\cdots$ |  $\cdots$ / $\cdots$ | 2MJ07092139-1029344 | 10,11,12
PV Cep | 20:45:53.94 |  $+$67:57:38.7 | 356.5 | L1158 | N | HH315/Y |  $\cdots$ | 13
V350 Cep | 21:43:00.01 |  $+$66:11:28.0 | 872.5 | NG7129 |  $\cdots$ |  $\cdots$ / $\cdots$ |  $\cdots$ | 5

Note. — 1: Distances taken from GAIA EDR3 (Gaia Collaboration (2016, 2020). a
assumed; References: [1] Coffey et al. 2004; [2] Jensen et al. 2007; [3]
Leinert et al. 1993; [4] Banzatti et al. 2019; [5] Herbig 2008; [6] Kun et al.
2016; [7] Arce et al. 2010; [8] Young et al. 2015; [9] Akeson et al. 2019;
[10] Joncour et al. 2017; [11] Li et al. 2015; [12] Banzatti et al. 2014; [13]
Principe et al. 2018.

Table 2:  Photometric variability

ID | mag${}_{min}-$magmax | banda | Bursts | Ref
---|---|---|---|---
| (mag) |  |  |
EXors
XZ Tau | 16.6$-$10.4 | V | many/last 2014$-$2016 | 1,2
UZ Tau E | 18.5$-$12.5 | V | many/last 2006 | 3,4
VY Tau | 15.3$-$9.0 | V | 1960-1972/last 2014 | 2,3,5,6
NY Ori |  $>$16$-$13.3 | V | many/last 2001 | 7,8
V1143 Ori | 16.7$-$13.5 | V | many/last 1993 | 7,8
EXor candidates
DR Tau | 16.0$-$10.5 | V | 1960$-$1980 | 3,9,10,11
ASASSN-13db |  $>$17$-$13.2 | V | 2013,2014-2017 | 12,13
V1647 Ori | 23.0$-$17.8 | r′ | 1966,2003, 2008 | 14,15,16,17
iPTF 15afq | 19.0$-$16.7 | R | 2015,2019 | 18,19
PV Cep | 16$-$11 | R | 1977,1979,2008 | 20,21,22
V350 Cep |  $>$21$-$16.5 | B | 1978-2005 | 7,22,23

Note. — References: [1] Coffey et al. 2004; [2] Dodin et al. 2016; [3] Herbig
& Bell 1988; [4] Jensen et al. 2007; [5] Herbig 1990; [6] Stone 1983; [7]
Herbig 2008; [8] Jurdana-Šepić et al. 2017; [9] Chavarria 1979; [10] Kenyon et
al. 1994; [11] Grankin et al. 2007; [12] Holoien et al.2014; [13] Sicilia-
Aguilar et al. 2017; [14] Briceño et al. 2004; [15] Aspin et al. 2006; [16]
Aspin et al. 2008; [17] Aspin & Reipurth 2009; [18] Miller et al. 2015; [19]
Hillenbrand 2019; [20] Cohen et al. 1981; [21] Caratti o Garatti et al. 2013;
[22] Lorenzetti et al. 2011; [23] Jurdana-Šepić et al. 2018.

Table 3:  Stellar parameters

Name | AV | Lbol | SpT | $L_{\mathrm{*}}$ | $T_{\mathrm{eff}}$ | $M_{\mathrm{*}}$ | $L_{\mathrm{acc}}$ | $\dot{M}_{\mathrm{acc}}$ | Class | Ref
---|---|---|---|---|---|---|---|---|---|---
| (mag) | (L⊙) |  | (L⊙) | (K) | (M⊙) | (L⊙) | (10-7M⊙ yr-1) |  |
EXors
XZ Tau | 1.4$-$1.5 | 4.2$-$10.7 | M3 | 0.57 | 3410 | 0.37 | 0.347 | 1$-$10, 0.14 | II | 1,2,3,4
UZ Tau E | 0.5$-$1.5 | 1.7 | M3 | 0.62 | 3410 | 0.3 | 0.025 | 1$-$3 | II | 1,4,5,6
VY Tau | 0.5$-$0.85 | 0.72 | M1.5 | 0.39 | 3640 | 0.39 |  $\cdots$ |  $\cdots$ | II | 1,7,8
NY Ori | 0.3 |  $\cdots$ | K4 |  $\cdots$ | 4430 | 1.0 |  $\cdots$ |  $\cdots$ |  $\cdots$ | 9,10
V1143 Ori | 0 |  $\cdots$ | M2 | 0.20 | 3500 | 0.5 |  $\cdots$ |  $\cdots$ | II | 7
EXor candidates
DR Tau | 0.45$-$1.6 | 1.2 | K7 | 0.32 | 4070 | 0.75 | 0.74 | 4$-$9 | II | 1,4,11,12,13
ASASSN-13db | 0.04 | 0.1 | M5 | 0.03 | 3070-4855 | 0.15 |  $<$4 10${}^{-4}-$0.2 | 0.01$-$3.3 | II | 14,15
V1647 Ori | 10$-$20 | 9.5$-$90 | M0 | 5.20 | 3800 | 0.8 | 3.5$-$16 | 0.2$-$38 | I/f | 16,17,18,19,20
iPTF 15afq |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | I | 21
PV Cep | 10.8$-$12 | 41$-$100 | G4 | 17.0 | 5500 | 2.4$-$4 | 5.7$-$80 | 2$-$50 |  $\cdots$ | 22,23,24,25,26,27
V350 Cep |  $\cdots$ |  $\cdots$ | M0 |  $\cdots$ | 3560-3900 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | 28

Note. — References: [1] Herczeg & Hillenbrand 2014; [2] Osorio et al. 2016;
[3] Hartigan & Kenyon 2003; [4] Csépány, et al. 2017; [5] Yang et al. 2012;
[6] Lorenzetti et al. 2007; [7] Sipos & Kóspál 2014; [8] Cabrit et al. 1990;
[9] Hillenbrand 1997; [10] Kounkel et al. 2019; [11] Muzerolle et al. 2003;
[12] Lorenzetti et al. 2009; [13] Petrov et al. 2011; [14] Holoien et al.
2014; [15] Sicilia-Aguilar et al. 2017; [16] Aspin 2011; [17] Andrews et al.
2004; [18] Gramajo et al. 2014; [19] Moody & Stahler 2017; [20] Aspin et al.
2008; [21] Fischer et al. 2016; [22] Kun et al. 2011; [23] Hamidouche 2010;
[24] Magakian & Movsesian 2001; [25] Cohen et al. 1981; [26] The et al. 1994;
[27] Caratti o Garatti et al. 2013; [28] Kun et al. 2009.

Table 4:  Journal of observations.

Source | MODS | LUCI
---|---|---
| Date | MJD | Tint (B/R) | Phota | Filter | Date | MJD | Tint (zJ/HK) | Phota | Filter
| (yy/mm/dd) |  | (s) | (mag) |  | (yy/mm/dd) |  | (s) | (mag) |
XZ Tau | 2014/12/21 | 57012 | 600/600 | 12.54 | r |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
| 2018/03/21 | 58198 | 1200/1200 | 13.68 | r |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
UZ Tau E | 2014/12/22 | 57013 | 600/600 | 12.21 | r | 2016/01/25 | 57412 | 120/88 | 7.94 | [FeII]1.64
VY Tau | 2014/12/21 | 57012 | 600/600 | 12.88 | r | 2016/03/04 | 57451 | 160/80 | 9.43 | [FeII]1.64
NY Ori | 2014/12/21 | 57012 | 600/600 | 11.94 | r |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
| 2016/01/31 | 57418 | 600/600 | 14.05 | r |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
V1143 Ori | 2014/12/21 | 57012 | 1200/1200 | 16.23 | r |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
DR Tau | 2014/10/31 | 56961 | 150/150 | 12.311 (-1) | V | 2018/03/23 | 58200 | 30/30 | 9.08 | $J$
| 2018/03/19 | 58196 | 1200/1200 | 11.92 | r |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
ASASSN-13db | 2014/12/21 | 57012 | 1200/1200 | 13.94 | g | 2015/02/10 | 57063 | 1200/1200 | 10.65 | $H$
| 2015/10/01 | 57296 | 1200/1200 | 13.95 | g |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
| 2017/02/21 | 57805 | 1200/1200 | 17.66 | r |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
V1647 Ori | 2016/01/26 | 57413 |  $\cdots$ /1200 | 18.68 | r | 2016/01/26 | 57413 | 1200/600 | 12.10 | $J$
iPTF 15afq | 2016/01/27 | 57414 |  $\cdots$ /1200 | 19.7 | r | 2018/01/28 | 58146 | 1200/200 | 14.12 | $J$
| 2017/02/15 | 57799 |  $\cdots$ /1200 | 19.2 | r | 2019/12/02 | 58819 | 400/220 | 9.73 | $K$
PV Cep | 2015/06/24 | 57197 | 1200/1200 | 16.34 | g | 2016/12/08 | 57730 | 1200/240 | 9.06 | $J$
V350 Cep | 2019/07/04 | 58668 | 2400/2400 | 15.682 (0) | r | 2015/11/08 | 57334 | 800/480 | 11.783 (+6) | $H$
| 2019/10/02 | 58758 | 2400/2400 | 15.602 (0) | r | 2018/09/29 | 58390 | 800/ $\cdots$ | 12.96 | $J$
| 2019/10/17 | 58773 | 2400/2400 | 15.792 (-1) | r |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$

aafootnotetext: Photometry obtained in the indicated filter during the same
night as spectroscopy. Exceptions are flagged (1 = ASA-SN; 2 = ZTF, 3 = Campo
Imperatore Telescope, Italy) with in parenthesis the number of days before (-)
or after (+) the spectroscopic observation.

Table 5:  AV, and veiling derived from optical and near-infrared spectra.

Source | Optical | Near-infrared
---|---|---
| Date | Continuum | Lines | Date | Col-Col | Lines
|  | r${}_{710}^{a}$ | AV | Nl | AV |  | AV | Nl | AV
| (yy/mm/dd) |  | (mag) |  | (mag) | (yy/mm/dd) | (mag) |  | (mag)
XZ Tau | 2014/12/21 | 0.8 | 1.0$-$2.0 | 22 | 2.0$-$3.0 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
| 2018/03/21 | 0.0 | 1.4$-$3.4 | 16 | 2.8$-$4.4 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
UZ Tau E | 2014/12/22 | 0.7 | 1.6$-$3.0 | 25 | 2.0$-$3.0 | 2016/01/25 | 0.0$-$1.5 | 3 | 0.0$-$1.5
VY Tau | 2014/12/21 | 0.4 | 0.0$-$2.0 | 9 | 1.6$-$4.0 | 2016/03/04 |  $\cdots$ | 0 |  $\cdots$
NY Ori | 2014/12/21 | 0.2b |  $\cdots$ | 11 | 0.4$-$2.4 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
| 2016/01/31 | 0.1b |  $\cdots$ | 8 | 0.5$-$2.9 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
V1143 Ori | 2014/12/21 | 0.0 | 0.0$-$1.2 | 16 | 0.0$-$1.6 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
DR Tau | 2014/10/31 |  $>$0.5b |  $\cdots$ | 22 | 1.7$-$2.6 | 2018/03/23 | 0.5${}^{b}-$2.5 | 27c | 2.0$-$2.9
| 2018/03/19 |  $>$0.5b |  $\cdots$ | 27c | 2.0$-$2.9 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
ASASSN-13db | 2014/12/21 |  $>$3 |  $\cdots$ | 3 | 0.0$-$1.5 | 2015/02/10 | 0.0$-$1.5 | 1 |  $\cdots$
| 2015/10/01 |  $>$3 |  $\cdots$ | 3 | 0.0$-$1.5 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
| 2017/02/21 | 0.4 | 0.0$-$0.2 | 11 | 0.0$-$1.0 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
V1647 Ori | 2016/01/26 |  $>$2.5 |  $\cdots$ | 5 | 7.7$-$10.6 | 2016/01/26 | 12.0$-$14.0 | 3 | 12.0$-$14.0
iPTF 15afq | 2016/01/27 |  $\cdots$ |  $\cdots$ | 7 | 4.2$-$7.7 | 2018/01/28 |  $\cdots$ | 4 | 2.0$-$5.0
| 2017/02/15 |  $\cdots$ |  $\cdots$ | 5 | 5.9$-$8.9 | 2019/12/02 | 0.3$-$2.3 | 0 |  $\cdots$
PV Cep | 2015/06/24 | 0.0b |  $\cdots$ | 11 | 4.3$-$6.3 | 2016/12/08 | 0.0$-$1.0 | 5 | 1.8$-$5.1
V350 Cep | 2019/07/04 | 0.1 | 1.4$-$3.0 | 15 | 1.3$-$2.9 | 2015/11/08 | 0.8$-$2.8 | 4 | 0.0$-$1.5
| 2019/10/02 | 0.2 | 1.0$-$3.0 | 14 | 0.6$-$2.2 | 2018/09/29 |  $\cdots$ | 3 | 0.0$-$1.5
| 2019/10/17 | 0.0 | 1.3$-$3.3 | 16 | 0.7$-$2.3 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$

aafootnotetext: Typical error on r710 is 0.3 (Appendix B). bDerived by
assuming the AV value fitted from the accretion lines (see Appendix B for more
details). cOptical and near-infrared lines have been combined.

Table 6:  Accretion luminosity and mass accretion rate. Source | Date (yy/mm/dd) |  $L_{\mathrm{acc}}$/L⊙ | log($\dot{M}_{\mathrm{acc}}$/M⊙ yr-1)
---|---|---|---
XZ Tau | 2014/12/21 | 0.30 $-$ 0.90 | -7.16 $-$ -6.68
| 2018/03/21 | 0.10 $-$ 0.37 | -7.06 $-$ -7.63
UZ Tau E | 2014/12/22 | 0.40 $-$ 0.55 | -6.92 $-$ -6.78
| 2016/01/25 | 0.28 $-$ 0.50 | -7.07 $-$ -6.83
VY Tau | 2014/12/21 | 0.004 $-$ 0.04 | -9.13 $-$ -8.16
NY Ori | 2014/12/21 | 0.60 $-$ 3.60 |  $\cdots$
| 2016/01/31 | 0.03 $-$ 0.33 |  $\cdots$
V1143 Ori | 2014/12/21 | 0.006 $-$ 0.01 | -9.23 $-$ -9.11
DR Tau | 2014/10/31 | 1.31 $-$ 3.40 | -7.11 $-$ -6.69
| 2018/03/19-23 | 0.70 $-$ 1.74 | -7.38 $-$ -6.98
ASASSN-13db | 2014/12/21 | 0.21 $-$ 0.41 | -7.48 $-$ -7.17
| 2015/02/10 | 0.11 $-$ 0.14 | -7.76 $-$ -7.64
| 2015/10/01 | 0.20 $-$ 0.41 | -7.48 $-$ -7.18
| 2017/02/21 | 0.003 $-$ 0.009 | -9.26 $-$ -8.85
V1647 Ori | 2016/01/26 | 7.24 $-$ 20.89 | -5.72 $-$ -5.26
iPTF 15afq | 2016/01/27 | 0.72 $-$ 3.54 |  $\cdots$
| 2017/02/15 | 2.18 $-$ 11.20 |  $\cdots$
| 2018/01/28 | 0.58 $-$ 1.47 |  $\cdots$
PV Cep | 2015/06/24 | 4.78 $-$ 18.60 | -6.54 $-$ -5.95
| 2016/12/08 | 2.63 $-$ 8.04 | -6.79 $-$ -6.31
V350 Cep | 2015/11/08 | 0.49 $-$ 0.87 |  $\cdots$
| 2018/09/29 | 0.19 $-$ 0.32 |  $\cdots$
| 2019/07/04 | 0.82 $-$ 0.89 |  $\cdots$
| 2019/10/02 | 0.66 $-$ 1.09 |  $\cdots$
| 2019/10/17 | 0.67 $-$ 1.31 |  $\cdots$
Table 7:  Hydrogen physical parameters derived from Balmer, Paschen, and
Brackett line ratios.

Source | Date (yy/mm/dd) | Bright. level | Series | Nlines | Log[nH(cm-3)] | T (K)
---|---|---|---|---|---|---
XZ Tau | 2014/12/21 | H | Ba | 11 | 9.8 | 7500-8750
| 2018/03/21 | L | Ba | 9 | 9.4 | all
UZ Tau E | 2014/12/22 | L | Ba | 11 | 9.8 | 7500-8750
| 2016/01/25 | I | Pa+Br | 4 | 10.6-10.8 | 5000-12500
VY Tau | 2014/12/21 | L | Ba | 7 | 9.2-9.4 | all
NY Ori | 2014/12/21 | H | Pa | 4 | 9.0-10.4 | 5000-12500
V1143 Ori | 2014/12/21 | I | Ba | 11 | 9.2-9.4 | all
DR Tau | 2014/10/31 | L | Ba+Pa | 9 | 10.0-10.8 | 5000-12500
| 2018/03/19-23 | L | Ba+Pa+Br | 20 | 10.6-11.0 | 8750-12500
ASASSN-13db | 2017/02/21 | L | Ba | 5 | 9.4 | all
V1647 Ori | 2016/01/26 | ? | Ba+Pa | 4 | 9.0-10.0 | 8750-12500
iPTF 15afq | 2016/01/27 | L? | Pa | 5 | 9.0-10.6 | all
| 2018/01/28 | L? | Pa+Br | 4 | 10.4-11.0 | all
PV Cep | 2015/06/24 | H | Pa | 4 | 11.0 | 8750-12500
| 2016/12/08 | H | Pa+Br | 6 | 9.0-10.4 | all
V350 Cep | 2015/11/08 | I | Pa+Br | 5 | 10.6-10.8 | all
| 2019/07/04 | I | Ba+Pa | 8 | 9.2-9.8 | all
| 2019/10/02 | I | Ba+Pa | 8 | 9.2-9.8 | all
| 2019/10/17 | I | Ba+Pa | 8 | 9.2-9.8 | all

aafootnotetext: Brightness level as retrieved from photometric data. H:high;
I:intermediate; L:low.

##  Appendix A Appendix A

Fluxes of the lines used in the analysis. For optical lines ($\lambda<$ 1
$\mu$m) we give the air wavelength, for near-IR lines ($\lambda>$ 1 $\mu$m)
the vacuum wavelength.

Table 8:  XZ Tau: observed lines. |  | 2014/12/21 | 2018/03/21
---|---|---|---
Line | $\lambda_{air}$ | F$\pm\Delta$F
| (Å) | (10-14 erg s-1 cm-2)
H15 | 3711.977 | 3.2$\pm$0.2 |  $\cdots$
H14 | 3721.938 | 3.5$\pm$0.2 |  $\cdots$
H13 | 3734.368 | 3.6$\pm$0.2 |  $\cdots$
H12 | 3750.151 | 4.1$\pm$0.2 | 0.18$\pm$0.02
H11 | 3770.630 | 3.7$\pm$0.2 | 0.22$\pm$0.02
H10 | 3797.898 | 5.3$\pm$0.2 | 0.36$\pm$0.02
H9 | 3835.384 | 6.3$\pm$0.3 | 0.55$\pm$0.02
H8 | 3889.049 | 7.4$\pm$0.2 | 0.88$\pm$0.02
CaII K | 3933.660 | 51.3$\pm$0.1 | 4.4$\pm$0.02
CaII H/H7 | 3968.470 | 47.9$\pm$0.1 | 4.3$\pm$0.02
H$\delta$ | 4101.702 | 13.8$\pm$0.2 | 1.5$\pm$0.02
H$\gamma$ | 4340.464 | 20.4$\pm$0.2 | 3.0$\pm$0.05
H$\beta$ | 4861.325 | 50.2$\pm$0.3 | 7.6$\pm$0.05
HeI | 5015.678 | 9.2$\pm$0.2 | 0.6$\pm$0.03
HeI | 5875.621 | 7.1$\pm$0.2 | 0.9$\pm$0.08
[OI] | 6300.304 | 12.4$\pm$0.6 | 12.7$\pm$0.09
[OI] | 6363.776 | 3.7$\pm$0.6 | 4.2$\pm$0.09
H$\alpha$ | 6562.800 | 364$\pm$0.6 | 152$\pm$0.09
HeI | 6678.151 | 3.7$\pm$0.6 |  $\cdots$
[SII] | 6730.82 | 4.1$\pm$0.6 | 6.1$\pm$0.09
OI | 7773.055 | 2.6$\pm$0.9 |  $\cdots$
OI | 8446.360 | 19.4$\pm$1.2 | 6.0$\pm$2.0
CaII | 8498.020 | 217.3$\pm$1.2 | 44.1$\pm$1.0
CaII | 8542.090 | 207.2$\pm$1.2 | 46.3$\pm$1.0
CaII | 8662.140 | 181.9$\pm$1.2 | 38.1$\pm$1.0
Pa9 | 9229.015 | 15.2$\pm$1.8 |  $\cdots$
Table 9:  UZ Tau E: observed lines. 2014/12/22 | 2016/01/25
---|---
Line | $\lambda_{air}$ | F$\pm\Delta$F | Line | $\lambda_{vac}$ | F$\pm\Delta$F
| (Å) | (10-14 erg s-1 cm-2) |  | (Å) | (10-14 erg s-1 cm-2)
H15 | 3711.977 | 1.7$\pm$0.6 | Pa$\delta$ | 10052.128 | 59.3$\pm$1.6
H14 | 3721.938 | 2.5$\pm$0.6 | He I | 10833.2 | 68.8$\pm$1.0
H13 | 3734.368 | 2.5$\pm$0.6 | Pa$\gamma$ | 10941.090 | 92.0$\pm$0.9
H12 | 3750.151 | 3.2$\pm$0.6 | Pa$\beta$ | 12821.59 | 131$\pm$0.5
H11 | 3770.630 | 3.7$\pm$0.6 | Br12 | 16411.674 | 16.0$\pm$2.3
H10 | 3797.898 | 5.0$\pm$0.6 | Br11 | 16811.111 | 29.7$\pm$4.3
H9 | 3835.384 | 6.9$\pm$0.9 | Br10 | 17366.850 | 17.4$\pm$2.1
H8 | 3889.049 | 8.8$\pm$0.7 | CO 2-0 | 22930 | 14.5$\pm$2.7
CaII K | 3933.660 | 20.8$\pm$0.7 | CO 3-1 | 23230 | 18.5$\pm$2.7
CaII H/H7 | 3968.470 | 23.1$\pm$0.7 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 4026.191 | 1.4$\pm$0.7 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\delta$ | 4101.702 | 14.8$\pm$0.3 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\gamma$ | 4340.464 | 20.0$\pm$0.5 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 4471.480 | 3.1$\pm$0.5 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\beta$ | 4861.325 | 42.7$\pm$0.3 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 4921.931 | 4.9$\pm$0.4 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 5015.678 | 4.3$\pm$0.4 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 5875.621 | 7.9$\pm$0.4 |  $\cdots$ |  $\cdots$ |  $\cdots$
[OI] | 6300.304 | 14.0$\pm$0.4 |  $\cdots$ |  $\cdots$ |  $\cdots$
[OI] | 6363.776 | 3.3$\pm$0.4 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\alpha$ | 6562.800 | 357$\pm$0.4 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 6678.151 | 3.6$\pm$0.4 |  $\cdots$ |  $\cdots$ |  $\cdots$
[SII] | 6730.82 | 2.1$\pm$0.4 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 7065.190 | 3.6$\pm$0.9 |  $\cdots$ |  $\cdots$ |  $\cdots$
OI | 8446.360 | 22.4$\pm$0.9 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8498.020 | 92.4$\pm$1.0 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8542.090 | 81.2$\pm$1.0 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8662.140 | 78.8$\pm$1.0 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa9 | 9229.015 | 17.9$\pm$1.3 |  $\cdots$ |  $\cdots$ |  $\cdots$
Table 10:  VY Tau: observed lines. 2014/12/21 | 2016/03/04
---|---
Line | $\lambda_{air}$ | F$\pm\Delta$F | Line | $\lambda_{vac}$ | F$\pm\Delta$F
| (Å) | (10-14 erg s-1 cm-2) |  | (Å) | (10-14 erg s-1 cm-2)
H10 | 3797.898 | 0.16$\pm$0.04 | He I | 10833.2 | 0.5$\pm$0.1
H9 | 3835.384 | 0.19$\pm$0.08 | CO 2-0 | 22930 | -9.2$\pm$0.7
H8 | 3889.049 | 0.33$\pm$0.12 | CO 3-1 | 23230 | -3.9$\pm$0.7
CaII K | 3933.660 | 1.94$\pm$0.12 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII H/H7 | 3968.470 | 2.55$\pm$0.12 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\delta$ | 4101.702 | 0.47$\pm$0.06 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\gamma$ | 4340.464 | 0.97$\pm$0.08 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\beta$ | 4861.325 | 2.83$\pm$0.13 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\alpha$ | 6562.800 | 11.3$\pm$0.28 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa9 | 9229.015 | 2.37$\pm$0.68 |  $\cdots$ |  $\cdots$ |  $\cdots$
Table 11:  NY Ori: observed lines. |  | 2014/12/21 | 2016/01/31
---|---|---|---
Line | $\lambda_{air}$ | F$\pm\Delta$F
| (Å) | (10-14 erg s-1 cm-2)
H15 | 3711.977 | -1.7$\pm$0.4 | -0.09$\pm$0.02
H14 | 3721.938 | -2.2$\pm$0.4 | -0.15$\pm$0.02
H13 | 3734.368 | -2.4$\pm$0.6 | -0.23$\pm$0.02
H12 | 3750.151 | -3.4$\pm$0.4 | -0.22$\pm$0.02
H11 | 3770.630 | -3.0$\pm$0.4 | -0.23$\pm$0.02
H10 | 3797.898 | 1.7(-3.6)$\pm$0.4 | -0.44$\pm$0.02
H9 | 3835.384 | -5.1$\pm$0.4 | -0.76$\pm$0.02
H8 | 3889.049 | 2.6(-4.1)$\pm$0.4 | 0.1(-0.4)$\pm$0.02
CaII K | 3933.660 | 14.7$\pm$0.1 | 0.6$\pm$0.02
CaII H/H7 | 3968.470 | 11.0(-3.4)$\pm$0.1 | 0.4(-0.1)$\pm$0.02
HeI | 4026.191 | 1.4$\pm$0.1 | 0.07$\pm$0.03
H$\delta$ | 4101.702 | 3.7(-5.4)$\pm$0.3 | 0.3(-0.6)$\pm$0.03
H$\gamma$ | 4340.464 | 4.9(-5.1)$\pm$0.3 | 0.4(-0.5)$\pm$0.03
HeI | 4471.480 | 1.7$\pm$0.3 |  $\cdots$
H$\beta$ | 4861.325 | 9.1(-3.6)$\pm$0.3 | 0.9(-0.2)$\pm$0.03
HeI | 4921.931 | 2.0(-1.5)$\pm$0.1 |  $\cdots$
HeI | 5015.678 | 2.0(-2.4)$\pm$0.1 |  $\cdots$
HeI | 5875.621 | 6.5$\pm$0.4 | 0.4$\pm$0.04
[OI] | 6300.304 | 0.9$\pm$0.3 | 0.3$\pm$0.04
H$\alpha$ | 6562.800 | 188$\pm$0.4 | 28.7$\pm$0.04
HeI | 6678.151 | 2.8$\pm$0.3 | 0.3$\pm$0.04
OI | 8446.360 | 4.7(-2.3)$\pm$0.2 | (-0.8)$\pm$0.1
Pa17 | 8467.253 |  $\cdots$ | (-0.4)$\pm$0.1
CaII | 8498.020 | 52.7$\pm$0.2 | 2.4$\pm$0.1
Pa16 | 8502.483 |  $\cdots$ | (-0.3)$\pm$0.1
CaII | 8542.090 | 42.4$\pm$0.2 | 1.2$\pm$0.1
Pa15 | 8545.383 |  $\cdots$ | (-0.5)$\pm$0.1
Pa14 | 8598.392 | 2.3$\pm$0.3 | (-0.8)$\pm$0.1
CaII | 8662.140 | 39.0$\pm$0.2 | 1.0$\pm$0.1
Pa13 | 8665.018 |  $\cdots$ | -0.8$\pm$0.1
Pa12 | 8750.472 | 3.7$\pm$0.2 | -1.1$\pm$0.2
Pa11 | 8862.783 | 4.2$\pm$0.2 | -1.3$\pm$0.2
Pa10 | 9014.910 | 6.1$\pm$0.3 | -1.6$\pm$0.2
Pa9 | 9229.015 | 10.2$\pm$0.3 | -2.0$\pm$0.4
Table 12:  V1143 Ori: observed lines. 2014/12/21
---
Line | $\lambda_{air}$ | F$\pm\Delta$F
| (Å) | (10-14 erg s-1 cm-2)
H14 | 3721.938 | 0.04$\pm$0.01
H13 | 3734.368 | 0.05$\pm$0.01
H12 | 3750.151 | 0.055$\pm$0.005
H11 | 3770.630 | 0.067$\pm$0.005
H10 | 3797.898 | 0.111$\pm$0.006
H9 | 3835.384 | 0.161$\pm$0.006
H8 | 3889.049 | 0.209$\pm$0.006
CaII K | 3933.660 | 0.371$\pm$0.005
CaII H/H7 | 3968.470 | 0.540$\pm$0.005
H$\delta$ | 4101.702 | 0.322$\pm$0.006
H$\gamma$ | 4340.464 | 0.370$\pm$0.006
HeI | 4471.480 | 0.043$\pm$0.006
H$\beta$ | 4861.325 | 0.486$\pm$0.005
HeI | 4921.931 | 0.058$\pm$0.006
HeI | 5015.678 | 0.056$\pm$0.007
HeI | 5875.621 | 0.11$\pm$0.01
H$\alpha$ | 6562.800 | 1.55$\pm$0.01
Table 13:  DR Tau: observed lines. |  | 2014/10/31 | 2018/03/19 | 2018/03/23
---|---|---|---|---
Line | $\lambda_{air}$ | F$\pm\Delta$F | Line | $\lambda_{vac}$ | F$\pm\Delta$F
| (Å) | (10-14 erg s-1 cm-2) |  | (Å) | (10-14 erg s-1 cm-2)
H18 | 3691.554 |  $\cdots$ | 1.04$\pm$0.06 | Pa$\delta$ | 10052.128 | 70.1$\pm$1.8
H17 | 3697.151 |  $\cdots$ | 1.35$\pm$0.06 | [SII] | 10339.24 | 4.3$\pm$1.8
H16 | 3703.852 |  $\cdots$ | 2.57$\pm$0.08 | He I | 10833.2 | (-35.7)12.1$\pm$0.1
H15 | 3711.977 | 3.0$\pm$0.2 | 2.90$\pm$0.06 | Pa$\gamma$ | 10941.090 | 106.0$\pm$1.8
H14 | 3721.938 | 4.3$\pm$0.3 | 4.1$\pm$0.1 | Pa$\beta$ | 12821.59 | 173.1$\pm$1.8
H13 | 3734.368 | 3.5$\pm$0.4 | 4.7$\pm$0.1 | Br14 | 15884.880 | 10.5$\pm$1.6
H12 | 3750.152 | 4.4$\pm$0.5 | 4.8$\pm$0.1 | Br13 | 16113.714 | 15.4$\pm$1.6
H11 | 3770.630 |  $\cdots$ | 4.4$\pm$0.2 | Br12 | 16411.674 | 22.1$\pm$1.6
H10 | 3797.897 |  $\cdots$ | 4.9$\pm$0.2 | Br11 | 16811.111 | 31.5$\pm$1.6
H9 | 3835.384 |  $\cdots$ | 5.7$\pm$0.1 | Br10 | 17366.850 | 36.3$\pm$1.6
CaII K | 3933.660 | 53.5$\pm$0.3 | 18.1$\pm$0.1 | Br$\gamma$ | 21661.20 | 31.4$\pm$1.8
CaII H/H7 | 3968.470 | 25.3$\pm$0.3 | 14.0(-0.14)$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 4026.191 | 2.5$\pm$0.4 | 1.6$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\delta$ | 4101.702 | (-2.5)3.1$\pm$0.3 | (-1.4)5.1$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\gamma$ | 4340.464 | (-1.6)6.0$\pm$0.3 | (-1.7)5.0$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 4471.480 | 7.5$\pm$0.7 | 4.1$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeII | 4685.804 | 1.9$\pm$0.7 | 1.5$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 4713.146 |  $\cdots$ | 0.6$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\beta$ | 4861.325 | (-1.5)30.8$\pm$0.1 | (-2.2)16.5$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 4921.931 | 15.4$\pm$0.3 | 4.5(-0.4)$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 5015.678 | 17.0$\pm$0.3 | 5.5(-0.5)$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 5875.621 | 15.3$\pm$0.3 | 11.4$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
[OI] | 6300.304 | 1.8$\pm$0.4 | 2.4$\pm$0.1 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\alpha$ | 6562.800 | 407.5$\pm$0.5 | 260.5$\pm$0.3 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 6678.151 | 7.2$\pm$0.4 | 5.7$\pm$0.3 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 7065.190 | 6.8$\pm$0.4 | 4.1$\pm$0.3 |  $\cdots$ |  $\cdots$ |  $\cdots$
OI | 7773.055 | 12.3$\pm$0.4 | 2.9(-3.7)$\pm$0.4 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa25 | 8323.424 |  $\cdots$ | 0.8$\pm$0.2 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa24 | 8333.782 | 4.3$\pm$0.4 | 1.2$\pm$0.2 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa23 | 8345.552 | 5.0$\pm$0.4 | 1.4$\pm$0.2 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa22 | 8359.003 | 5.7$\pm$0.5 | 2.0$\pm$0.2 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa21 | 8374.475 | 8.4$\pm$0.6 | 2.2$\pm$0.2 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa20 | 8392.396 | 10.7$\pm$0.6 | 2.8$\pm$0.2 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa19 | 8413.317 | 10.7$\pm$0.5 | 4.1$\pm$0.2 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa18 | 8437.955 | 15.3$\pm$0.6 | 5.3$\pm$0.2 |  $\cdots$ |  $\cdots$ |  $\cdots$
OI | 8446.360 | 44.3$\pm$0.4 | 17.9$\pm$0.6 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa17 | 8467.253 | 14.3$\pm$0.5 | 5.7$\pm$0.6 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8498.020 | 237.0$\pm$0.4 | 74.7$\pm$0.3 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8542.090 | 253.5$\pm$0.4 | 81.8$\pm$0.3 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa14 | 8598.392 | 23.8$\pm$0.5 | 13.0$\pm$0.3 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8662.140 | 214.8$\pm$0.4 | 70.8$\pm$0.5 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa12 | 8750.472 | 31.3$\pm$0.5 | 14.9$\pm$0.8 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa11 | 8862.783 | 34.4$\pm$0.5 | 19.1$\pm$0.8 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa10 | 9014.910 | 39.4$\pm$0.5 | 20.6$\pm$0.8 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa9 | 9229.015 | 53.0$\pm$0.5 | 26.5$\pm$0.8 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa8 | 9545.971 | 63.1$\pm$2 | 34.7$\pm$1.3 |  $\cdots$ |  $\cdots$ |  $\cdots$
Table 14:  ASASSN-13db: observed lines. |  | 2014/12/21 | 2015/10/01 | 2017/02/21 | 2015/02/10
---|---|---|---|---|---
Line | $\lambda_{air}$ | F$\pm\Delta$F | Line | $\lambda_{vac}$ | F$\pm\Delta$F
| (Å) | (10-14 erg s-1 cm-2) |  | (Å) | (10-14 erg s-1 cm-2)
H9 | 3835.384 |  $\cdots$ |  $\cdots$ | 0.10$\pm$0.02 | He I | 10833.2 | (-13.3)$\pm$0.1
CaII K | 3933.660 | (-1.9)1.2$\pm$0.2 | (-1.3)6.2$\pm$0.3 | (-0.01)0.21$\pm$0.01 | Pa$\beta$ | 12821.59 | 5.7$\pm$1.5
CaII H/H7 | 3968.470 | (-1.0)0.9$\pm$0.2 | (-1.5)3.6$\pm$0.3 | (-0.01)0.32$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\delta$ | 4101.702 |  $\cdots$ |  $\cdots$ | 0.25$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\gamma$ | 4340.464 |  $\cdots$ |  $\cdots$ | 0.31$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\beta$ | 4861.325 | -2.6$\pm$0.2 | -3.0$\pm$0.3 | 0.47$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 4921.931 |  $\cdots$ |  $\cdots$ | 0.08$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 5015.678 |  $\cdots$ |  $\cdots$ | 0.07$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 5875.621 |  $\cdots$ |  $\cdots$ | 0.08$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
[OI] | 6300.304 |  $\cdots$ |  $\cdots$ | 0.09$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
[OI] | 6363.776 |  $\cdots$ |  $\cdots$ | 0.04$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\alpha$ | 6562.800 | (-1.4)11.5$\pm$1.0 | (-1.4)27.3$\pm$0.4 | 1.21$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 6678.151 |  $\cdots$ |  $\cdots$ | 0.06$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
[SII] | 6716.44 |  $\cdots$ |  $\cdots$ | 0.08$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
[SII] | 6730.82 |  $\cdots$ |  $\cdots$ | 0.18$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8498.020 | 13.1$\pm$0.3 | 12.8$\pm$0.4 | 0.23$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8542.090 | 13.5$\pm$0.3 | 13.2$\pm$0.4 | 0.24$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8662.140 | 10.2$\pm$0.3 | 9.8$\pm$0.4 | 0.19$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$
Table 15:  V1647 Ori: observed lines. 2016/01/26
---
Line | $\lambda_{air/vac}$ | F$\pm\Delta$F
| (Å) | (10-14 erg s-1 cm-2)
[OI] | 6300.304 | 0.22$\pm$0.06
H$\alpha$ | 6562.800 | (-0.04)0.60$\pm$0.02
OI | 7773.055 | 0.03$\pm$0.04
OI | 8446.360 | 0.18$\pm$0.04
CaII | 8498.020 | 1.09$\pm$0.04
CaII | 8542.090 | 1.20$\pm$0.04
CaII | 8662.140 | 1.09$\pm$0.04
Pa$\delta$ | 10052.128 | 2.8$\pm$0.3
He I | 10833.2 | (-29.8)$\pm$1.0
Pa$\gamma$ | 10941.090 | 10.3$\pm$1.0
Pa$\beta$ | 12821.59 | 73.2$\pm$2.0
Table 16:  iPTF15afq: observed lines. |  | 2016/01/27 | 2017/02/15 |  |  | 2018/01/28 | 2019/12/02
---|---|---|---|---|---|---|---
Line | $\lambda_{air}$ | F$\pm\Delta$F | Line | $\lambda_{vac}$ | F$\pm\Delta$F
| (Å) | (10-14 erg s-1 cm-2) |  | (Å) | (10-14 erg s-1 cm-2)
[OI] | 6300.304 | 0.35$\pm$0.01 | 0.21$\pm$0.01 | Pa$\delta$ | 10052.128 | 0.46$\pm$0.02 | -0.6$\pm$0.1
[OI] | 6363.776 | 0.11$\pm$0.01 | 0.08$\pm$0.01 | [SII] | 10320.49 | 0.10$\pm$0.02 |  $\cdots$
H$\alpha$ | 6562.800 | 1.46$\pm$0.01 | 1.14$\pm$0.01 | [SII] | 10336.41 | 0.16$\pm$0.02 |  $\cdots$
[NII] | 6583.46 | 0.03$\pm$0.01 |  $\cdots$ | He I | 10833.2 | 0.62$\pm$0.02 | -2.2$\pm$0.1
[SII] | 6716.44 | 0.04$\pm$0.01 | 0.03$\pm$0.01 | Pa$\gamma$ | 10941.090 | 0.67$\pm$0.02 | -1.6$\pm$0.2
[SII] | 6730.82 | 0.08$\pm$0.01 | 0.06$\pm$0.01 | [FeII] | 12570.2068 | 0.26$\pm$0.02 |  $\cdots$
[FeII] | 7155.1742 | 0.14$\pm$0.01 | 0.09$\pm$0.01 | Pa$\beta$ | 12821.59 | 1.61$\pm$0.02 | -1.9$\pm$0.2
[FeII] | 7167.9228 | 0.04$\pm$0.01 |  $\cdots$ | Br8 | 19450.871 | 0.6$\pm$0.1 |  $\cdots$
[FeII] | 7288.9899 | 0.05$\pm$0.01 |  $\cdots$ | H2 1-0S(1) | 21218. | 0.12$\pm$0.02 |  $\cdots$
[CaII] | 7323.89 | 0.10$\pm$0.01 |  $\cdots$ | Br$\gamma$ | 21661.20 | 0.54$\pm$0.04 |  $\cdots$
[NiII] | 7377.83 | 0.08$\pm$0.01 |  $\cdots$ | CO 2-0 | 22930 | 2.9$\pm$0.2 |  $\cdots$
[FeII] | 7388.1673 | 0.05$\pm$0.01 |  $\cdots$ | CO 3-1 | 23230 | 3.8$\pm$0.2 | (-4.0)3.6$\pm$0.5
[NiII] | 7411.61 | 0.04$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
[FeII] | 7452.5611 | 0.06$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
[FeII] | 8387.181 | 0.09$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
[FeII] | 8446.360 | 5.13$\pm$0.01 | 0.08$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa17 | 8467.253 | 0.04$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8498.020 | 8.15$\pm$0.01 | 1.01$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8542.090 | 9.84$\pm$0.01 | 1.07$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa14 | 8598.392 | 0.03$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8662.140 | 0.89$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa12 | 8750.472 | 0.04$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa11 | 8862.783 | 0.06$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa10 | 9014.910 | 0.09$\pm$0.01 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa9 | 9229.015 | 0.18$\pm$0.02 | 0.12$\pm$0.04 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa8 | 9545.971 | 0.14$\pm$0.02 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$
Table 17:  PV Cep: observed lines. |  | 2015/06/24 |  |  | 2016/12/08
---|---|---|---|---|---
Line | $\lambda_{air}$ | F$\pm\Delta$F | Line | $\lambda_{vac}$ | F$\pm\Delta$F
| (Å) | (10-14 erg s-1 cm-2) |  | (Å) | (10-14 erg s-1 cm-2)
H15 | 3711.977 | -0.05$\pm$0.02 | Pa8 | 9548.590 | 32$\pm$2
H14 | 3721.938 | -0.13$\pm$0.02 | Pa$\delta$ | 10052.128 | 48$\pm$2
H13 | 3734.368 | -0.10$\pm$0.02 | [SII] | 10320.49 | 36$\pm$3
H12 | 3750.151 | -0.10$\pm$0.02 | [NI] | 10400.59 | 20$\pm$3
H11 | 3770.630 | -0.10$\pm$0.02 | He I | 10833.2 | -43$\pm$1
H10 | 3797.898 | -0.15$\pm$0.02 | Pa$\gamma$ | 10941.090 | 61$\pm$3
H9 | 3835.384 | -0.22$\pm$0.02 | Pa$\beta$ | 12821.59 | 210$\pm$2
H8 | 3889.049 | -0.38$\pm$0.02 | Br14 | 15884.880 | 34$\pm$5
CaII K | 3933.660 | (-0.04)0.73$\pm$0.02 | Br13 | 16113.714 | 50$\pm$5
CaII H/H7 | 3968.470 | -0.28$\pm$0.02 | Br12 | 16411.674 | 60$\pm$5
H$\delta$ | 4101.702 | -0.23$\pm$0.02 | Br11 | 16811.111 | 65$\pm$5
H$\gamma$ | 4340.464 | -0.39$\pm$0.02 | Br10 | 17366.850 | 72$\pm$10
H$\beta$ | 4861.325 | (-0.39)1.58$\pm$0.02 | Br$\gamma$ | 21661.20 | 49$\pm$8
HeI | 4921.931 | 0.61$\pm$0.04 | C0 2-0 | 22930 | 161$\pm$20
HeI | 5015.678 | 0.47$\pm$0.04 |  $\cdots$ |  $\cdots$ |  $\cdots$
[OI] | 6300.304 | 4.68$\pm$0.07 |  $\cdots$ |  $\cdots$ |  $\cdots$
[OI] | 6363.776 | 1.61$\pm$0.07 |  $\cdots$ |  $\cdots$ |  $\cdots$
H$\alpha$ | 6562.800 | 52.8$\pm$0.04 |  $\cdots$ |  $\cdots$ |  $\cdots$
HeI | 6678.151 | 0.65$\pm$0.04 |  $\cdots$ |  $\cdots$ |  $\cdots$
OI | 7773.055 | 2.11$\pm$0.09 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa19 | 8413.317 | 2.00$\pm$0.09 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa18 | 8437.955 | 2.51$\pm$0.09 |  $\cdots$ |  $\cdots$ |  $\cdots$
OI | 8446.360 | 13.9$\pm$0.09 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa17 | 8467.253 | 3.97$\pm$0.09 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8498.020 | 79.6$\pm$0.09 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8542.090 | 69.4$\pm$0.09 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa14 | 8598.392 | 4.41$\pm$0.09 |  $\cdots$ |  $\cdots$ |  $\cdots$
CaII | 8662.140 | 61.9$\pm$0.09 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa12 | 8750.472 | 5.81$\pm$0.09 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa11 | 8862.783 | 6.34$\pm$0.09 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa10 | 9014.910 | 7.22$\pm$0.13 |  $\cdots$ |  $\cdots$ |  $\cdots$
Pa9 | 9229.015 | 8.16$\pm$0.13 |  $\cdots$ |  $\cdots$ |  $\cdots$
Table 18: V350 Cep: observed lines. |  | 2015/11/08 | 2018/09/29 |  |  | 2019/07/04 | 2019/10/02 | 2019/10/17
---|---|---|---|---|---|---|---|---
Line | $\lambda_{vac}$ | F$\pm\Delta$F | Line | $\lambda_{air}$ | F$\pm\Delta$F
| (Å) | (10-14 erg s-1 cm-2) |  | (Å) | (10-14 erg s-1 cm-2)
Pa$\delta$ | 10052.128 | 2.8$\pm$0.2 | 1.1$\pm$0.1 | CaII K | 3933.660 | 1.48$\pm$0.08 | 2.17$\pm$0.03 | 2.24$\pm$0.02
[SII] | 10320.49 | 1.2$\pm$0.2 |  $\cdots$ | CaII H/H7 | 3968.470 | 0.85$\pm$0.07 | 0.67$\pm$0.03 | 1.03$\pm$0.02
He I | 10833.2 | -2.2$\pm$0.1 | -1.2$\pm$0.1 | H$\delta$ | 4101.702 | 0.27$\pm$0.07 | 0.48$\pm$0.08 | 0.48$\pm$0.03
Pa$\gamma$ | 10941.090 | 2.9$\pm$0.1 | 1.5$\pm$0.1 | H$\gamma$ | 4340.464 | 0.32$\pm$0.05 | 0.65$\pm$0.08 | 0.58$\pm$0.03
Pa$\beta$ | 12821.59 | 4.1$\pm$0.1 | 1.6$\pm$0.1 | H$\beta$ | 4861.325 | 0.78$\pm$0.04 | (-0.09)1.50$\pm$0.02 | 1.45$\pm$0.03
Br12 | 16411.674 | 1.4$\pm$0.2 |  $\cdots$ | HeI | 4921.931 | 0.39$\pm$0.05 | 0.61$\pm$0.02 | 0.63$\pm$0.03
Br11 | 16811.111 | 1.9$\pm$0.5 |  $\cdots$ | HeI | 5015.678 | 0.35$\pm$0.03 | 0.54$\pm$0.04 | 0.55$\pm$0.03
Br10 | 17366.850 | 1.0$\pm$0.2 |  $\cdots$ | [OI] | 6300.304 | 0.88$\pm$0.04 | 0.96$\pm$0.01 | 1.24$\pm$0.04
H2 1-0S(1) | 21218. | 0.38$\pm$0.02 |  $\cdots$ | [OI] | 6363.776 | 0.30$\pm$0.04 | 0.37$\pm$0.01 | 0.36$\pm$0.04
Br$\gamma$ | 21661.20 | 0.97$\pm$0.07 |  $\cdots$ | H$\alpha$ | 6562.800 | 9.76$\pm$0.04 | 16.1$\pm$0.01 | 15.0$\pm$0.04
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | HeI | 6678.151 | 0.14$\pm$0.04 | 0.19$\pm$0.01 | 0.17$\pm$0.04
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | OI | 7773.055 | 0.46$\pm$0.09 | 0.43$\pm$0.01 | 0.49$\pm$0.04
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | Pa19 | 8413.317 | 0.07$\pm$0.02 | 0.06$\pm$0.02 |  $\cdots$
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | Pa18 | 8437.955 | 0.10$\pm$0.02 |  $\cdots$ |  $\cdots$
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | OI | 8446.360 | 0.59$\pm$0.02 | 0.92$\pm$0.06 | 0.77$\pm$0.05
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | Pa17 | 8467.253 | 0.33$\pm$0.02 | 0.32$\pm$0.06 | 0.35$\pm$0.05
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | CaII | 8498.020 | 7.04$\pm$0.02 | 6.95$\pm$0.06 | 7.93$\pm$0.05
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | CaII | 8542.090 | 8.17$\pm$0.02 | 8.14$\pm$0.06 | 9.27$\pm$0.05
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | Pa14 | 8598.392 | 0.35$\pm$0.04 | 0.41$\pm$0.06 | 0.40$\pm$0.05
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | CaII | 8662.140 | 7.48$\pm$0.04 | 7.30$\pm$0.06 | 8.16$\pm$0.05
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | Pa12 | 8750.472 | 0.46$\pm$0.04 | 0.56$\pm$0.06 | 0.46$\pm$0.05
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | Pa11 | 8862.783 | 0.47$\pm$0.04 | 0.50$\pm$0.06 | 0.55$\pm$0.05
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | Pa10 | 9014.910 | 0.6$\pm$0.1 | 0.63$\pm$0.06 | 1.04$\pm$0.06
$\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | Pa9 | 9229.015 | 1.0$\pm$0.1 | 0.97$\pm$0.07 | 1.4$\pm$0.2

##  Appendix B Appendix B: Veiling determination

Figure 12: Examples of BT-Settl models best matching the TiO bands. Spectra
are normalized to the peak of the band. For each panel, object name, date of
observation, temperature of the BT-Settl model, and veiling are reported. In
the third panel TiO bands are indicated. Figure 13:  Residuals on
$T_{\mathrm{eff}}$ fitting in UZ Tau E, assuming r=0.7. The BTSettl model that
best matches the MODS spectrum corresponds to $T_{\mathrm{eff}}$ = 3600 K, the
same as that found with high-spectral resolution observations.

In this Section we briefly discuss the strategy we adopted to estimate the
veiling ($r$) in our sample. For a given source, we have firstly selected a
number of BTSettl templates with temperature within 500 K of the
$T_{\mathrm{eff}}$ value given in the literature (Table 3) and log$g$ between
3.5 and 4.5. A veiling (increasing in steps of 0.1) is then added to each
template to find the value that best matches the spectral features observed in
our sample of stars.

For M-type stars, which are the majority of our sources, we have fitted the
portion of the continuum in the range 700-715 nm, where prominent TiO
absorption bands are located (Fig. 12). In all sources but two, our fit
consistently provides a $T_{\mathrm{eff}}$ in agreement with that of the
literature, and determined $r$ with a typical error of 0.3, corresponding to a
variation of $T_{\mathrm{eff}}$ less than 100 K. The two exceptions are UZ Tau
E and V1143 Ori ($T_{\mathrm{eff}}$ of literature 3410 K and 3500 K,
respectively). In the case of UZ Tau E we have accurate determinations of
$T_{\mathrm{eff}}$ = 3600 K and $r$ = 0.7 from high-resolution observations
obtained with the HARPS-N instrument at the Telescopio Nazionale Galileo
(Gangi et al. in preparation). Having assumed the veiling value, we have
fitted $T_{\mathrm{eff}}$ in the MODS spectrum, which is in perfect agreement
with that derived from the HARPS-N spectrum (Figure 13). In the case of V1143
Ori, we fit $T_{\mathrm{eff}}$ = 3200 K and $r$ = 0. Since for $r>$0
$T_{\mathrm{eff}}$ is even lower, the literature value cannot be reproduced.
For the few sources that do not present the TiO bands because of their SpT (in
particular NY Ori, DR Tau, and PV Cep), we applied an indirect method to
evaluate the veiling at 710 nm. We firstly reddened the BTSettl template by
assuming the AV derived from the ’line method’ (Section 6.1.1). By dividing
the observed spectrum with the reddened template, we have derived the shape of
the excess continuum flux. The ratio at 710 nm was then assumed as the veiling
value, reported in Table 6.

We also looked for TiO spectral variations in the three sources of this sub-
sample that were observed more than once (XZ Tau, V350 Cep and ASASSN13-db,
see Fig. 14). We have found that XZ Tau shows an increase in the depth of the
TiO bands between 2014 and 2018 which corresponds to a decrease in $r$ from
0.8 to 0.0. TiO bands do not appear in the spectra of ASASSN13-db during the
burst phases of 2014 and 2015 while they are present during the out-of-burst
phase of 2017 with a corresponding veiling of 0.4. The absence of TiO bands in
2014 and 2015 indicates that during the outburst the veiling increased to
values larger than 3. Finally, the veiling measured in V350 Cep does not show
variations that exceed the quoted uncertainty of 0.3.

Figure 14: TiO bands strength variability. Spectra are normalized to the peak
of the band and vertically shifted for a better visualization. Date of
observation, temperature of the best matching BT-Settl model and veiling are
also labelled.