# Chandra X-ray Analysis of Herbig Ae/Be Stars

Hema Anilkumar,1 Blesson Mathew,1 V. Jithesh,1 Sreeja S. Kartha,1 P. Manoj,2
Mayank Narang3 and Mahathi Chavali,4
1Department of Physics & Electronics, CHRIST (Deemed to be University),
Bangalore 560029, India
2Department of Astronomy and Astrophysics, Tata Institute of Fundamental
Research Homi Bhabha Road, Colaba, Mumbai 400005, India
3Academia Sinica Institute of Astronomy & Astrophysics, 11F of Astro-Math
Bldg., No. 1, Sec. 4, Roosevelt Road, Taipei 10617, Taiwan, Republic of China
4Max Planck Institute for Radio Astronomy, Auf dem Hügel 69, 53121 Bonn,
Germany E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

Herbig Ae/Be (HAeBe) stars are intermediate-mass pre-main sequence stars,
characterized by infrared excess and emission lines. They are observed to emit
X-rays, whose origin is a matter of discussion and not settled yet. X-ray
emission is not expected in HAeBe stars, as they lack the sub-surface
convective zone. In this study, we retrieved observations from the Chandra
archive for 62 HAeBe stars, among which 44 sources (detection fraction
$\sim$71%) were detected in X-rays, with 7 being new detections. We use this
sample as a test bed to conduct a comparative analysis of the X-ray properties
of HAeBe stars and their low-mass counterparts, T Tauri Stars (TTSs). Further,
we compare the X-ray properties of HAeBe stars and TTSs with optical and IR
properties to constrain the X-ray emission mechanism in HAeBe stars. We found
no correlation between X-ray emission and disk properties of HAeBe stars,
confirming that X-rays are not related to accretion shocks. About 56% of HAeBe
stars without any known sub-arcsec companions have lower plasma temperatures
(kT $\leq$ 2 keV). We observe flaring/variability in HAeBe stars with
confirmed low-mass companions. These stars show plasma temperatures > 2 keV,
similar to TTSs. Guided by this information we discuss the role of a T Tauri
companion for X-ray emission seen in our sample of HAeBe stars. From the
results obtained in this paper, we suggest that X-ray emission from HAeBe
stars may not be related to accretion shocks or hidden TTS, but rather can be
due to magnetically driven coronal emission.

###### keywords:

stars: coronae – stars: flare – stars: low-mass – stars: massive – stars: pre-
main-sequence – techniques: imaging spectroscopy

††pubyear: 2023††pagerange: Chandra X-ray Analysis of Herbig Ae/Be Stars–D

##  1 Introduction

Herbig Ae/Be (HAeBe) stars are intermediate mass (2 M⊙ $\leq$ M⋆ $\leq$ 8 M⊙)
young stellar objects. They are the higher mass counterparts of the T Tauri
stars (TTS) (e.g., Brittain et al., 2023). Like TTSs, HAeBe stars also show
H$\alpha$ emission and infrared (IR) excess from a dusty circumstellar
material. In addition to this, they are also observed to emit UV and X-rays
(Zinnecker & Preibisch, 1994; Damiani et al., 1994; Malfait et al., 1998). The
high energy radiation in low-mass main-sequence (MS) stars and pre-main-
sequence (PMS) stars is produced predominantly due to the involvement of the
magnetic field. However, the subsurface convection zone required for the
production of magnetic field (through $\alpha$-$\omega$ dynamo mechanism) is
lacking in HAeBe stars and hence the reason for high energy emission is an
unsettled problem (Appenzeller, 1994; Mayne et al., 2007; Mayne, 2010).

Many studies have been carried out to understand the origin of X-rays from
HAeBe stars. Some of the initial studies to understand the origin of X-ray
emission from HAeBe stars were performed by Zinnecker & Preibisch (1994),
Damiani et al. (1994), Skinner et al. (2004) and Hamaguchi et al. (2005) using
observations from ASCA, ROSAT, and Einstein X-ray Observatory. These studies
could not provide a proper explanation for the origin of X-rays in HAeBe stars
due to the limited resolution of the telescope. The advent of the Chandra
X-ray Observatory (CXO) in 1999 opened a new window (thanks to its high
spatial resolution) for studying X-rays from stars. Stelzer et al. (2006)
studied a sample of 17 HAeBe stars, of which 13 of them were X-ray emitters.
Based on the high plasma temperature obtained from the spectral analysis, this
study ruled out radiative winds as the sole emission mechanism and suggested
that either a huge fraction of HAeBe stars are intrinsic X-ray emitters or
most of the companions to HAeBe stars are yet to be discovered. Hamidouche et
al. (2008) studied a sample of 22 HAeBe stars observed with Chandra and found
X-ray emission from 14 stars, i.e., the sample has a detection rate of
$\sim$64%. They found the X-ray luminosity to be higher than that of TTSs
($10^{30}$ – $10^{31}$ erg s-1). Comparing the X-ray properties to other
stellar parameters they suggested that the X-ray emission might be intrinsic
to HAeBe stars, generated due to star-disk magnetic interaction. Further,
Stelzer et al. (2009a) studied a sample of 9 HAeBe stars using the CXO with a
detection rate of $\sim$100%. Confirming the results obtained from Stelzer et
al. (2006), Stelzer et al. (2009a) suggested that unresolved (multiple stars
not resolved in Chandra) coronally active low-mass companions as the possible
reason for X-ray emission in HAeBe stars. Nuñez et al. (2021) studied a small
sample of radiative intermediate-mass PMS (R-IMPS) stars and AB-type stars and
suggested a similar explanation as Stelzer et al. (2006, 2009a) for the
emission of X-rays in HAeBe stars.

In this work, we present the results obtained by analyzing a sample of HAeBe
stars showing X-ray emission and evaluate the possible mechanisms for X-ray
emission. Though X-ray emission in TTSs and HAeBe stars has been studied
previously, a comparative study with a large sample has not been done,
especially with the data from a telescope such as Chandra which has an
excellent sub-arcsec spatial resolution (Tsunemi et al., 2001). The paper is
structured as follows: the details about the data sample of HAeBe stars and
TTSs considered for this study are described in Section 2. Section 3 describes
the data acquisition process and analysis of the X-ray spectra and
lightcurves. Section 4 provides a detailed description of the observed X-ray
properties of our sample of HAeBe stars and TTSs in the ONC. Further, the
X-ray properties of HAeBe stars and TTSs are compared with their IR and
optical properties. In Section 5, we discuss the role of binaries to
understand the X-ray emission from HAeBe stars and explore the possible X-ray
emission mechanisms in the light of our study. The main results from this
study are concluded in Section 6.

##  2 Sample

###  2.1 Chandra data of HAeBe stars

We searched for the X-ray observations in the "Chandra Data Archive" with a
list of 279 HAeBe stars compiled from Hamaguchi et al. (2005), Stelzer et al.
(2006), Manoj et al. (2006), Hamidouche et al. (2008), Stelzer et al. (2009a),
Chen et al. (2016), Vioque et al. (2018), and identified X-ray observations
for 62 HAeBe stars in the field of view (FOV) of Advanced CCD Imaging
Spectrometer (ACIS) of Chandra. 19 of the 62 sources are presented for the
first time in this paper. They are marked with “*” in Table
LABEL:TableA1:xrayobservations. 7 of these 19 sources show X-ray emission. 28
of the 62 sources have multi-epoch data. In this paper, for spectral analysis,
we consider the observation with the highest exposure and maximum number of
counts for stars having multi-epoch data. We also study the temporal behavior
of the stars with multi-epoch observations. Of the 62 observations
corresponding to 62 sources, 44 are detected in X-rays. Our sample contains a
mixture of single and multiple-star systems. The ACIS images of the 44 sources
can be classified into 3 categories, i.e., resolved HAeBe star systems (11),
unresolved HAeBe star systems (14), and single HAeBe stars (19). We discuss
the multiple-star systems to understand the X-ray emission from HAeBe stars
and rule out the possibility of TTSs being solely responsible for the observed
X-ray emission. Other stellar parameters and the extinction ($A_{\mathrm{V}}$)
along the line of sight are compiled from the literature and tabulated
together in Table LABEL:TableA1:xrayobservations. Table
LABEL:TableA1:xrayobservations also provides details on the X-ray status of
the source, i.e., the sources detected in X-rays are denoted using “X”. The
multiple systems resolved by Chandra are denoted using “X2”. The distances for
all the sources are taken from the GAIA EDR3 distance catalog by Bailer-Jones
et al. (2021).

###  2.2 COUP catalog : T Tauri Stars

The data for TTSs were compiled from the COUP survey program (Getman et al.,
2005). The COUP program is the deepest and longest X-ray observation ever made
of a young stellar cluster providing a rich and unique data set for a wide
range of science studies (e.g., Getman et al., 2005; Preibisch et al., 2005;
Preibisch & Feigelson, 2005; Stelzer et al., 2005). The CXO was pointed at ONC
between 8 January 2003 and 21 January 2003 using the imaging configuration of
the ACIS-I, which gives a FOV of 17’x 17’. The total exposure time of the COUP
image was 838 ks, corresponding approximately to 9.7 days. The spatial
resolution of ACIS is better than 1" over most of the FOV. The very low
background allows the reliable detection of sources with as little as $\sim$5
source counts. The final COUP source catalog lists 1616 individual X-ray
sources. The ACIS-I chip has a very good point-spread function and a high
accuracy of the aspect solution which allows for a clear and unambiguous
identification of nearly all X-ray sources with optical or near-infrared
counterparts (Preibisch et al., 2005). We cross-matched the sources from the
COUP catalog (Getman et al., 2005), whose X-ray luminosities are available,
with the optical catalog from Da Rio et al. (2012) and obtained around 900
sources. The cross-matched list was further filtered to contain only those
sources having a membership probability $\geq$ 50%. These sources are
considered to be the bona fide members of ONC, thereby reducing the sample
size to 716. The number of known populations of stellar cluster members with
precisely determined properties significantly increased in this study. So, the
number of X-ray emitting TTSs in the sample considered here is more than those
listed in Preibisch et al. (2005).

##  3 X-ray data REDUCTION AND ANALYSIS

The data reduction was carried out using the CIAO (Fruscione et al., 2006)
software package version 4.13, in combination with the calibration database
(CALDB)111https://cxc.cfa.harvard.edu/caldb/downloads/index.html 4.9.4. We
followed the standard data analysis procedure recommended by the Chandra X-ray
Center (CXC)222 https://cxc.cfa.harvard.edu/ciao/guides/. The raw observation
files from the archive are reprocessed to obtain the new level-2 event file,
which applies the latest calibration and cleans the event file for good time
intervals (GTI), removing any spurious or bright pixels (cosmic ray). We then
apply an energy filter of 0.3 $-$ 8 keV to this event file, which is the final
event file used for all the scientific analyses. Further, we check for any
flaring events in the background lightcurve of the observations. If present
then, we follow the standard procedure as recommended by the CXC2 to remove
the time bins during which the flaring occurred. This is performed to keep the
data uncontaminated due to the background flares. We used the source detection
algorithm “`wavdetect`” (Freeman et al., 2002), as it is well suited for
separating closely spaced point sources. The wavelet scales are set between 1
and 16 in steps of $2^{n}$ (where, n = 0,1,2,3,4), and a detection
significance of $10^{-6}$ is set to avoid spurious or false detections. This
threshold is lowered to $10^{-4}$ for the stars with the closest companions
(Stelzer et al., 2009a). The position (RA & Dec) of the sources was then set
to the value obtained from the “`wavdetect`” algorithm. We then extract the
lightcurves and spectra from the event file. For the extraction of the
lightcurves and the spectra, we chose a circular region for both the point
source and the background. The size of the source region is set to the radius
that includes 90% of the point spread function (PSF) at 1.5 keV obtained from
the `wavdetect` algorithm, and for the background, we choose a source-free
region within the same CCD chip. For the sources that are partially resolved
such as in Figure 6, the extraction regions of the primary and the secondary
source were carefully chosen to not overlap with one another using appropriate
PSF fraction. In cases where the extraction regions of the primary and the
secondary source overlap the common region between them was excluded to
minimize the source contamination (Skinner & Güdel, 2014). This contamination
was not significant in any of the cases.

The targets AK Sco, BD+30 549, BP PSc, HD 100546, R CrA, HBC 442, and HD 36982
are very faint and have been observed at different epochs. To obtain their
spectrum with a better SNR, we co-added the spectra observed at different
epochs, using the script “`combine_spectra`” (Table
LABEL:TableA1:xrayobservations). This was performed only if the sources did
not show variability at the different epochs. We checked the variability for
these 7 sources by calculating the hardness ratio (HR) at different epochs
using Equation (1). More details with respect to HR are provided in Section
3.1. The sample size in our work is quite large (44 X-ray emitting sources)
compared to the previous studies. The X-ray parameters of 37 out of 44 HAeBe
stars in our sample have previously been studied in the literature. Further,
we found new X-ray detections for 7 HAeBe stars (See Table
LABEL:TableA1:xrayobservations). In order to maintain homogeneity with respect
to the data analysis, we re-analyzed the X-ray parameters for all the sources
in our sample. The X-ray parameters obtained from our analysis agree well with
the literature.

###  3.1 Temporal Analysis: X-ray Lightcurves

The 0.3 $-$ 8 keV X-ray lightcurves are extracted for the sources with
multiple observations. Only the lightcurves with good time resolution and
statistics are used for the scientific analysis. We binned the lightcurves of
all the detected sources to a fixed bin time. They are binned carefully to not
contain any empty bins. The bin time used for the lightcurve of each source is
mentioned in Figures 7 and 8. The HEASARC ftool “`lcstats`” is used to measure
the “rms fractional variation” to check the variability of the lightcurve. The
HR for all the X-ray emitting sources in the energy range 0.3 $-$ 1 keV (S)
and 1 $-$ 8 keV (H) was calculated using equation (1).

$HR=\frac{H-S}{H+S}$ (1)

The following 9 sources, V892 Tau, HD 97300, TY CrA, HD 37062, V372 Ori, V380
Ori, HR 5999, HD 36939 and [DLM2010] EC 95 show flare-like features (Giardino
et al., 2004; Hamaguchi et al., 2005). We observe a rapid rise in the count
rate of these sources. To confirm this we calculated the HRs and observed a
clear hardening. The HR of these stars increases with the increase in the
count rate which is not observed for stars without variable/flaring
lightcurves. The sources HD 97300 and EC 95 are observed only once while the
rest of the 7 sources have multiple observations obtained over different
epochs. HR 5999 has 7 multi-epoch observations and it is resolved only in the
Obs ID: 8901, where HR 5999 was the prime target of the observation. HR 5999 A
is very faint in X-rays, so its lightcurve could not be extracted. Similarly,
the star V892 Tau has 2 multi-epoch observations and is resolved only in Obs
ID: 3364. The lightcurves showing flaring/variability are presented in
Appendix D. More details on the lightcurves are discussed in detail in Section
5.1.

###  3.2 Spectral Analysis: The X-ray Spectra

The X-ray spectra of HAeBe stars were extracted using the CIAO tool
“`specextract`”. The Spectrum of each source was binned to a minimum of 5 or
more counts/bin. The grouping was chosen based on the quality of the spectrum
(i.e., the total number of counts). The sources having total counts < 100 are
binned to contain a minimum of 5 counts/bin. For X-ray sources that have the
total number of counts $\geq$ 100, we binned the spectrum to contain a minimum
of 15 counts/bin.

The spectral fitting was performed using the tool XSPEC version 12.11.1d
(Arnaud, 1996) available with HEASOFT (version 6.30). We used the plasma
emission model “APEC” (Smith et al., 2001) along with “TBABS” for absorption
correction. The solar abundances were set to values from Grevesse & Sauval
(1998) for the plasma model and Wilms et al. (2000) for the absorption model.
We tried fitting the spectrum with single temperature (1T) and two temperature
(2T) models, considering four different model combinations on each spectrum
i.e., 1T and 2T models with fixed or variable absorption column (NH) from
interstellar dust and the circumstellar material. For fits with fixed
absorption, the NH of the “TBABS” component was set to the values derived from
the AV, using the following equation taken from Bohlin et al. (1978),

$\frac{N_{H}}{E(B-V)}=5.8\times 10^{21}cm^{-2}mag^{-1}$ (2)

Since our work deals with HAeBe stars, we calculated the NH considering an
extinction factor RV=5. A larger RV is used as the circumstellar disk of HAeBe
stars is dominated by grains larger in size than the average dust grain in the
diffuse interstellar medium (Hernández et al., 2004; Manoj et al., 2006). The
column density of the close companions is assumed to be equal to that of the
HAeBe stars (Stelzer et al., 2006). Global metal abundances were fixed to 0.3
times the solar value (see, Imanishi et al., 2003; Feigelson et al., 2002, for
details). For each source, the initial kT and the emission measure (EM) values
for model fitting were taken from the literature.

To obtain the values of unabsorbed flux we use the convolution model "`cflux`”
in the energy range 0.3 $-$ 8 keV. We tried the model fitting with initial
values of model parameters from the literature for four different model
combinations (fixed and free NH for 1T and 2T APEC models) for each source and
adopted the fit with the minimum $\chi^{2}$. For $\chi^{2}$ statistics the fit
quality is determined by null-hypothesis probability (PNULL). We evaluated the
quality of the fit based on the PNULL to check how well the spectrum is
described by the model and accept only those spectral fits whose PNULL > 5%.
First, we try the spectral fitting with the 1T model by fixing NH to a value
obtained from equation 2. We consider this model as the best fit if PNULL >
5%, else 1T model with free NH is considered. If none of the 1T models gives a
best fit, then the 2T model with fixed NH is considered. Finally, if none of
the models provide the best fit then we use the 2T APEC model with free NH.
The above method described is used only for $\chi^{2}$ statistics which works
best only for the bright sources (>100 counts) i.e., in our case, the stellar
spectrum with at least 15 counts/bin, as this follows a Gaussian distribution
(Bevington & Robinson, 2003). The faint sources (<100 counts), i.e., stellar
spectrum with less than 10 counts/bin are best explained by a simple
Poissonian distribution. Here we used cash statistics (c-statistics) as the
test statistic for such data. The quality of the fit for the c-statistics is
determined by the “`goodness`” command which plots a histogram of probability
versus $\chi^{2}$. If the statistic of observed data exceeds 90% of the
simulated statistic values, then the model is rejected at 90%
confidence333https://heasarc.gsfc.nasa.gov/docs/xanadu/xspec/XspecManual.pdf.
The unabsorbed X-ray flux (FX) is obtained from the model fitting and the
unabsorbed X-ray luminosity (LX) for all the sources is calculated using the
formula,

$\centering L_{X}={4\pi d^{2}F_{X}}\@add@centering$ (3)

where d is the distance to the source. All the X-ray parameters in the model
were obtained with a 90 percent confidence range. The X-ray parameters of the
sources obtained by the above-mentioned methods agree well with previous
studies (Stelzer et al., 2006, 2009a; Hamidouche et al., 2008).

The spectrum of sources MWC 297 A, MWC 297 B, HR 5999 A, T CrA and R Mon
having very few counts ($\leq$10 counts) could not be fit with a model. We
obtained the count rate of these sources using the CIAO tool “`aprates`”. The
LX of these sources was then obtained using “PIMMS”, assuming an optically
thin plasma of 10 MK ($\sim$1 keV), abundance of $\sim$ 0.3 and the NH derived
from the AV. The same procedure was used to estimate the upper limits of the
LX for the non-detected sources. The LX of the source R CrA (counts > 100) was
also obtained using PIMMS as the model parameters could not be constrained.
Table LABEL:TableB1:xrayparam gives the X-ray parameters of stars fit with 1T
and 2T APEC models for HAeBe sources and their companions. Table
LABEL:TableB2:xraydark gives the parameters for the sources that could not be
detected in X-rays by Chandra.

##  4 RESULTS

Among our sample of 62 HAeBe stars, 44 of them show X-ray emission. Chandra
has resolved 11 HAeBe stars among the 44, where the resolved companion is at a
position >1". Of the 11 resolved systems, there are 5 stars with sub-arcsec
companions. Apart from the resolved systems, there are 14 HAeBe stars with
unresolved sub-arcsec companions i.e., companion stars within 1" of the
primary which Chandra could not resolve. So, in total there are 19 out of 44
unresolved sub-arcsec multiple-star systems where i.) the primary and the
secondary are HAeBe stars (e.g., AK Sco, V892 Tau), ii.) the primary and
secondary are HAeBe stars with a TTS companion (e.g., R CrA, V380 Ori), and
iii.) the primary is a HAeBe star and its companions are TTSs (e.g., TY Cra,
HD 97300). In this paper, the group of 19 stars is identified as non-single
stars. We evaluated a sub-arcsec companion (stars with separation < 1")
fraction of 43.1$\%$ for the stars detected in X-rays in this paper. 25 out of
44 stars (56.8$\%$) are either confirmed to have no stellar companions or have
any evidence of the presence of a companion as of yet (we name this group as
‘single stars’). The details of the companion status of the HAeBe stars in our
sample are given in Table LABEL:TableC1:individualstars.

The X-ray parameters (kT, LX and NH) of HAeBe stars in our sample obtained
from X-ray spectral analysis are tabulated in Table LABEL:TableB1:xrayparam.
To understand the origin of the X-ray emission in HAeBe stars, we compare
their X-ray parameters with other stellar parameters such as bolometric
luminosity, temperature, infrared excess, and accretion. Further, we compare
the properties of HAeBe stars with those of TTSs. The parameters of TTSs in
the ONC are obtained from the COUP analysis (Getman et al., 2005), as
discussed in Section 2.2.

Figure 1: X-ray luminosity (LX) against bolometric luminosity (Lbol). Grey
dots $-$ TTSs (ONC), blue stars $-$ HAeBe stars(our sample), orange stars $-$
close companions to HAeBe stars (our sample), faint blue stars $-$ upper
limits (X-ray dark sources). Figure 2: Variation of fractional X-ray
luminosity against the temperature and spectral type. The dashed line denotes
the mean LX/Lbol of TTSs (10-3.6) and the dotted line that of main-sequence OB
type stars (10-7). Figure 3: (a). kT vs LX plot for stars fitted with 1T APEC
model. The light grey “+” symbol denotes TTSs fitted with 1T model and the
light blue “filled star” symbol denotes HAeBe stars fitted with 1T APEC model.
The orange “filled star” symbol denotes the resolved companions of HAeBe
stars, (b). kT vs LX plot for stars fitted with 2T APEC model. The grey “+”
and “o” symbols denote kT1 and kT2 of TTSs, respectively whereas dark blue
“filled star” and “filled cross” symbols denote kT1 and kT2 of HAeBe stars,
respectively. The red “filled star” and “filled cross” symbols denote kT1 and
kT2 of the resolved companions of the HAeBe stars. The blue and red lines
connecting these symbols represent the plasma temperatures of their respective
stars; The stars outlined in red represent the HAeBe stars with unresolved
companions. The red dashed line is marked at 2 keV denoting the plasma
temperature within which single HAeBe stars are found.

###  4.1 Estimation of LX/Lbol of HAeBe stars and their variation with
spectral type

The variation of X-ray luminosity with bolometric luminosity (Lbol) for HAeBe
stars (our sample) and the TTSs from the COUP sample is shown in Figure 1. We
compute the bolometric luminosity of HAeBe stars as follows. The observed
apparent magnitude (mV) is corrected for extinction (AV) and using the
distance modulus relation we calculated the absolute magnitude (MV). The
distances to the sources used for the calculation are mentioned in Table
LABEL:TableA1:xrayobservations. The mV for HAeBe stars were obtained from
Manoj et al. (2006), Mathew et al. (2018), and Vioque et al. (2018). From MV
we estimated the bolometric magnitude (Mbol) using the values of bolometric
correction corresponding to each spectral type, from the table listed in
Pecaut & Mamajek (2013). We then calculate the bolometric luminosity (Lbol) of
the HAeBe stars using the equation below,

$\log\left({\frac{L_{bol}}{L_{\odot}}}\right)=\left({\frac{M_{bol\odot}-M_{bol}}{2.5}}\right)$
(4)

where, the bolometric magnitude of the Sun Mbol,⊙ = 4.75. The bolometric
luminosity and the unabsorbed X-ray luminosity for the TTSs are obtained from
Getman et al. (2005).

From Figure 1, we see that the Lbol-LX distribution of HAeBe stars and TTSs
follow a similar trend, with TTSs being more steeper. The X-ray luminosities
of both the stellar groups are slightly different, ranging from around
$10^{28}$ to $10^{33}$ erg s-1 for HAeBe stars and $10^{27}$ to $10^{32}$ erg
s-1 for TTSs. We see there is a necessity to normalize the LX, as this gives
us an estimate of how bright or faint the HAeBe stars are in X-rays compared
to TTSs. Hence, we now consider the plot of LX/Lbol against effective
temperature/spectral type (Figure 2). The effective temperature (Teff) for
HAeBe stars were obtained from Manoj et al. (2006), Hamidouche et al. (2008)
and Vioque et al. (2018). For TTSs, we obtained the Teff by cross-matching the
final list of TTSs (see, Section 2.2) with Da Rio et al. (2012). The Teff was
converted to spectral type using the table listed in Pecaut & Mamajek (2013).

We see from Figure 1 that, although the LX of HAeBe stars and TTSs are not
substantially different, their LX/Lbol appear to greatly differ. It is
apparent from Figure 2 that the LX/Lbol of HAeBe stars is lower by about 3 to
4 orders of magnitude when compared to that of TTSs. In order to provide a
visual comparison in Figure 2, we have marked the average LX/Lbol of TTSs
(10-3.6) (Preibisch et al., 2005) with a horizontal dashed line and that of MS
OB stars ($10^{-7}$) (Nazé et al., 2011) with a horizontal dotted line. The
LX/Lbol of most of the HAeBe stars in our sample (86%) fall in between the
mean value of TTSs and OB stars. The LX/Lbol of HAeBe stars decrease with
increasing Teff, i.e., massive HAeBe stars emit lesser X-ray flux. We do not
observe such a trend for TTSs.

###  4.2 Variation of LX with kT

The X-ray spectra of PMS stars could fit well with a 1T or 2T absorbed thermal
plasma model (Getman et al., 2005). The plasma temperatures (kT) for the HAeBe
stars in our sample were estimated as discussed in Section 3.2 and it agrees
well with the literature (Hamaguchi et al., 2005; Stelzer et al., 2006,
2009a). While the coronal temperature is unlikely to be isothermal or be
approximated by 1T or 2T models, we still follow this approach as our aim is
to homogeneously characterize the plasma temperatures of HAeBe stars in our
sample. The plasma temperatures obtained from the model fitting of the X-ray
spectrum can be considered as a characteristic temperature (Peres et al.,
2004; Preibisch et al., 2005) and hence we use this parameter to see its
variation with LX. Figure 3 shows the kT versus LX for HAeBe stars (our
sample) and TTSs (COUP sample). Figure 3(a) represents the HAeBe stars and
TTSs that fit with the 1T APEC model and Figure 3(b) represents the HAeBe
stars and TTSs that fit with the 2T APEC model. In the case of TTSs fitted
with 1T model (Figure 3a), the plasma temperature of a majority of them is on
the hotter side i.e., kT ranging from 0.18 to 9.5 keV. However, in the case of
TTSs with 2T model fits (Figure 3b), we observe a cool component (kT1) ranging
from 0.14 to 1.37 keV and a hot component (kT2) ranging from 0.72 to 9.23 keV
(Preibisch et al., 2005). The cool component seems to be constant at $\sim$10
MK (kT $\sim$ 1 keV), while the hot component seems to increase with
increasing LX. An increase in kT with increasing LX is also observed for the
TTSs with 1T fit (Figure 3a). HAeBe stars with 1T fit which have no known
companions have plasma temperatures in the range 0.21 to 1.98 keV, and HAeBe
stars with sub-arcsec companions (stars outlined in red in Figure 3a) have
plasma temperatures in the range from 0.65 to 3.05 keV. In the case of HAeBe
stars with 2T fit having no known companions, the plasma temperature kT1
ranges from 0.12 to 0.20 keV and kT2 from 0.58 to 0.74 keV, whereas for HAeBe
stars with sub-arcsec companions (stars outlined in red in Figure 3b), the
plasma temperature kT1 ranges from 0.34 to 1.41 keV and kT2 from 1.29 to 8.25
keV. The sub-arcsec companions to HAeBe stars vary from being low mass to
similar mass stars (Table LABEL:TableC1:individualstars). We see hotter plasma
temperatures (kT > 2 keV) for TTSs than that for single HAeBe stars, which are
relatively cooler i.e., kT $\leq$ 2 keV (See Table LABEL:TableB1:xrayparam).
The plasma temperatures of HAeBe stars with sub-arcsec companions are hotter
(kT > 2 keV), similar to TTSs. About 50% of the non-single HAeBe stars fit
with the 1T model and 80% non-single HAeBe stars fit with the 2T model have kT
> 2 keV.

###  4.3 Estimation of IR spectral index and its variation with LX/Lbol

It is understood from the infrared studies of HAeBe stars (Manoj et al., 2006;
Chen et al., 2016) that the IR flux excess is due to the presence of dust in
their circumstellar disk. Lada (1987) showed that the IR spectral index is one
of the important criteria used to identify the presence of the disk in YSOs.
In this work, we used the IR spectral index $\eta_{(2-4.6)}$ as an indicator
of IR excess. A classification scheme for YSOs was quantified by Greene et al.
(1994) using the slope of the SEDs in the IR region. The YSOs are distributed
into 4 classes, described as Class I ($\eta$ $\geq$ 0.3), Flat (-0.3 $\leq$
$\eta$ < 0.3), Class II (-1.6 $\leq$ $\eta$ < -0.3), Class III ($\eta$ $\leq$
-1.6).

Figure 4: IR Spectral Index against LX/Lbol showing the classification scheme
by Greene et al. (1994). Class I ($\eta$ $\geq$ 0.3), Flat (-0.3 $\leq$ $\eta$
< 0.3), Class II (-1.6 $\leq$ $\eta$ < -0.3), Class III ($\eta$ $\leq$ -1.6).

We obtained the spectral indices in the KS-W2 IR band for our sample using the
method described in Manoj et al. (2011) and Arun et al. (2019). The spectral
index was calculated using the following equation,

$\eta_{\left(\lambda_{1}-\lambda_{2}\right)}=\frac{\log\left({\frac{\lambda_{2}F_{\lambda_{2}}}{\lambda_{1}F_{\lambda_{1}}}}\right)}{\log\left({\frac{\lambda_{2}}{\lambda{1}}}\right)}$
(5)

where, $\lambda_{1}$,$\lambda_{2}$ are the two wavelengths considered, and
F${}_{\lambda_{1}}$, F${}_{\lambda_{2}}$ are the flux values at
$\lambda_{1}$,$\lambda_{2}$.

Figure 4 shows the plot of $\eta_{(2-4.6)}$, an indicator of infrared excess
between KS and W2 bands, against the LX/Lbol for HAeBe stars and TTSs. The
excess NIR emission in HAeBe stars is from their inner disk (Dullemond &
Dominik, 2004a, b, 2005). Most of the stars in our sample fall under Class II
and flat group, which indicates that a majority of the stars in our sample
possess a circumstellar disk. We see in Figure 4 that similar to TTSs, HAeBe
stars do not follow a relation between LX/Lbol and $\eta_{(2-4.6)}$. We see
that the LX/Lbol of both samples does not seem to be affected by the presence
or absence of the disk. Further, we need to investigate if the accretion of
matter from the circumstellar disk to the star is responsible for the X-ray
emission in HAeBe stars, which is evaluated in Section 4.4.

###  4.4 Estimation of LHα/Lbol and its variation with LX/Lbol

We know that TTSs accumulate mass through the magnetospheric accretion
process. According to this paradigm, the magnetic field lines from the star
truncate the disk, and the material from the disk is transferred onto the
star. When the in-falling material shocks the photosphere, X-rays are emitted
(Calvet & Gullbring, 1998). It is known that the H$\alpha$ emission line is a
good tracer of accretion in PMS stars. Hence, to investigate if the accretion
mechanism in HAeBe stars is contributing to the emission of X-rays, we
calculate the H$\alpha$ luminosity (LHα) and plot the LHα/Lbol against LX/Lbol
(Figure 5).

Figure 5: LHα/Lbol of HAeBe stars (Blue star) and TTSs (grey dot) plotted
against the LX/Lbol. The Pearson’s coefficient for HAeBe stars and TTSs is r =
-0.09; p-value = .00051 and r = 0.23; p-value = 0.60167, respectively.

We calculated the H$\alpha$ flux using the method described in Mathew et al.
(2018). They used the extinction corrected R-band flux density as a proxy for
the continuum flux density underlying the H$\alpha$ line, from which the
continuum H$\alpha$ flux (Fν,cont(Hα)) is obtained using the equation 6,

$F_{\mathrm{{\nu,cont}(H\alpha)}}=F_{\mathrm{{\nu,0}}}\times
10^{\mathrm{{\left(\frac{-R_{0}}{2.5}\right)}}}$ (6)

where Fν,0 = 3.08 $\times$ 10-23 W m-2 Hz-1 and R0 is extinction corrected
magnitude. The extinction in R-band (AR) was estimated from AV using the
extinction curve of McClure (2009). From the Fν,cont(Hα), we estimated the
H$\alpha$ line flux using Fλ,line(Hα) = Fν,cont(Hα) $\times$
EW$\mathrm{(H\alpha)}$. The absorption corrected equivalent width
(EW$\mathrm{(H\alpha)}$) for HAeBe stars was obtained from the literature
(Manoj et al., 2006; Fairlamb et al., 2017; Vioque et al., 2018; Wichittanakom
et al., 2020). The H$\alpha$ luminosity (LHα) is then calculated using the
standard flux-luminosity relation, similar to Equation 3. Figure 5 shows the
plot of LHα/Lbol against LX/Lbol for both HAeBe stars and TTSs. We perform
Pearson’s correlation test to check for a linear relation between X-ray
emission and the accretion processes. For TTSs, we obtain Pearson’s co-
efficient r = 0.23 with a p-value = 0.0005 and for HAeBe stars we obtain r =
-0.09 with a p-value = 0.6017 at a significance ($\alpha$) level of 0.05.
These values suggest that, for TTSs and HAeBe stars, we do not observe any
correlation between LHα/Lbol and LX/Lbol. This strongly suggests that the
accretion process is not a major contributor to the production of X-rays in
TTSs and HAeBe stars.

(a)  (b)  (c)  (d)

Figure 6: 20" x 20" ACIS $-$ CCD images of the HAeBe stars having equal mass
companions (possible HAeBe stars). The images are binned to a pixel size of
0.25". the green circle denotes the photon extraction region of the primary
star and the red circle denotes the photon extraction region of the secondary
star. For the sources, BD+65 1638 A & B and V892 Tau B, the radius of
extraction regions was set to a PSF fraction of 0.84.

##  5 Discussion

Quite a few mechanisms have been proposed to explain the X-ray emission from
HAeBe stars, with the prominent ones being due to the contribution from a low-
mass (undetected) companion in the close proximity of the HAeBe star (Stelzer
et al., 2006, 2009a), colliding winds producing X-ray emission (similar to
what is seen in massive OB-type stars: Zinnecker & Preibisch, 1994; Damiani et
al., 1994), presence of a shear-powered dynamo in late-type HAeBe stars (Tout
& Pringle, 1995), the magnetically confined wind shock model (MCWS) (Swartz et
al., 2005; Telleschi et al., 2007b), and magnetospheric accretion shock model
(Calvet & Gullbring, 1998). From the results obtained in this paper, we
propose that the X-ray emission from HAeBe stars has a different origin than
TTSs. In the following sub-sections, we rule out the possibility of the
presence of an unseen low-mass companion as solely the origin of X-rays and
discuss the possible X-ray mechanisms that contribute to the emission of
X-rays in HAeBe stars.

###  5.1 Role of binaries

In this section, we discuss the details of the various types of companions
mentioned in Section 4 and present three observational details that support
intrinsic X-ray emission from HAeBe stars.

  1. 1.

Planetary studies – The presence of planetary companions around the X-ray
emitting HAeBe stars: Some stars in our sample were studied individually using
high-resolution instruments and adaptive optics. There were no strong
signatures of stellar companions, though planetary companions were discovered.
Among the 25 single HAeBe stars, 7 such systems are present. For example, the
HAe star AB Aurigae has been studied extensively (e.g., Boccaletti et al.,
2020; Currie et al., 2022; Zhou et al., 2022). Fukagawa et al. (2004) studied
AB Aur in the NIR region using the Coronographic Imager Adaptive Optics
systems on the Subaru telescope, in which spiral features within the
circumstellar matter were discovered. Tang et al. (2012) suggested the
presence of an unseen companion to AB Aur at 30 AU as the probable reason for
the observed spiral patterns in the disk. A follow-up paper by Tang et al.
(2017) identified two spirals in high-resolution ALMA 12 CO $J$ = 2 $-$ 1
emission observations and suggested that the spiral arms are due to several
protoplanets present at 60 $-$ 80 AU (0".4 $-$ 0".6 separation) and another
companion at 30 AU (0".16 $-$ 0".18 separation). Boccaletti et al. (2020)
presented VLT/SPHERE observations of AB Aur and suggested that the two spirals
observed in the system are due to the protoplanetary candidates at around 30
AU ($\sim$14$M_{\mathrm{Jup}}$) and 110 AU ($\sim$3$M_{\mathrm{Jup}}$) from
the center of the star. Similarly, the Herbig Be star HD 100546 was studied by
Pérez et al. (2020) using the Atacama Large Millimeter/submillimeter Array.
They found evidence for the signatures of 2 protoplanets
($\sim$3.1$M_{\mathrm{Jup}}$ & 8.5$M_{\mathrm{Jup}}$; Fedele et al., 2021) and
ruled out the possibility of stellar companions by their sparse-aperture
masking data. HD 163296 was studied by Kirwan et al. (2022) using spectro-
images and position velocity diagrams extracted from the MUSE data cube. They
ruled out the presence of any stellar companions within 1 AU $-$ 35 AU and
suggested the presence of either a brown dwarf or a massive planet. The
details of other stars supporting the evidence of protoplanetary candidates in
their inner disk regions are tabulated in Table LABEL:TableC1:individualstars
with references. These studies strongly rule out the presence of a stellar
companion. Also, the type of X-ray emission observed from the group of stars
identified as single stars is towards the cooler side (kT $\leq$ 2 keV), as
discussed in Section 4.2.

  2. 2.

Similar mass companions to the X-ray emitting HAeBe stars: In our sample of
stars, there are 4 stars that possess a similar mass/HAeBe star as companions,
with no other low-mass companion reported (see, Table
LABEL:TableC1:individualstars). Figure 6 shows the 20" $\times$ 20" ACIS
images binned to a pixel size of 0.25" of the HAeBe stars having companions
that are possibly HAeBe stars. The details of these 4 stars are discussed
below.
AK Sco: AK Sco (HD 152404) is a spectroscopic binary (SB2; Andersen et al.,
1989) system consisting of two equal-mass (total M⋆=2.49M⊙) F5-type stars,
with a period of 13.6 days and a large eccentricity (e = 0.47). The distance
of separation between the two stellar components is $\sim$ 0.16 AU and they
come as close as 11R⋆ at the periastron. From the high-resolution studies
(Esposito et al., 2020) there has been no detection of another low-mass star
in this system. The system has been observed thrice using CXO. Due to its very
low quality, we combined the data to obtain a good-quality spectrum (see,
Section 3.2). The details of the X-ray parameters are given in Table
LABEL:TableB1:xrayparam. We see a very cool plasma temperature (kT$\sim$0.6
keV) and no variability in the lightcurve.
BD+65 1638: BD+65 1638 is one of the three massive HBe stars in the reflection
nebula NGC 7129. It has a spectral type of B2-B3 (Racine, 1968; Matthews et
al., 2003). Dahm & Hillenbrand (2015) first suggested for the star to be
spectroscopic binary. From the source’s L⋆ and Teff, the star is placed very
close to the birth-line of a 6M⊙ object (Matthews et al., 2003). In Chandra
ACIS data the star is resolved into 2 X-ray bright sources with a separation
of $\sim$1.3" (Figure 6(b)). This system is one of the bright HBe stars in our
sample. The X-ray spectra of both the stars fit well with 1T APEC plasma model
with plasma temperature kT = 1.98 keV for BD+65 1638 A and kT = 1.13 keV for
BD+65 1638 B. The plasma temperature of component A is slightly higher than
that of component B. The LX for both the stars is within the range for B-type
stars. We do not observe any variability/flaring in the lightcurve of both the
sources. The stellar system seems to be unusually in a quiescent state, as the
star has just emerged out of the gaseous envelope.
PDS 69: PDS 69/H4636 is associated with the reflection nebula NGC 5367. This
is a double-star consisting of two B-type stars of spectral types B7 (PDS 69
A) and B4 (PDS 69 B). The system is a wide binary at an angular separation of
$\sim$3.7". Both the stars show IR excess and the northern component i.e., PDS
69 B shows H$\alpha$ in emission (Williams et al., 1977; Maheswar et al.,
2002). From their positions in the CMD, PDS 69 A and PDS 69 B appear to have
masses of about 4.5 and 8 M⊙, respectively (Williams et al., 1977). The wide
binary is well resolved in X-rays (Chandra data). The ACIS CCD image is shown
in figure 6 (c). The X-ray spectrum is fitted with 1T APEC model for PDS 69 A,
with plasma temperature kT = 1.56 keV and LX=1.58$\times$1030 erg s-1. As PDS
69 B is very faint, we used PIMMS to obtain the LX=1.17$\times$1030 erg s-1.
The X-ray lightcurve of PDS 69 A does not show any signs of
variability/flaring or any hard X-ray emission. This double-star system is one
example of two resolved HAeBe stars emitting X-rays with no known third
unresolved companion.
V892 Tau: V892 Tau (Elias 3-1) is a late HAeBe star in the Taurus-Aurigae
star-forming region (Smith et al., 2005). It has a spectral type ranging from
B9 to A6 (Alonso-Albi et al., 2009; van der Marel et al., 2016). It is a
triple star system, consisting of a central binary of equal mass ($\sim$6M⊙,)
separated by $\sim$ 7AU and an M5-type TTS at $\sim$500 AU ($\sim$ 4") from
the central binary (Long et al., 2021). We denote the central binary as V892
Tau A and the M-type companion as V892 Tau B. The star has been studied
previously in X-rays by Stelzer et al. (2006) and Giardino et al. (2004). We
analyzed the Chandra data available for 2 epochs. The system could be resolved
from the M-type companion only in one of the observations (Obs ID: 3364) as
the coordinates of V892 Tau was near the aim-point of the observation (Figure
6 (d)). In our analysis, the Chandra X-ray spectrum was fit with a 2T APEC
model (kT1 = 1 keV & kT2 = 2 keV). A flare is observed in the $\sim$ 18ks
X-ray lightcurve, from the region of the central binary, and no
variability/flare in the position of the M-type low-mass companion was seen.
We also analyzed the observation (Obs ID: 2522) where the source was not
resolved, and the X-ray lightcurve showed some variability with a count rate
similar to that of the central binary in Obs ID: 3364 (see, Appendix 7). The
origin of the flare seems to be coming from the two nearly equal mass HAeBe
stars (Giardino et al., 2004). Similar characteristics are observed in V372
Ori, which is a spectroscopic binary with an A-type PMS companion. However,
many details are not known about this star and it is a good candidate for
future studies.

  3. 3.

HAeBe stars which present flaring/variability: Appendix D shows the
lightcurves of the 9 HAeBe stars stellar systems unresolved with Chandra. Of
the 44 stars detected in X-rays, we see X-ray lightcurve variability (short
term and long term) and flaring in 9 stars (see, Section 3.1; Appendix D). 6
of these have low-mass companions (TY CrA, V380 Ori, HR 5999, HD 97300,
[DLM2010] EC 95 and HD 37062), 2 have a similar mass companion (V892 Tau, V372
Ori) and 1 (HD 36939) have no companions as of now. X-ray flares in low-mass
PMS like TTSs are very common and are characterized by rapid enhancement in
the count rate which is followed by slow decay Caramazza et al. (2007). Low-
mass PMS stars show thermal emission with a kT $\sim$ 1 – 8 keV and a strong
variability with occasional rapid flares consistent with the scenario of
enhanced solar-type activity, attributable to magnetic dynamo processes
(Imanishi et al., 2003; Preibisch et al., 2005). Such characteristics are
observed in the HAeBe stars with low-mass companions. TY CrA shows a large
flare lasting for more than 10ks. While V380 Ori, HR 5999, HD 97300, and HD
37062 also show flare-like properties, their observation does not contain the
details of the whole flare (rise and decay). [DLM2010] EC 95, a proto-Herbig
star (Dzib et al., 2010) shows short term variability. The plasma temperature
of these stars is very high, on the hotter side ranging from kT = 2.75 $-$
8.25 keV (Figure 3a & b). V892 Tau, and V372 Ori are the only two HAeBe stars
known to show flaring activity, with a similar flare shape (Figure 7). This is
unusual for HAeBe stars as they do not possess convective layers like low-mass
stars. The plasma temperature of V892 Tau is very close to what we see for
single HAeBe stars in this paper (kT $\sim$ 2 keV; Section 4.2). V372 Ori has
a plasma temperature (kT1 = 0.94 keV; kT2 = 3.75 keV) and flare
characteristics as seen in low-mass stars. It is identified as a spectroscopic
binary (Abt et al., 1991) consisting of a B9.5+A0 stellar system. The
lightcurve of HD 36939 has a very low count-rate ($\sim$0.002 cts/s) compared
to other stars, but we see variability in the count-rate in only one (Obs ID:
4395) of the 14 observations (spanned over a period of 13 years). More high-
resolution studies are required to understand the properties of these stellar
systems.

The above-mentioned are details of types of sub-arcsec companions to HAeBe
stars in our sample (i.e., planetary, similar mass and low-mass companions).
The temperature structure of single HAeBe stars and HAeBe stars with
companions seems to be different. We see that single HAeBe have cooler plasma
temperatures compared to the non-single stars. From the discussion, we can see
that in HAeBe stars, X-ray emission is not confined to one mechanism. In
Section 5.3, we discuss the possible emission mechanisms that can explain the
origin of X-ray emission in HAeBe stars.

###  5.2 Magnetic fields in Herbig Ae/Be stars

Magnetic field has been detected in very few HAeBe stars. Hubrig et al. (2009)
studied 21 HAeBe stars using FORS1 mounted on the 8m Kueyen telescope of the
VLT. They detected magnetic fields for all of the sources in their sample with
a confidence level of $\sim$1$\sigma$. Among the 21 HAeBe stars, 10 were
detected with a significance level > 3$\sigma$. The magnetic fields varied
from less than 10 G to $\sim$500 G. On cross-matching the X-ray emitting stars
in our sample with Hubrig et al. (2009), we find 11 stars (HD 100453, HD
100546, HD 135344B, HD 144432, HD 144668, HD 150193, AK Sco, HD 163296, HD
169142, HD 190073, HD 97300) to possess magnetic field (see, Table 3 Hubrig et
al., 2009). Alecian et al. (2009) studied the multiple-star system V380 Ori
using high-resolution spectropolarimeters ESPaDOnS, installed on the 3.6 m
CFHT, and Narval, installed on the 1.9 m TBL. They measured a magnetic field
strength of $\sim$500 G for the primary. No magnetic field was detected for
the secondary component. Hubrig et al. (2011) studied HD 31648 using high-
resolution polarimetric spectra obtained with SOFIN spectrograph installed at
Nordic Optical Telescope. A variable longitudinal magnetic field of the order
of a few hundred Gauss up to one kG (-952$\pm$177 G) was measured. Another
study by Hubrig et al. (2013), measured magnetic fields on 6 HAeBe stars using
the HARPS in spectropolarimeter mode. They measured a weak magnetic field of
$\sim$65 G for HD 104237. Järvinen et al. (2015), was able to measure the
magnetic field of the primary ($\sim$13 G) and the secondary ($\sim$129 G) of
the HD 104237 system using HARPS spectrograph. The primary is a HAe star and
the secondary is a TTS (Table LABEL:TableC1:individualstars). Petit et al.
(2008) detected a magnetic field strength of -240$\pm$70 G was detected for HD
36982. This was further confirmed by Alecian et al. (2013) with moderate
detection. We see that the Herbig Ae/Be stars possess much weaker magnetic
fields than their lower mass counterpart T Tauri stars which have magnetic
fields in kG (Järvinen et al., 2019b). The 15 HAeBe stars mentioned in the
above discussion emit X-rays whose origin is intimately linked to the presence
of the magnetic field. Of these 15 stars, 11 are single HAeBe stars and 4 are
non-single HAeBe stars. We find that from our sample of 25 single HAeBe stars,
11 are confirmed to have magnetic fields. Apart from these 11 confirmed
magnetic field detections, magnetic fields were measured on 4 single HAeBe
stars. However, only upper limits were obtained. For the remaining 10 single
HAeBe stars magnetic field detections have not been studied. Among the 19 non-
single HAeBe stars, 4 stars are confirmed to have magnetic field detections.
Upper limits were measured for 3 non-single HAeBe stars. The magnetic fields
of the remaining 12 non-single HAeBe stars have not yet been studied. Further
studies on measuring the magnetic fields of X-ray-emitting HAeBe stars would
provide more insights.

###  5.3 Possible X-ray emission mechanisms in HAeBe stars

Sun-like stars contain outer convection zones and the X-rays from the corona
are linked to a solar-like dynamo. A similar mechanism is thought to be the
origin of X-rays in TTSs (Preibisch et al., 2005; Telleschi et al., 2007a).
Magnetospheric accretion and coronal X-ray are mainly the reason for X-ray
emission in TTSs. Massive OB-type stars produce X-rays from radiation-driven
instability shocks and magnetically confined wind shocks (Babel & Montmerle,
1997b; Zhekov & Palla, 2007; Montmerle, 2007). Radiative wind models are also
ruled as the temperatures of HAeBe stars are not hot enough for the winds to
be generated. From the results obtained in this study, we see that the LX with
respect to the Lbol decreases towards earlier type stars (Figure 1 & 2),
suggesting a decrease in the X-ray activity levels towards early-type stars.
Wichittanakom et al. (2020) showed that the mass-accretion rates and accretion
luminosity of HAeBe stars increase with the increase in mass. Our studies show
that LHα, which is a good tracer of mass accretion rate, does not correlate
with the LX of HAeBe stars, strongly suggesting that in-falling material from
the disk is not responsible for X-ray emission (Figure 5). Stelzer et al.
(2006) and Stelzer et al. (2009a) suggest that X-ray emission in HAeBe stars
could be attributed to hidden low mass companions, as the binarity rate of
their sample was high and also the X-ray properties were very similar to that
of other low-mass stars. In Section 5.1, we have documented high-resolution
studies that confirm the presence of planetary companions rather than stellar
companions in our HAeBe sample (Table LABEL:TableC1:individualstars).
Similarly, for HAeBe stars with confirmed unresolved low mass companions, some
of which are observed at multi-epochs, we observe long-term and short-term
variability with flaring in their X-ray lightcurves. This result along with
the soft/cool X-ray temperature of single HAeBe stars and similar mass HAeBe
systems suggests that the X-ray emission in HAeBe stars may not be solely due
to hidden TTSs but is intrinsic to the star itself. But, what could be the
emission mechanism involved in the production of X-rays in HAeBe stars?

Telleschi et al. (2007b) and Günther & Schmitt (2009) studied the high-
resolution X-ray spectra of HAeBe stars AB Aurigae, and HD 163296,
respectively, which are obtained from the XMM-Newton Reflection Grating
Spectrometers (RGS) and EPIC instruments. These stars are confirmed to be
single and not known to have any sub-arcsec companions from the high-
resolution studies discussed in Section 5.1. The X-ray spectra of both stars
are very soft with low photoelectric absorption (the spectrum falls off
rapidly above 1 keV) pointing to a cool source. AB Aurigae has temperature
components of kT1 = 0.21${}^{+0.03}_{-0.03}$ keV, kT2 =
0.6${}^{+0.03}_{-0.04}$ keV, and LX = 4 $\times$ 1029 erg s-1 in the energy
range 0.3 $-$ 10.0 keV (Telleschi et al., 2007b). While, HD 163296 has
temperature components kT1 = 0.21${}^{+0.03}_{-0.01}$ keV, kT2 =
0.51${}^{+0.1}_{-0.03}$ keV, kT3 = 2.7${}^{+1.5}_{-0.8}$ keV, and LX = 4.1
$\times$ 1029 erg s-1 in the energy range 0.2 $-$ 8.0 keV (Günther & Schmitt,
2009). About 35% of the emission measure was found at 0.2 keV for AB Auriage,
whereas it is $\sim$ 60% for HD 163296. Further, Günther et al. (2013) studied
the Chandra data of the source HD 163296, using the same method as in Günther
& Schmitt (2009) and obtained similar model parameters with kT1 =
0.19${}^{+0.23}_{-0.02}$ keV, kT2 = 0.60${}^{+0.05}_{-0.04}$ keV, kT3 =
2.0${}^{+1.0}_{-0.5}$ keV, and LX = 3.8 $\times$ 1029 erg s-1 for Obs ID: 3733
and kT1 = 0.16${}^{+0.01}_{-0.01}$ keV, kT2 = 0.60${}^{+0.2}_{-0.02}$ keV, kT3
= 2.4${}^{+0.5}_{-0.3}$ keV and LX = 6.3 $\times$ 1029 erg s-1 for Obs ID:
12359. These values within the error bars are consistent with our study.
Telleschi et al. (2007b) and Günther & Schmitt (2009) analyzed the density-
sensitive He-like line triplets of O vii, observed in the high-resolution
spectra of AB Aurigae, and HD 163296. The flux ratio of the forbidden (f) to
intercombination (i) lines (i.e., $\mathcal{R}=f/i$ ratio) at 22.1 Å, and 21.8
Å, respectively, point to the low-density regions (such as stellar coronae; ne
< 1011 cm-3) as the location for the production of majority the of X-ray
emission. The lightcurve of AB Aurigae shows a period of $\sim$ 42 hours which
is consistent with the modulation period of Mg ii and He i lines formed in the
chromosphere. No such period was found in the lightcurve of HD 163296, but a
short-term variability was observed. Both studies reject the accretion shock
model and companion hypothesis. The accretion shock model is rejected based on
the high f / i line ratios obtained from the triplets which suggest a low-
density environment as the formation region of X-rays, whereas accretion
shocks are generated in regions close to the stellar photosphere, in the high-
density regions. HD 163296 drives a jet HH 409, which could be the reason for
the excessive soft X-ray emission, while the $\sim$ 2 keV component could be
from the coronae. The above-mentioned characteristics observed in the high-
resolution X-ray spectra of the HAeBe stars AB Aurigae, and HD 163296 clearly
point to X-ray emission in the low-density regions such as coronae which
supports the presence of magnetic fields in these stars. Based on the evidence
discussed above, Telleschi et al. (2007b) and Günther & Schmitt (2009)
suggested either a wind shock model or magnetically confined winds could be
the mechanism responsible for the emission of X-rays.

From the above discussion on individual stars AB Aurigae, and HD 163296 and
the magnetic field detections of the HAeBe stars in Section 5.2. It is
plausible to say that the X-ray emission in HAeBe stars is due to the presence
of magnetic fields. However, the origin of the magnetic field remains unknown.
In the study by Hubrig et al. (2009), the magnetic fields were found to be
stronger in younger HAeBe stars. They also observe that the strength of the
magnetic field seems to be increasing with an increase in X-ray luminosity for
HAe stars in their sample. The HAe stars are young ($\sim$ 1 $-$ 2 Myr),
implying that the X-ray emission decays over time as the star reaches the Main
Sequence. This agrees well with the model proposed by Tout & Pringle (1995),
where a non-solar dynamo operates in rapidly rotating A-type stars based on
rotational sheer energy. Other possible mechanisms for X-ray emission from
magnetic activity could be due to fossil fields or magnetically confined wind
shocks (Babel & Montmerle, 1997a). The HAeBe stars identified as single stars
in this paper (Table LABEL:TableC1:individualstars), are less studied and do
not have enough data to understand their environment and companion status.
This can be further explored with multi-wavelength photometry and spectrometry
along with high-resolution spectro-polarimetric data to better understand the
X-ray emission mechanism responsible for these stars. More evidence is
required to support the origin of X-ray emission due to magnetic fields in
HAeBe stars.

##  6 Conclusions

We present the X-ray analysis of a large sample of HAeBe stars, investigated
with the Chandra X-ray telescope having a remarkable $\sim$1" spatial
resolution. The main results obtained from this study are given below :

  1. 1.

Out of 62 HAeBe sources, 44 have been detected in X-rays, i.e., a detection
rate of $\sim$70.9% is observed. We present new X-ray detections of 7 HAeBe
stars in this work. The LX of HAeBe stars ranges between 1028 to 1033 erg s-1
slightly higher than TTSs. The fractional X-ray luminosity of HAeBe stars
ranges from $\log$(LX/Lbol) = -3.5 to -10, most of them falling between
$\log$(LX/Lbol) = -4 to -7, which is in between the value observed for TTSs
and OB-type MS stars.

  2. 2.

In our sample, 25 out of 44 X-ray emitting HAeBe stars (56.2%) do not have any
reported companions (the single status of some of the sources has been
confirmed from previous high-resolution imaging studies). Some of the multi-
star systems show variability and flaring, usually observed in TTSs. This
evidence rules out the presence of low-mass companions as solely being the
X-ray emitters in our sample of HAeBe stars.

  3. 3.

The plasma temperatures of HAeBe stars without any known companions are in the
soft energy regime (kT $\leq$ 2 keV), differing from TTSs in which the hot
component (kT > 2 keV) dominates. About 50% of non-single HAeBe stars fit with
the 1T model and 80% of non-single HAeBe stars fit with the 2T model show hot
X-ray emission, i.e., kT > 2 keV.

  4. 4.

We do not find any correlation between the X-ray luminosity and the disk
properties i.e., IR excess ($\eta_{2-4.6}$) and fractional accretion
luminosity (LHα/Lbol) for the HAeBe stars in our sample, implying that the
X-ray emission is not related to the dust content in the disk or accretion
shocks.

  5. 5.

From the X-ray spectrum and lightcurve analysis of the HAeBe stars in our
sample, we see that the stars with confirmed low-mass companions have a plasma
temperature on the hotter side following a similar pattern as TTSs. 9 stars
are found to show variability/flaring features in the X-ray lightcurve. Such
X-ray variability/flaring is not observed in single HAeBe stars and equal mass
HAeBe star systems, except V892 Tau and V372 Ori.

From the results obtained, we propose that the X-ray emission from HAeBe is
intrinsic and not due to the presence of a close low-mass companion. The
X-rays from the HAeBe stars need to be studied further in-depth to constrain
the emission mechanism. Studies by Telleschi et al. (2007b) and Günther &
Schmitt (2009) suggest wind shock models and magnetically confined winds as
possible mechanisms for X-ray emission. The validation of these models
requires the presence of magnetic fields in HAeBe stars. Based on the previous
studies on magnetic fields in literature we found that in single HAeBe stars
(i.e., stars with no known sub-arcsec companion), structured magnetic fields
could be present. This favors the sheer dynamo theory (Tout & Pringle, 1995),
fossil field hypothesis, or magnetically confined wind shocks (Babel &
Montmerle, 1997a). Further studies with multi-wavelength photometry and
spectrometry along with high-resolution spectro-polarimetric data are
suggested to support the magnetically driven coronal emission model in these
stars.

## Acknowledgements

We thank the referee for providing constructive comments and suggestions that
have helped enhance the quality of the paper. This research has made use of
data obtained from the Chandra Data Archive and the software provided by the
Chandra X-ray Center (CXC) in the application package “CIAO”. We also made use
of the software “XSPEC” provided by the High Energy Astrophysics Science
Archive Research Center (HEASARC), which is a service of the Astrophysics
Science Division at NASA/GSFC. We thank the help desk at Chandra and HEASARC
for clarifying doubts regarding the X-ray data analysis and software issues.
This study has used the Gaia EDR3 data to obtain the distances for the sample
of stars in this paper. Hence, we express our gratitude to the Gaia
collaboration. We thank the SIMBAD database and VizieR online library service
for helping with the relevant literature survey. We would like to thank our
colleagues Dr. Savithri H Ezhikode, Ms. Nidhi Sabu, and Dr. Arun Roy for
providing help with the X-ray and optical analysis part of the paper. We
acknowledge the support given by the Center for Research, CHRIST (Deemed to be
University), Bangalore, India. VJ and SSK thank the Inter-University Centre
for Astronomy and Astrophysics (IUCAA), Pune, India, for the Visiting
Associateship.

## Data Availability

The data used in this work was accessed from the Chandra Data Archive (CDA;
https://cda.harvard.edu/chaser/mainEntry.do).

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##  Appendix A Tabulated parameters of Chandra Observations of Herbig Ae/Be
stars

Table 1: Observation details of Herbig Ae/Be stars SI No. | Object | X-ray Iden. | Obs ID | Instrument | RA (J2000) | Dec (J2000) |  A${}_{\mathrm{V}}^{(a)}$ |  Distance(pc)(b) | No. of Obs. | Ref.
---|---|---|---|---|---|---|---|---|---|---
1 | AB Aurigae | X | 3755 | ACIS-S | 04:55:45.86 | 30:33:04.49 | 0.39 | $155.05\pm 0.81$ | 1 | 10
2 | AK Sco | X |  combined444AK Sco: 983+13274+13275 | ACIS-I | 16:54:44.85 | -36:53:18.55 | 0.35 | $139.12\pm 0.57$ | 3 | 4,10,23
3 | AS 310 | X | 6399 | ACIS-S | 18:33:21.19 | -04:58:05.78 | 6.06 | $2390.61\pm 77.58$ | 1 | 13
4 | BD+30 549 | X |  combined555BD+30 549: 642+6436+6437 | ACIS-I | 03:29:19.81 | 31:24:56.81 | 2.30 | $284.78\pm 1.93$ | 3 | 4
5 | AS 477 | X2 | 6401 | ACIS-S | 21:52:34.10 | 47:13:43.60 | 2.00 | $756.28\pm 7.05$ | 1 | 13
6 | V361 Cep | X | 6400 | ACIS-S | 21:42:50.22 | 66:06:35.08 | 2.78 | $882.79\pm 11.53$ | 1 | 13
7 | BD+65 1638 | X2 | 6400 | ACIS-S | 21:42:58.84 | 66:06:10.21 | 3.39 | $1150\pm-$ | 1 | 13
8 | BP Psc | X |  combined666BP PSc: 10856+8900 | ACIS-S | 23:22:24.71 | -02:13:41.76 | 1.34 | $211.39\pm 23.48$ | 2 | 18
9 | PDS 69 | X2 | 6423 | ACIS-I | 13:57:43.92 | -39:58:47.27 | 2.58 | $674.34\pm 7.63$ | 1 | 11
10 | V892 Tau | X2 | 3364 | ACIS-S | 04:18:40.62 | 28:19:15.54 | 11.99 | $133.06\pm 1.69$ | 2 | 4,10
11 | HD 100453 | X2 | 6429 | ACIS-S | 11:33:05.63 | -54:19:28.95 | 0.10 | $103.68\pm 0.22$ | 1 | 15
12 | HD 100546 | X |  combined777HD 100546: 2403+3427 | ACIS-I | 11:33:25.39 | -70:11:41.43 | 0.44 | $107.98\pm 0.42$ | 2 | 4,10
13 | HD 104237 | X2 | 7326 | ACIS-S | 12:00:04.99 | -78:11:34.66 | 0.59 | $106.48\pm 0.53$ | 6 | 4,10
14 | HD 141569 | $-$ | 981 | ACIS-I | 15:49:57.70 | -03:55:17.00 | 0.50 | $111.33\pm 0.37$ | 1 | 4
15 | HD 144432 | X2 | 6398 | ACIS-S | 16:06:57.93 | -27:43:10.09 | 0.53 | $154.16\pm 0.57$ | 1 | 13
16 | HD 147889 | $-$ | 618 | ACIS-I | 16:25:24.32 | -24:27:56.60 | 5.35 | $135.65\pm 0.42$ | 1 | 4,10
17 | HD 150193 | X2 | 982 | ACIS-I | 16:40:17.95 | -23:53:44.96 | 2.08 | $149.89\pm 0.41$ | 1 | 4,10
18 | HD 163296 | X | 12359 | ACIS-S | 17:56:21.29 | -21:57:22.38 | 0.24 | $100.50\pm 0.34$ | 2 | 4,10,20
19 | HD 169142 | X | 6430 | ACIS-S | 18:24:29.79 | -29:46:49.52 | 0.60 | $114.44\pm 0.30$ | 1 | 7
20 | HD 176386 | $-$ | 3499 | ACIS-I | 19:01:38.90 | -36:53:27.00 | 0.61 | $154.37\pm 0.64$ | 8 | 9
21 | HD 259431 | X2 | 6397 | ACIS-S | 06:33:05.21 | 10:19:20.16 | 2.05 | $642.26\pm 11.33$ | 1 | 13
22 | HD 31648 | X | 8939 | ACIS-S | 04:58:46.27 | 29:50:36.64 | 0.55 | $155.08\pm 1.16$ | 1 | 16
23 | HD 97300 | X | 1867 | ACIS-I | 11:09:50.08 | -76:36:47.65 | 2.15 | $188.78\pm 1.25$ | 1 | 4
24 | V373 Cep | $-$ | 6400 | ACIS-S | 21:43:06.68 | 66:06:54.60 | 5.14 | $870.73\pm 14.91$ | 1 | 13
25 | LkHa 25 | $-$ | 2550 | ACIS-I | 06:40:44.60 | 09:48:02.00 | 1.98 | $694.02\pm 49.67$ | 3 | 4
26 | HD 36982 | X |  combined888HD 36982: 4395+3744+4373+4374+4396+3498+17735 | ACIS-I | 05:35:09.84 | -05:27:52.95 | 1.32 | $404.00\pm 3.82$ | 54 | 3
27 | MR Ori | $-$ | 4373 | ACIS-S | 05:35:16.97 | -05:21:45.33 | 2.58 | $395.34\pm 2.49$ | 86 | 10
28 | MWC 297 | X2 | 1883 | ACIS-I | 18:27:39.54 | -03:49:51.93 | 12.46 | $407.41\pm 5.19$ | 1 | 4,10
29 | PDS 144S | $-$ | 11012 | ACIS-S | 15:49:15.32 | -26:00:54.70 | 0.92 | $1174.13\pm 119.29$ | 1 | *
30 | PDS 581 | $-$ | 4504 | ACIS-S | 19:36:18.90 | 29:32:50 | 4.24 | $1574.38\pm 208.99$ | 1 | *
31 | R CrA | X |  combined 999R CrA: 19+3499+4475+5402+5403+5404+5405+19709 | ACIS-I | 19:01:53.66 | -36:57:08.01 | 5.45 | $126.02\pm 8.19$ | 9 | 6
32 | TY CrA | X | 19709 | ACIS-I | 19:01:40.80 | -36:52:34.01 | 2.84 | $159.13\pm 4.20$ | 9 | 9
33 | T CrA | X |  combined101010T CrA: 19+5404 | ACIS-I | 19:01:58.79 | -36:57:50.33 | 3.00 | $130.00\pm 0.00$ | 7 | 9
34 | HR 6000 | X | 9921 | ACIS-I | 16:08:34.55 | -39:05:34.18 | 0.16 | $157.43\pm 1.05$ | 13 | 13
35 | HD 37062 | X | 4374 | ACIS-I | 05:35:31.42 | -05:25:16.12 | 0.65 | $405.49\pm 6.45$ | 67 | 1,2
36 | V372 Ori | X | 6418 | ACIS-I | 05:34:46.96 | -05:34:15.15 | 0.84 | $447.95\pm 10.77$ | 7 | 10
37 | V380 Ori | X | 12391 | ACIS-S | 05:36:25.40 | -06:42:57.99 | 4.41 | $379.64\pm 15.06$ | 2 | 4
38 | HR 5999 | X2 | 8901 | ACIS-I | 16:08:34.27 | -39:06:18.34 | 0.54 | $157.67\pm 0.72$ | 7 | 13
39 | Z CMa | X | 10845 | ACIS-S | 07:03:43.16 | -11:33:06.24 | 7.89 | $639.94\pm 222.48$ | 2 | 4,14
40 | HBC 217 | X | 2540 | ACIS-I | 06:40:42.18 | 09:33:37.26 | 0.29 | $703.07\pm 6.98$ | 7 | 25
41 | HBC 442 | X |  combined111111HBC 442: 6416+6417+8936 | ACIS-S | 05:34:14.16 | -05:36:54.38 | 0.38 | $383.52\pm 2.58$ | 3 | *
42 | HD 135344 B | X | 9927 | ACIS-S | 15:15:48.44 | -37:09:16.34 | 0.37 | $134.55\pm 0.46$ | 1 | 22
43 | HBC 222 | X | 2550 | ACIS-S | 06:40:51.20 | 09:44:46.00 | 0.18 | $704.58\pm 6.86$ | 3 | 21
44 | HD 135344 | $-$ | 9927 | ACIS-S | 15:15:48.90 | -37:08:56.00 | 0.36 | $134.18\pm 0.58$ | 1 | *
45 | LkHa 257 | $-$ | 15723 | ACIS-I | 21:54:18.80 | 47:12:10.00 | 2.28 | $786.10\pm 8.13$ | 1 | *
46 | LkHa 260 | $-$ | 978 | ACIS-I | 18:19:09.40 | -13:50:41.00 | 5.16 | $1556.18\pm 89.85$ | 5 | *
47 | LkHa 339 | $-$ | 12392 | ACIS-I | 06:10:57.80 | -06:14:40.00 | 4.15 | $824.55\pm 10.46$ | 1 | *
48 | MWC 930 | $-$ | 23646 | ACIS-I | 18:26:25.20 | -07:13:18.00 | 14.06 | $7129.11\pm 2189.23$ | 1 | *
49 | RR Tau | $-$ | 8242 | ACIS-I | 05:39:30.50 | 26:22:27.00 | 3.62 | $798.19\pm 18.67$ | 1 | *
50 | GSC 3975-0579 | $-$ | 16309 | ACIS-S | 21:38:08.50 | 57:26:48.00 | 1.26 | $901.56\pm 13.67$ | 2 | *
51 | HD 36939 | X | 4396 | ACIS-S | 05:34:55.29 | -05:30:22.09 | 0.81 | $416.48\pm 7.24$ | 45 | 1,2
52 | HD 250550 | X | 21185 | ACIS-S | 06:02:00.00 | 16:30:57 | 0.70 | $748.41\pm 25.88$ | 1 | *
53 | R Mon | X | 21183 | ACIS-S | 06:39:09.95 | 08:44:09.55 | 5.66 | $800.00\pm 0.00$ | 1 | *
54 | HD 244604 | X | 21187 | ACIS-S | 05:31:57.23 | 11:17:41.52 | 1.69 | $397.93\pm 3.33$ | 1 | *
55 | V346 Ori | $-$ | 21186 | ACIS-S | 05:24:42.80 | 01:43:48.00 | 0.31 | $336.39\pm 2.39$ | 2 | *
56 | [DLM2010] EC 95a | X | 4479 | ACIS-I | 18:29:57.89 | 01:12:46.15 | 36.00 | $429.20\pm 2.00$ | 1 | 8
57 | T Ori | $-$ | 2567 | ACIS-S | 05:35:50.50 | -05:28:35.00 | 2.54 | $398.95\pm 4.90$ | 25 | *
58 | MWC 953 | X | 2298 | ACIS-I | 18:43:28.42 | -03:46:17.03 | 5.84 | $1985.43\pm 68.25$ | 1 | *
59 | PDS 37 | $-$ | 7519 | ACIS-S | 10:10:00.32 | -57:02:07.35 | 9.37 | $1626.95\pm 63.06$ | 1 | *
60 | HD 190073 | X | 21184 | ACIS-S | 20:03:02.51 | 05:44:16.65 | 0.34 | $824.88\pm 21.93$ | 1 | *
61 | HD 245906 | X | 8242 | ACIS-I | 05:39:30.47 | +26:19:55.15 | 1.11 | $900.58\pm 296.13$ | 1 | *
62 | AFGL 961 | X | 12142 | ACIS-I | 06:34:37.74 | 04:12:44.20 | 32.00 | $1500.00\pm 0.00$ | 1 | 21

“$-$” in column (3) denotes non-detections. “*” in column (11) denotes new
stars which are not previously studied.
Table LABEL:TableA1:xrayobservations: (a) The AV values were compiled from
Preibisch (1999), Hamaguchi et al. (2005), Stelzer et al. (2006), Manoj et al.
(2006), Stelzer et al. (2009a), Dahm & Hillenbrand (2015), Fairlamb et al.
(2015), Mathew et al. (2018), Vioque et al. (2018) and Arun et al. (2019).
(b) Distance measurements were compiled from Preibisch (1999), Manoj et al.
(2006), Sandell et al. (2011), (Dahm & Hillenbrand, 2015) and Bailer-Jones et
al. (2021).
References: (1) Feigelson et al. (2002), (2) Feigelson et al. (2003), (3)
Getman et al. (2005), (4) Stelzer et al. (2006), (5) Flaccomio et al. (2006),
(6) Forbrich et al. (2006), (7) Grady et al. (2007), (8) Giardino et al.
(2007), (9) Forbrich & Preibisch (2007), (10) Hamidouche et al. (2008), (11)
Getman et al. (2008), (12) Testa et al. (2008), (13) Stelzer et al. (2009a),
(14) Stelzer et al. (2009b), (15) Collins et al. (2009), (16) Grady et al.
(2010), (17) Winston et al. (2010), (18) Kastner et al. (2010), (19) Forbrich
et al. (2011), (20) Günther et al. (2013), (21) Kuhn et al. (2013), (22)
McJunkin et al. (2014), (23) Getman et al. (2016), (24) Guarcello et al.
(2017), (25) Townsley et al. (2019).

##  Appendix B X-ray parameters of HAeBe Stars

Table 2: X-ray parameters of HAeBe Stars from the Model Fitting Name |  NH [1022 cm-2] | kT1 [keV] | kT2 [keV] | statistic |  Pnull | GOF | $\chi^{2}$ | ${\chi_{\mathrm{red}}}^{2}$ | DOF |  $\log$ LX [erg s-1]
---|---|---|---|---|---|---|---|---|---|---
AB Aurigae | 0.05 |  0.78$\pm$0.13 |  | c |  | 53.00% |  |  | 10 |  29.11$\pm$0.10
AK Sco | 0.04 |  0.59$\pm$0.14 |  | c |  | 31.00% |  |  | 8 |  29.10$\pm$0.09
AS 310 | 0.70 |  0.57$\pm$0.33 |  | c |  | 52.00% |  |  | 4 |  31.00$\pm$0.24
BD+30 549 |  2.22$\pm$1.04 |  0.23$\pm$0.14 |  | $\chi^{2}$ | 5.87E-01 |  | 1.93 | 0.64 | 3 |  31.29$\pm$1.22
V361 Cep | 0.32 |  1.54$\pm$0.56 |  | c |  | 22.00% |  |  | 12 |  30.56$\pm$0.10
BP Psc | 7.94 |  2.07$\pm$19.15 |  | c |  | 3.00% |  |  | 7 |  29.17$\pm$0.49
HD 100546 | 0.05 |  1.40$\pm$0.44 |  | c |  | 9.00% |  |  | 9 |  29.21$\pm$0.10
HD 163296 |  0.18$\pm$0.05 |  0.12$\pm$0.04 |  0.58$\pm$0.04 | $\chi^{2}$ | 1.40E-01 |  | 83.94 | 1.18 | 71 |  29.89$\pm$0.14
HD 169142 | 0.07 |  0.21$\pm$0.02 |  | c |  | 39.00% |  |  | 13 |  29.19$\pm$0.09
HD 31648 |  0.52$\pm$0.25 |  0.48$\pm$0.16 |  | $\chi^{2}$ | 1.74E-01 |  | 7.69 | 1.538 | 5 |  29.91$\pm$0.36
HD 97300 |  0.12$\pm$0.06 |  2.50$\pm$0.38 |  | $\chi^{2}$ | 2.45E-02 |  | 77.49 | 1.41 | 55 |  29.87$\pm$0.03
HD 36982 | 0.15 |  1.06$\pm$0.22 |  | $\chi^{2}$ | 2.79E-01 |  | 10.96 | 1.22 | 9 |  28.58$\pm$0.07
R CrA | 0.63 |  |  | PIMMS |  |  |  |  |  |  29.85$\pm$0.06
TY CrA | 0.33 |  1.41$\pm$0.15 |  8.25$\pm$3.23 | $\chi^{2}$ | 2.21E-01 |  | 341.2 | 1.06 | 322 |  31.45$\pm$0.02
HR 6000 | 0.02 |  0.74$\pm$0.09 |  1.29$\pm$0.24 | $\chi^{2}$ | 2.01E-01 |  | 72.15 | 1.15 | 63 |  30.15$\pm$0.02
HD 37062 |  0.23$\pm$0.03 |  1.01$\pm$0.07 |  2.75$\pm$0.22 | $\chi^{2}$ | 2.61E-02 |  | 177.63 | 1.24 | 143 |  31.12$\pm$0.02
V372 Ori | 0.10 |  0.94$\pm$0.21 |  3.47$\pm$0.80 | $\chi^{2}$ | 4.64E-01 |  | 71.41 | 1.01 | 71 |  30.96$\pm$0.03
V380 Ori |  1.05$\pm$0.13 |  0.45$\pm$0.11 |  4.32$\pm$2.95 | $\chi^{2}$ | 7.46E-02 |  | 123.27 | 1.21 | 102 |  31.91$\pm$0.18
Z CMa | 0.92 |  0.66$\pm$0.20 |  | $\chi^{2}$ | 2.90E-01 |  | 6.17 | 1.23 | 5 |  30.62$\pm$0.34
HBC 217 | 0.03 |  0.98$\pm$0.11 |  | $\chi^{2}$ | 5.14E-01 |  | 6.23 | 0.89 | 7 |  29.95$\pm$0.06
HBC 442 | 0.04 |  0.86$\pm$0.12 |  | $\chi^{2}$ | 1.73E-01 |  | 16.4 | 1.37 | 12 |  29.80$\pm$0.06
HD 135344B | 0.04 |  0.20$\pm$0.14 |  0.74$\pm$0.07 | $\chi^{2}$ | 7.94E-01 |  | 29.79 | 0.81 | 37 |  29.55$\pm$0.03
HD 36939 | 0.09 |  0.93$\pm$0.11 |  | $\chi^{2}$ | 2.65E-01 |  | 13.46 | 1.22 | 11 |  29.65$\pm$0.06
HD 250550 | 0.08 |  1.02$\pm$0.54 |  | c |  | 38.00% |  |  | 3 |  30.15$\pm$0.30
HD 244604 | 0.20 |  |  | PIMMS |  |  |  |  |  |  29.18$\pm$0.26
HD 245906 | 0.13 |  0.65$\pm$0.35 |  | $\chi^{2}$ | 5.92E-01 |  | 1.91 | 0.64 | 3 |  30.75$\pm$0.35
R Mon | 0.66 |  |  | PIMMS |  |  |  |  |  |  30.76$\pm$0.15
AFGL 961 |  22.89$\pm$15.31 |  1.13$\pm$1.49 |  | $\chi^{2}$ | 4.54E-01 |  | 2.62 | 0.87 | 3 |  33.12$\pm$1.39
[DLM2010] EC 95a |  3.59$\pm$0.18 |  3.05$\pm$0.26 |  | $\chi^{2}$ | 7.90E-01 |  | 175.99 | 0.92 | 192 |  31.79$\pm$0.03
HBC 222 | 0.02 |  0.65$\pm$0.32 |  | c |  | 16.00% |  |  | 9 |  30.08$\pm$0.11
MWC 953 | 0.68 |  |  | PIMMS |  |  |  |  |  |  30.64$\pm$0.27
HD 190073 | 0.04 |  |  | PIMMS |  |  |  |  |  |  29.86$\pm$0.34
T CrA | 0.35 |  |  | PIMMS |  |  |  |  |  |  28.35$\pm$0.25
AS 477 A | 0.23 |  1.53$\pm$3.26 |  | c |  | 13.00% |  |  | 4 |  29.71$\pm$0.23
BD+65 1638 A | 0.39 |  1.98$\pm$0.59 |  | $\chi^{2}$ | 5.06E-01 |  | 9.28 | 0.928 | 10 |  31.23$\pm$0.05
PDS 69 A | 0.30 |  1.56$\pm$0.59 |  | c |  | 41.00% |  |  | 10 |  30.20$\pm$0.10
V892 Tau A |  1.06$\pm$0.17 |  2.07$\pm$0.30 |  | $\chi^{2}$ | 1.41E-01 |  | 80.61 | 1.19 | 68 |  30.61$\pm$0.06
HD 100453 A | 0.01 |  0.27$\pm$0.05 |  | c |  | 3.00% |  |  | 6 |  28.68$\pm$0.12
HD 104237 A |  0.36$\pm$0.11 |  0.34$\pm$0.05 |  1.91$\pm$0.30 | $\chi^{2}$ | 6.03E-02 |  | 82.43 | 1.29 | 64 |  30.91$\pm$0.17
HD 144432 A | 0.06 |  0.57$\pm$0.14 |  | c |  | 13.00% |  |  | 8 |  29.04$\pm$0.10
HD 150193 A | 0.24 |  0.87$\pm$0.49 |  | c |  | 13.00% |  |  | 3 |  29.47$\pm$0.26
HD 259431 A | 0.24 |  2.51$\pm$0.91 |  | $\chi^{2}$ | 3.17E-01 |  | 10.43 | 1.16 | 9 |  30.98$\pm$0.07
MWC 297 A | 1.45 |  1.84$\pm$2.39 |  | c |  | 2.00% |  |  | 8 |  29.56$\pm$0.18
HR 5999 A | 0.06 |  |  | PIMMS |  |  |  |  |  |  28.91$\pm$0.23
AS 477 B | 0.23 |  1.48$\pm$0.72 |  | c |  | 47.00% |  |  | 14 |  30.37$\pm$0.10
BD+65 1638 B | 0.39 |  1.13$\pm$0.10 |  | $\chi^{2}$ | 6.40E-02 |  | 26.61 | 1.57 | 17 |  31.40$\pm$0.04
PDS 69 B | 0.30 |  |  | PIMMS |  |  |  |  |  |  30.07$\pm$0.16
V892 Tau B | 1.39 |  1.16$\pm$0.26 |  | $\chi^{2}$ | 3.84E-01 |  | 8.53 | 1.07 | 8 |  29.99$\pm$0.12
HD 100453 B | 0.01 |  0.73$\pm$0.22 |  | c |  | 36.00% |  |  | 7 |  28.68$\pm$0.10
HD 104237 B | 0.07 |  0.64$\pm$0.15 |  | c |  | 13.00% |  |  | 9 |  29.17$\pm$0.11
HD 144432 B |  0.03$\pm$0.02 |  0.85$\pm$0.11 |  3.23$\pm$0.65 | $\chi^{2}$ | 7.53E-02 |  | 102.17 | 1.23 | 83 |  30.35$\pm$0.04
HD 150193 B | 0.24 |  1.40$\pm$0.26 |  | $\chi^{2}$ | 3.61E-01 |  | 7.69 | 1.10 | 7 |  30.29$\pm$0.07
HD 259431 B | 0.24 |  1.02$\pm$0.30 |  | c |  | 27.00% |  |  | 7 |  30.42$\pm$0.12
MWC 297 B | 1.45 |  |  | PIMMS |  |  |  |  |  |  29.69$\pm$0.21
HR 5999 B | 0.06 |  1.04$\pm$0.10 | >5.75 | $\chi^{2}$ | 1.73E-01 |  | 40.51 | 1.23 | 33 |  30.36$\pm$0.05
HD 176386 B | 0.07 |  0.88$\pm$0.08 |  1.68$\pm$0.34 | $\chi^{2}$ | 1.50E-01 |  | 100.68 | 1.16 | 87 |  30.30$\pm$0.02
V373 Cep B | 0.60 |  |  | PIMMS |  |  |  |  |  |  30.29$\pm$0.18
Table 3: Upper-limits for LX of non-detected sources in X-rays SI. No | Name |  NH [1022 cm-2] |  $\log$LX [erg s-1]
---|---|---|---
1 | HD 141569 A | 0.06 | < 28.75
2 | HD 147889 | 0.62 | < 28.65
3 | V373 Cep A | 0.60 | < 29.93
4 | LkHa 25 | 0.23 | < 29.26
5 | MR Ori | 0.30 | < 29.15
6 | PDS 144 S | 0.11 | < 27.79
7 | PDS 581 | 0.49 | < 30.40
8 | PDS 37 | 1.09 | < 31.17
9 | HD 176386 A | 0.07 | < 27.92
10 | HD 135344 | 0.04 | < 28.11
11 | LkHa 257 | 0.27 | < 29.17
12 | LkHa 260 | 0.60 | < 30.32
13 | LkHa 339 | 0.48 | < 29.15
14 | MWC 930 | 1.63 | < 32.01
15 | RR Tau | 0.42 | < 29.58
16 | GSC 3975-0579 | 0.15 | < 30.29
17 | V346 Ori | 0.04 | < 28.94
18 | T Ori | 0.30 | < 29.83

Table LABEL:TableB1:xrayparam: Column(1) provides the name of the object,
Column(2) the absorption column density in 1022 cm-2 units, Column(3) & (4)
the plasma temperature in keV, Column(5) identifies the test statistic used
for model fitting, Column(6)(7)(8)(9) & (10) gives the null-hypothesis value,
goodness of fit, $\chi^{2}$, ${\chi_{\mathrm{red}}}^{2}$ and degrees of
freedom, finally Column(11) gives LX in log scale. Table
LABEL:TableB2:xraydark: “<” denotes that the value is upper-limit.

##  Appendix C Details of the companion status of X-ray emitting stars

Table 4: Details of HAeBe stars detected in X-rays Object | SpT | Binary |  | Resolved in
---
Chandra (Sep")
| Details of companion
---
Reference
AB Aurigae | A0 | No | $-$ |  | Planetary companion (<1")
---
69, 73
AK Sco | F5 | Yes | No |  | Stellar companion
---
(F5-type; <1")
12, 39, 57
AS 310 | B1 | No | $-$ |  | $-$
---
75
BD+30 549 | B8 | No | $-$ |  | $-$
---
V361 Cep | B2 | No | $-$ |  | $-$
---
BP Psc | G9 | Yes | No |  | Stellar companion (<1")
---
31
HD 100546 | B9 | No | $-$ |  | Planetary companion (<1")
---
50, 60, 62, 63
HD 163296 | A0 | No | $-$ |  | Planetary companion (<1")
---
47, 70, 71
HD 169142 | B9 | No | $-$ |  | Planetary companion (<1")
---
43, 44, 48, 51
HD 31648 | A5 | No | $-$ |  | $-$
---
75
HD 97300 | B9V | Yes | No |  | Stellar companion
---
M3.5-type (<1")
7, 36, 37
HD 36982 | B2 | No | $-$ |  | $-$
---
R CrA | A5 | Yes | No |  | Stellar companions
---
equal mass star ($\sim$2.3M⊙; <1")
+M-type ($\sim$0".156)
54, 55, 58, 66
TY CrA | B9 | Yes | No |  | Quadruple star system
---
K2+F-type (<1.2 AU; <1")
M4-type ($\sim$41 AU; $\sim$0.3")
4, 13, 21, 27
T CrA | F0 | Yes | No |  | Stellar companion
---
($\sim$1M⊙; < 1")
74, 76
HR 6000 | A1.5 | $-$ | $-$ |  | Unknown companion
---
($\sim$ 52MJup; <1")
53
HD 37062 | B4 | Yes | No |  | Triple star system
---
2 stars (LMS+A-type; <1")
8, 46
V372 Ori | B9.5 | Yes | No |  | Stellar companion
---
Spectroscopic binary; A0.5-type
1,10
V380 Ori | A1 | Yes | No |  | Quadruple system
---
TTS+IR source (<1")
M5 type (>1")
26, 42
Z CMa | B5+F5 | Yes | No |  | Stellar companion
---
FU Ori (<1")
11, 67
HBC 217 | G0 | No | $-$ |  | $-$
---
75
HBC 442 | F8 | No | $-$ |  | $-$
---
HD 135344B | A0 | No | $-$ |  | Planetary companion (<1")
---
41, 45
HD 36939 | B8-B9 | No | $-$ |  | $-$
---
HD 250550 | B9 | Yes | No |  | Stellar companion
---
LMS ($\sim$0.8M⊙; <1")
59, 72
HD 244604 | A0 | No | $-$ |  | $-$
---
75
HD 245906 | A6 | Yes | No |  | Stellar companion
---
G5-type (<1")
33
R Mon | B8 IIIe | Yes | No |  | Stellar companion
---
TTS (<1")
5, 20
AFGL 961 | B2-B3/B5 | Yes | $-$ |  | Stellar companion
---
early B-type (>1")
35
[DLM2010] EC 95a | K2 | Yes | No |  | Stellar companion
---
TTS (<1")
30
HBC 222 | F8 | No | $-$ |  | $-$
---
75
MWC 953 | B2Ve | No | $-$ |  | $-$
---
75
V1295 Aql | A2IVe | No | $-$ |  | Planetary companion (<1")
---
72
AS 477 | A0 | Yes | Yes (1.4”) |  | Stellar companion
---
Visual companion (>1")
9, 28
BD+65 1638 | B2 | Yes | Yes (1.332”) |  | Stellar companion
---
B-type (>1")
28, 40
PDS 69 | B7 | Yes | Yes (3.774”) |  | Stellar companion
---
B4-type (possible HAeBe ; >1")
2, 15, 22, 32
V892 Tau | A0 | Yes | Yes (4.523”) |  | Stellar companion
---
A-type (<1") + M3-type (>1")
24, 61
HD 100453 | A9Ve | Yes | Yes (1.007”) |  | Planetary companion(<1")
---
\+ M-dwarf ($\sim$1")
18, 29, 49
HD 104237 | A4 | Yes | Yes (1.482”) |  | Stellar companion
---
K3 (<1") + M3/M4-type (>1")
16, 25, 38, 52, 68
HD 144432 | A8 | Yes | Yes (1.400”) |  | Stellar companion
---
K7+M1-type (>1")
17, 23, 34
HD 150193 | A0 | Yes | Yes (1.406”) |  | Stellar companion
---
K4-type (>1")
2, 14, 44
HD 259431 | B6 | Yes | Yes (3.112”) |  | Stellar companion
---
$\sim$0.6 M⊙ source (<1")
Visual companion (>1")
19, 33, 72
MWC 297 | B1 | Yes | Yes (3.532”) |  | Stellar companion
---
low mass star (<1"; $\sim$0.1-0.5M⊙)
Visual companion (>1")
56, 64, 72
HR 5999 | A7 | Yes | Yes (1.574”) |  | Stellar companion
---
TTS (>1")
3, 6, 65, 72

References: (1) Abt et al. (1991), (2) Reipurth & Zinnecker (1993), (3)
Stecklum et al. (1995), (4) Corporon et al. (1996), (5) Ghez et al. (1997),
(6) Close et al. (1997), (7) Leinert et al. (1997), (8) Corporon & Lagrange
(1999), (9) Maheswar et al. (2002), (10) Manoj et al. (2002), (11) Millan-
Gabet & Monnier (2002), (12) Alencar et al. (2003), (13) Chauvin et al.
(2003), (14) Fukagawa et al. (2003), (15) (Maheswar et al., 2004), (16) Böhm
et al. (2004), (17) Pérez et al. (2004), (18) Chen et al. (2006), (19) Baines
et al. (2006), (20) Fuente et al. (2006), (21) Forbrich & Preibisch (2007),
(22) Haikala & Olberg (2007), (23) Carmona et al. (2007), (24) Monnier et al.
(2008), (25) Testa et al. (2008), (26) Alecian et al. (2009), (27) Boersma et
al. (2009), (28) Stelzer et al. (2009a), (29) Collins et al. (2009), (30) Dzib
et al. (2010), (31) Kastner et al. (2010), (32) Haikala & Reipurth (2010),
(33) Wheelwright et al. (2010), (34) Müller et al. (2011), (35) Sandell et al.
(2011), (36) Kóspál et al. (2012), (37) Daemgen et al. (2013), (38) Garcia et
al. (2013), (39) Gómez de Castro et al. (2013), (40) Dahm & Hillenbrand
(2015), (41) Bae et al. (2016), (42) Rodríguez et al. (2016), (43) Fedele et
al. (2017), (44) Monnier et al. (2017), (45) Cazzoletti et al. (2018), (46)
GRAVITY Collaboration et al. (2018), (47) Isella et al. (2018), (48) Ligi et
al. (2018), (49) Wagner et al. (2018), (50) Brittain et al. (2019), (51)
Gratton et al. (2019), (52) Järvinen et al. (2019a), (53) Kervella et al.
(2019), (54) Mesa et al. (2019), (55) Sissa et al. (2019), (56) Ubeira-
Gabellini et al. (2020), (57) Gómez de Castro et al. (2020), (58) Launhardt et
al. (2020), (59) Laws et al. (2020), (60) Fedele et al. (2021), (61) Long et
al. (2021), (62) Pérez et al. (2020), (63) Pyerin et al. (2021), (64) Sallum
et al. (2021), (65) Panić et al. (2021), (66) Sandell et al. (2021), (67) Dong
et al. (2022), (68) Godoy et al. (2022), (69) Currie et al. (2022), (70)
D’Angelo & Marzari (2022), (71) Leiendecker et al. (2022), (72) Rich et al.
(2022), (73) Zhou et al. (2022), (74) Rigliaco et al. (2023), (75) Thomas et
al. (2023), (76) Whelan et al. (2023).

##  Appendix D HAeBe Stars: Variable lightcurves

Figure 7: Chandra ACIS lightcurves of 8 stars in the energy range 0.3 $-$ 8.0
keV that present variability and flaring. The fixed bin time in seconds for
each source is mentioned at the bottom right of the plot.

Figure 8: Chandra ACIS lightcurves of HD 37062 in the energy range 0.3 $-$ 8.0
keV presenting variability and flaring over a period of 20 years. The fixed
bin time in seconds for each source is mentioned at the bottom right of the
plot.