their mass-loss rate increase dramatically, to values as high as
$\dot{M}=10^{-4}$ M⊙ yr-1. This produces a (typically) spherical CSE, which
eventually can obscure the central star. For such embedded stars, referred to
as OH/IR objects, (reliable) optical (e.g. in particular Gaia) distance
determination is not possible. For the shorter period Miras, the so called
Period-Luminosity (PL) relation can be used to infer their distances, but it
has been shown that this relation does not hold for variable stars having
period greater than $\sim$450 days Whitelock et al. (1991) and hence cannot be
used for OH/IR stars which periods are typically ranging between 1 to 5-6
years. Other distance determination such as kinematic distances and maser
parallaxes are in theory possible but, the former can be substantially
imprecise due to peculiar motions (Reid et al. 2009), while maser parallaxes
requires the persistence of (compact) maser spots over the duration of a
period which can be quite difficult to achieve especially for the objects
exhibiting (very) long periods.

OH/IR stars have the property of emitting a very strong (typically tens of Jy)
1612-MHz maser emission. The layer where this emission comes from in the
radially expanding CSE is such that acceleration is not longer expected and
the maser layer itself is of very small extent compared to the diameter of the
shell itself. These basic properties of the CSE in terms of velocity field and
structure and the need of velocity coherence for amplification of the maser
emission to be detected in a given line of sight, produce the characteristic
double-peak spectrum observed towards these objects with the “blue(-shifted)”
peak emanating from the front cap of the CSE while the “red(-shifted)” peak
from the back cap.

Phase-lag distance determinations rely on 2 independent measurements of the
diameter of the OH-maser shell and are possible because of the very nature of
the CSE itself and the fact that the OH maser emission, being radiatively
pumped, follow the (optical) stellar pulsation albeit with a small delay
(typically of 15-20%).

The first ingredient: the phase-lag is the measurement of the delay in the
light curves of the “red” peak with respect to the “blue” peak simply due to
the light travel time across the CSE, which ranges typically between several
days to months for the most sizable OH/IR stars. The second ingredient: the
angular diameter is obtained by (a single epoch) interferometric mapping,
providing the angular diameter of the shell (cf. Fig. 11).

Figure 11: _t_ op: eMERLIN map of the integrated emission over the inner part
of the spectrum of OH 83.4-0.9 covering the velocity range $\sim$[$-55$;$-23$]
km s-1. The magenta full and dotted circles delineate the outer part of the
shell. middle: “raw” blue-peak and red-peak light curves of OH 83.4-0.9
obtained from the NRT monitoring over a period of 7 years. bottom: scaled and
shifted light curves for phase-lag determination (Etoka et al. 2018).

The first phase-lag measurements were performed by Schultz et al. (1978),
while the first attempt at actual distance determinations were made in the
1980’s by Herman & Habing (1985) and van Langevelde et al. (1990).

The method has been recently revisited with the phase-lag measurements of a
sample of 20 such objects performed with the Nançay Radio Telescope (NRT)
(Etoka et al. 2018; Engels et al. 2015). A detailed analysis of the method,
and the identified caveats of its use, show that an accuracy of $<10$% in the
distance determination can be achieved (Etoka & Engels 2022). In order to
achieve this precision, a well-sampled, i.e., regularly spaced with an
adequate number of measurement points to determine the shape of the light
curves with enough precision to infer the delay is needed along with a high-
sensitivity interferometric map.

Single-dish monitoring of these objects alone not only provides the
determination of their phase-lag, and hence the linear diameter of their CSEs
but also an accurate measurement of their periods for free. Radio data are
crucial to determine the stellar pulsation periods of OH/IR stars which are
heavily obscured by their CSEs and invisible in the optical and even in near
infrared bands. To measure the period, it is essential that the monitoring
program shall cover at least a few pulsation periods (typically 5 years or
more). Taking plotting on period-luminosity diagrams of long period variables
as a target for the period determination accuracy to be achieved, 7 days is
enough for a star with a pulsation period greater than 500 days.

Beyond this, the monitoring cadence must be such that the light curves are
well sampled so that their shape is well defined as it is paramount for the
lag-determination. It should consequently be a fraction (i.e., ideally less
than 50%) of the lag to be measured.

Due to its location, the TNRT monitoring in single-dish mode will allow such
measurements for OH/IR stars at lower declinations than achievable with in
particular the NRT and hence in a complementarity fashion especially with
respect to the Galactic center.

To determine the distance of these objects punctual interferometric
observations are needed.

####  5.1.2 Maser Monitoring to Investigate the Shaping and Driving of Evolved
Star Winds

It is not well-understood how mass is lost from the atmospheres of cool,
evolved stars, although once dust has formed its role in driving the wind is
well-modelled (Bladh et al. 2015). Mass loss may be affected by exceptionally
deep pulsations, thermal pulses (Richards et al. 2012; Homan et al. 2021) and
episodic ejecta (Humphreys et al. 2021). Current mass loss models differ by
orders of magnitude (e.g. Saberi et al. 2019 ; Decin 2021). Subtle asymmetries
developing during the AGB lead to the shaping of post-AGB water fountains and
PNe, see section 5.2.2. Evolutionary changes may be seen even in a few years
e.g. (Vlemmings et al. 2014).

The kinematics inside $\sim$5 $R{\star}$, traced by SiO masers, are pulsation
dominated and complex (Bladh et al. 2019). Shocks further out in the 22-GHz
maser shell ($\sim$5$R{\star}$ to tens $R{\star}$) might be caused by rapid
expansion of denser regions leading to a pressure difference. 22 GHz masers
are concentrated in clumps where the dust-gas collision rate and heating are
more effective. OH mainline masers appear less strongly accelerated than
22-GHz maser clumps at apparently similar radii, consistent with tracing
cooler, less dense, interleaving gas (e.g. Richards et al. 1999). Very high-
velocity OH 1667 MHz and/or H2O 22-GHz emission is sometimes seen (exceeding
the 1612 GHz expansion speed), possibly related to exceptional, directed mass
loss. Some sources show an equatorial density enhancement and biconical
outflow as modelled by Zijlstra et al. (2001); potential origins include
magnetic fields and interactions with companions (planetary mass upwards) e.g.
Decin et al. (2020), Homan et al. (2020)

Maser variability arising from different causes will propagate away from the
star at different speeds depending on the cause: direct, radiative effects;
the timescale of heating and cooling (e.g. ionisation/recombination), shock
speed, the wind velocity and localised inhomogeneities. Out to 5 $R{\star}$
the light-travel time is $\sim$1 hr for an AGB star of radius 1 au but will be
slowed by factors of a few in IR-optically thick clumps. The timescale for
shock propagation at $\leq$10 km s-1 is $\leq$2.5 yr and the net wind outflow
timescale could be as long as 50 yr (Assaf 2018) due to non-linear
trajectories (see e.g. Harper & Linsky (2001) for timescales around RSG).
Further out, most O-rich winds follows radial streamlines with slow
acceleration. 22-GHz maser fluctuations may follow direct heating by the IR
period (lagging the optical) or via locally-induced shocks propagating through
clumps and density inhomogeneities at sound speed (up to years), or as long-
term variations in the mass loss rate affect the wind density (e.g.
’superperiods’, Lekht et al. 2001). Disruption of masing by planetary
companions around a solar-mass star would occur every (few) year(s) at a few
$R{\star}$. Monitoring OH and H2O close in time is needed to reveal features
with coordinated variability (ruling out projection effects) and investigate
the small-scale structure of the wind in the region where it attains escape
velocity (e.g. Colom et al. 2014). Monitoring will catch brief high-velocity
outflows and establish whether H2O activity is connected to extreme velocity
OH.

Monitoring tracks velocity drifts due to complex morphology and dynamics. A
red-blue velocity offset at one epoch may appear in an opposite sense at
another epoch hinting at asymmetric mass loss, or show velocity drifts
indicating a preferred axis of mass loss, distinguishing between an equatorial
density enhancement versus a rotating disc. For example, asymmetries were
deduced from single-dish monitoring for VX Sgr, Pashchenko & Rudnitskii
(1999), refined by interferometry (Murakawa et al. 2003).

Polarisation monitoring of SiO (Herpin et al. 2006) and OH (e.g. Wolak et al.
2013) will investigate the magnetic field dependence on radius (hence
distinguishing between dipole, solar and toroidal configurations) and its
variability and role in shaping the wind or confining clumps. There are also
polarisation phenomena at the clump scale, probed by VLBI, where rotations of
the electric vector position angle through 90 degrees, apparently within a
single clump, have been observed in TX Cam (Tobin et al. 2019) and R Cas
(Assaf et al. 2013). In-depth studies of specific objects e.g. examining the
effects of low-mass companions or the aftermath of a thermal pulse, can be
complemented by interferometry (sec 5.2). H2O also survives around objects
which have only recently become C-rich e.g. Zhao et al. (2020) and references
therein and SiO masers also occur around S-type (C/O $\sim$ 1) stars. The
89-GHz HCN maser is found close to C-rich stars (Smith et al. (2014) and
references therein) and, like SiO, shows rapid variability. Monitoring and
comparisons with O-rich stars are of great interest in examining how the
chemical composition of the star and nascent dust may fundamentally determine
mass loss dynamics and wind acceleration (e.g. Gottlieb et al. 2022).

Single-dish monitoring observations of circumstellar masers are one of the key
topics where the TNRT will have high impact. A database of OH, H2O and SiO
masers in O-rich CSEs has recently been compiled (Sobolev et al. 2019).
H2O-SiO studies have been made using Yonsei and KVN antennas (e.g. Kim et al.
2016). These cover several hundred lower-mass stars, mostly single-epoch. A
large-scale SiO maser survey such as the Bulge Asymmetries and Dynamical
Evolution (BAaDE) project (see e.g. Stroh et al. 2018; Lewis et al. 2020) has
observed over 28,000 colour-selected red giants in the Galactic plane to
explore the dynamical structures of the inner Galaxy, which provides a wealth
of information on the statistics of SiO maser transitions and their sources.
Selected targets can thus be investigated further with single-dish
observations. Long-term monitoring e.g. using Medicina and Pushchino at 22 GHz
or Nançay at 1.6 GHz is only coordinated for a few objects. Interferometric
observations are not practical for frequent monitoring of large samples, rapid
($<$few hr) variability or rapid responses and much flux is often resolved out
on 100s-km baselines.

Dedicated TNRT monitoring programs will provide consistent, multi-frequency
coverage of a large sample of evolved star masers, at 1 Jy sensitivity,
allowing the statistical analysis crucial for testing the hypothesis mentioned
above. This kind of intensive monitoring at three frequencies is globally
unique apart from a connected project using the Yebes 40 m telescope, where
$\sim$800 hours will be spent per year. The TNRT can observe stars down to
Dec. –55∘ at 15∘ elevation, for example covering almost all the ATOMIUM (Decin
et al. 2020) and NESS (Scicluna et al. 2022) samples, and giving easy access
to high mass stars concentrated towards the Galactic plane.

Sampling intervals are determined by sensitivity, rates of maser variability
and survival of distinct features, depending on location with respect to the
star, typically from a few weeks in the SiO maser zone to months at greater
distances, but the TNRT could also investigate poorly-understood hints at much
shorter maser flickering. 1-2 Jy sensitivity, sufficient to monitor bright SiO
and H2O sources, can be achieved in minutes. Detection experiments, and
polarimetric observations (also requiring high spectral resolution) may need a
few hr for individual targets. Significant monitoring results will emerge in
$<$ a year for short-period stars; several years or more will build up a
clearer picture for longer-period objects and reveal changes in objects
undergoing rapid evolution. Occasional interferometric imaging (e.g. per
stellar period) will locate features with respect to the star and inform the
interpretation of more frequent spectral monitoring (e.g. Richards et al.
2012) and can be coordinated with the TNRT in VLBI (Sec. 5.2).

####  5.1.3 Signposts/Signatures of transitional stages on the AGB

In optically-thin-CSE evolved stars, dramatic flaring events can occur in OH.
The first report of an OH maser flaring event was reported by Jewell et al.
(1981). A census and statistical analysis of such events was presented by
Etoka & Le Squeren (1997) solely on the basis of long-term monitoring which
constrained several aspects of these flaring events. These objects are all
Miras, with CSE properties putting them in the “(OH) maser” and “(OH) non-
maser” O-rich AGB transitional phase, hence heralding changes in the CSE. The
most important/striking characteristics are that this type of emission is
substantially polarised and that their velocities point to restricted
locations close to the star. The latter points were indeed confirmed recently
by VLBI mapping; additionally, confrontation between contemporaneous OH and
water maser monitoring indicates that OH flaring seems to be associated with
fainter water maser emissivity(Etoka et al. 2017).

While OH/IR stars show a very low percentage of polarisation (typically less
than 10%), highly polarised emission (up to 100%) is sometimes observed
towards objects with an intermediary-thick CSE, reminiscent of what is
observed toward the optically-thin-CSE flaring Miras. The strong polarisation
observed towards these objects higher up on the AGB, is suspected to be linked
with them being in the transitional phase Mira $\rightarrow$ OH/IR star (Etoka
& Le Squeren 2004). How this highly polarised emission (i.e., the influence of
magnetic field) is linked to the evolutionary changes undergone by the CSE
when an object evolved towards the [P]PN stage is not yet understood.

The TNRT capability of monitoring in both species transitions “quasi
simultaneously” and retrieve at the same time the polarimetric information,
will allow to investigate the possible relation between the OH and water
emissivity of transitional objects (i.e., both in the “early” Mira stage and
in the Mira $\rightarrow$ OH/IR star stage) so as to identify the CSE changes
that the OH maser flaring/strong polarisation “flashes” are the sign post of.

This requires a dual monitoring approach: a ”searching” low-cadence monitoring
(i.e., typically monthly or even every other month) of the selected candidates
(at a sensitivity of a few 100 mJy) and an ”event-recording” high-cadence
monitoring (i.e., as fast as a few days – weekly) towards those objects
showing the aforementioned transitional signatures.

Punctual inteferometric imaging is needed in order to localised the impacted
regions of the CSE while the monitoring itself provide the crucial dymanical
evolution of these events.

####  5.1.4 Episodic Events of Circumstellar Masers

Here they are defined as events of the CSE masers that cannot be explained by
repeated or stable patterns in spectra and interferometry maps, different from
those that are likely attributed to regular stellar pulsations and stable
radial accelerations in the CSEs. They may include episodic appearances and
flares of the masers on time scales shorter than a typical pulsation period of
long-period variable stars (100–2000 d). Some of them might be associated with
non-recurrent events such as a nova (e.g. V407 Cyg, Deguchi et al. 2011).
Including such an extreme case and recurrent events over a stellar pulsation
period, thye may be associated with binarity of an evolved stellar system
forming a “common envelope” such as a symbiotic star. Khouri et al. (2021)
suggest that a group of the stellar systems hosting high velocity ($\geq$100
km s-1) H2O maser sources, so-called “water fountain” sources (WFs, e.g.,
Gómez et al. 2017) should be experiencing such a common envelope evolution
during a very short period ($\leq$100 yr). Therefore, it is also expected to
see a secular evolution of SiO/H2O/OH maser sources in the same stellar group
over a human lifetime, namely a discovery of a new maser source and a
permanent death of an existing maser source though the evolution of a fast
bipolar molecular jet (e.g. Tafoya et al. 2020).

One of the science goals in this topic is to reliably determine the timescales
of key events described above, namely the period of the recurrent events and
total duration of the WF phase. Catching new highest velocity components of
H2O masers associated with such events may need a high cadence (an epoch
spacing of 1–2 weeks) monitoring program such as FLASHING (Finest Legacy
Acquisitions of SiO-/H2O-maser Ignitions by Nobeyama Generation, Imai et al.
2020). Tracing decadal evolutions of the individual WFs, even a small number
($\sim$15) of such sources, is a tough work for an observatory because of the
weakness (as weak as 0.1 Jy) of the H2O masers, especially the highest
velocity components. Detections of SiO ($v=1$ $J=1\rightarrow 0$) masers from
the WFs (Imai et al. 2005; Amada et al. 2022) are challenging because of their
rareness and faintness (as weak as 0.2 Jy) but scientifically brilliant
because they should locate the central stellar system and trace the spatio-
kinematics of the root of the WF jet. Eventually, such single-dish monitoring
programs need collaborations of large telescopes including TNRT, similarly to
M2O (Maser Monitoring Organization).

It has been expected that the episodic event of H2O masers should be
associated with a new launch of a fast jet (Imai et al. 2020) and/or an
intercation between the jet decelerated and the ambient CSE entrained by the
jet (Tafoya et al. 2020). This may be able to be traced by finding some
systematic trends of line-of-sight velocity drifts in the H2O maser spectral
peaks. Breaking a record of the jet’ top speed will improve our understanding
about how deep gravitational potential can launch such a fast jet.

Definitely, collimated jets found in WFs should be driven by a magneto-
hydrodynamical force (Vlemmings et al. 2006). Measurement of linear
polarization and the Zeeman effect of H2O masers should provide key
information on the configuration and the strength of the magnetic field in the
jet (e.g. Vlemmings 2019). Because the degree of circular polarization is
quite low ($<<$1%), a spectral stacking using the monitoring data may be
useful for the polarization detection.

###  5.2 Circumstellar Dynamics and Galactic Kinematics Revealed in VLBI

Thanks to compactness of the individual features of CSE masers and their
associations with unique objects in astrophysics, they are interesting targets
for VLBI animation synthesis and high-accuracy astrometry. Note that, because
of dominant flux contributions from more extended structures of the CSE
masers, their mapping also needs relatively short baselines (100–1000 km).

####  5.2.1 Propagation of Shock and Heat Waves in CSEs

Animation synthesis of CSE maser sources is a much more straightfoward way to
reveal the dynamics of the CSE than the spectroscopic monitoring mentioned
above (e.g. Gonidakis et al. 2013). One of the key issues about the CSEs is
the mechanism of a stellar wind acceleration that may be enhanced by
propagation of pulsation-driven shock waves. Another is the origin of
asymmetric mass loss. The maser animations definitely visualize the possible
accelerations/decelerations of maser gas clumps (e.g. Imai et al. 2003) and
will shed light on these issues.

However, the change in the physical conditions of maser excitation should be
taken into account when intepreting the observed spatio-kinematics. This is
tracable in only simultaneous monitoring of multiple maser lines, especially
SiO masers ($J=1\rightarrow 0$, $2\rightarrow 1$, and $3\rightarrow 2$)
located closely to each other so as to shape a common maser ring (e.g. Yoon et
al. 2018). Based on these motivations, large projects for intensively
monitoring CSE masers have been conducted. As their first phase, the KaVA
ESTEMA (Expanded Study on Stellar Masers)151515See KaVA Large Programs website
at https://radio.kasi.re.kr/kava/large_programs.php. was conducted during
2015–2018, aiming to statistically explore CSE masers towards $\sim$80 stars
and find the targets suitable for the intensive VLBI monitoring (Imai 2018).
The second phase of ESTEMA (EAVN Synthesis of Stellar Maser Animations) has
realized the required intensive (a time spacing shorter than 1/20 of the
stellar pulsation period) VLBI monitoring observations of H2O and SiO masers
in four frequency bands (22/43/86/129 GHz) for two stars with different
periods (BX Cam and NML Cyg) since 2018.

The number of telescopes participating in the monitoring is always a key
factor for its success. Taking over six hours per mapping session every two
weeks is challenging in terms of time allocation and scheduling. Imaging
simulations (Imai et al. in private communication) suggest that a snapshot
image needs at least a few consecutive hours with seven telescopes of KaVA.
Unless the EAVN telescopes, except KVN, can simultaneously observe H2O and SiO
masers, the requested hours should be doubled as seen in ESTEMA for the first
two stars. Even KaVA can fully provide seven telescopes in a limited fraction
of the whole year. Therefore, other EAVN telescopes including TNRT are always
desired to join such monitoring observations, especially for the shorter-
period Mira variables (including Mira a.k.a. $o$ Ceti) with pulsation periods
$\leq$350 d.

####  5.2.2 Ignitions and Secular Evolutions of Water Fountain jets/flows

Mapping H2O (especially very high velocity components) and SiO masers
associated with WFs is challenging without large telescopes such as TNRT.
Internal proper motions of the maser features in WFs will determine the speed
and the kinetic timescale of the jet. These parameters are a key to elucidate
the mechanism of the jet launch and predict the final CSE evolution leading to
the variety of shapes observed in planetary nebulae. The spatio-kinematics of
the H2O masers have been revealed for most of the WFs with interferometers.
However, it is crucial to revisit those sources. In subsequent interferometric
observations, we may able to locate new maser components that are possibly
associated with the episodic events mentioned above. We also may be able to
trace the secular evolution of the WFs over decades (e.g. Imai et al. 2020).

####  5.2.3 The Astrometry for CSE Masers around the Galactic Center

Trigonometric observations of CSE masers provide a strong support for the
determination of physical parameters, such as the stellar luminosity and the
mass-loss rate, of the accompanying stars. In addition, the period–luminosity
relation of long-period variable stars in the Milky Way Galaxy will be more
accurately determined with dozens of the trigonometric and infrared
photometric samples. All of OH/H2O/SiO masers can be used for the
trigonometry. OH masers, in particular, will become the main targets in the
SKA era for high accuracy astrometry (Rioja & Dodson 2020). Orosz et al.
(2017) demonstrated such astrometry for CSE OH masers. TNRT may build a VLBI
intermediate length ($\sim$1100 km) baseline with FAST in L(1.6 GHz)-band.
With other EAVN telescopes such as TMRT, XAO, Usuda, one can promote the OH
maser trigonometry.

The Galactic Center is a unique region including the Galactic Bulge and the
Nuclear Stellar Disk and Bulge, whose mass distribution and formation
histories will be elucidated through a large sample of the three-dimensional
velocities of associating stars. CSE maser sources associated with these stars
should be interesting astrometric targets because they are obscured by clouds
in the Galactic Plane and some of them are invisible in optical and near-
infrared. TNRT has a good access to the Galactic Center.

###  5.3 Thermal SiO Line and Stellar Continuum Emissions

The depletion of SiO into dust formation or later accretion has been variously
estimated at ’almost all’ to $\sim$50% (e.g. Sahai & Bieging 1993 and
references therein), as the dust chemistry is rather uncertain (e.g. Liu &
Jiang 2017). Accurate estimates of the mass of gas-phase SiO provide vital
constraints but the higher-J transitions, whilst good tracers of the local
radiation field and shocks, are thus less reliable for density measurements of
gas extending further from the star. The low surface brightness makes thermal
$J=2\rightarrow 1$ $v=0$ inaccessible to current interferometers and
$J=1\rightarrow 0$ is barely studied even by single dish.

In some objects, $J=5\rightarrow 4$ and $J=6\rightarrow 5$ thermal SiO
transitions reach greater apparent outflow velocities than CO, although they
have a smaller angular size, (Gottlieb et al. 2022), attributed to the wind
interaction with a companion. Measuring the $J=2\rightarrow 1$ and
$J=1\rightarrow 0$ velocities will investigate whether this is influenced by
excitation effects.

The TNRT is unique in having the sensitivity to search for the two lowest
thermal SiO transitions simultaneously, measuring the non-dust fraction and
revealing any velocity anomalies. In order to detect the line wings, a few hr
per star is needed for a small sample, providing immediate science results by
comparison with known dust properties. One repetition at a different stellar
phase would identify any maser contamination (weak but sometimes present in
$v=0$). These observations will provide an abundance profile for objects with
independent higher-J data and the kinematics can be compared with CO
properties (e.g. from the NESS and Atomium surveys).

Betelgeuse is 10–14 mJy at K-band and 28 mJy at Q-band, radio spectral index
$\sim$1.4. It is continuously monitored optically but only occasionally at
other wavelengths (e.g. O’Gorman et al. 2015). Radio emission represents the
flux density at the optically thick surface at that radius, increasing with
decreasing frequency (e.g. the C-band radius is 4-6 times the optical stellar
radius). The K-band emission is thought to emanate from just outside the
region where pulsation shocks are damped whilst higher frequencies come from
nearer the optical photosphere. In 2019 optical dimming appeared to arise from
an exceptional mass ejection, such that Betelgeuse was (partly) obscured by a
dust formation event and/or cooled due to internal processes (e.g. Montargès
et al. 2021; Dharmawardena et al. 2020 Dupree et al. 2022). K-band emission
would be unaffected by dust formation further from the star but unusual
variability (monitored every few months) would indicate internal changes,
potentially related to exceptional convective events. (e.g. Harper & Brown
2006) Coordinated K-Q-W band monitoring will be invaluable in measuring the
passage of disturbances out from the photosphere, changes in temperature and
propagation mechanisms (e.g. Harper & Linsky 2001).

The TNRT could be unique in providing radio continuum monitoring of the
closest supernova precursor. Typical radio fluctuations, whether due to shocks
or convection or other causes, are $\sim$10%, requiring $\sim$1 hr per visit,
including accurate flux scale calibration. Weekly or fortnightly monitoring
during the favourable months would sample $\sim$half Betelguese’s 400-day
period (predicted to re-appear following the Great Dimming). Moreover,
piggyback observations at high spectral resolution will reveal if the mass
loss event initiates H2O or SiO masers (seen weakly at shorter wavelengths).

##  6 Radio Survey and Monitoring of Emission in CP Stars

Led by Eugene Semenko & David Mkrtichian

In the upper Main Sequence, about 10-15% of A and B stars showing the
abnormally strengthened or weakened lines of some indicative elements in their
spectra are recognized as chemically peculiar Ap/Bp or CP stars. This class of
hot stars was divided into four subgroups according to the most characteristic
chemical elements (Preston 1974). Chemical anomalies in such stars are the
result of atomic diffusion carrying out in stabilized atmospheres. It is
believed that the slow axial rotation, tidal effects, and magnetic field are
the main factors of stabilization. The magnetic field plays a special role.

In spectropolarimetric observations, about 10-15% of all CP stars show the
presence of a strong magnetic field which often has a simple dipolar or
multipolar configuration and covers the whole surface of the star (Donati &
Landstreet 2009). The strength of the magnetic field of CP stars ranges from
few hundreds to above thirty thousand gauss. In contrast to the Sun and other
cool stars, the magnetic field of CP stars remains stable on a timescale of
decades, albeit the observed characteristics of the field vary with rotation.

A synchronous photometric, spectral and magnetic variability is a major
feature of peculiar A and B stars. This phenomenon was explained by Stibbs
(1950) within the framework of a model of oblique rotator. Within this model,
the magnetic dipolar or multipolar field is frozen in the stellar plasma and
rotates with star. The magnetic axis is inclined to the rotational axis of the
star at the angle $\beta$, while the rotational axis is inclined to the line
of sight at the angle $i$ (12). According to the study by Landstreet & Mathys
(2000), the angle $\beta$ in the slowly rotating magnetic stars with period
$P_{\mathrm{rot}}>25^{\mathrm{d}}$ tends to be close to zero, i.e the magnetic
and rotational axes are almost co-aligned. The explicit configuration
determines the character and amplitude of observed indicators.

Figure 12: Schematic view of a magnetic oblique rotator.

Apart from a regular variability, chemically peculiar stars are mostly quiet
in all spectral domains. An exception is a fraction of the hottest CP stars.
In their case, a strong magnetic field in combination with relatively fast
rotation and powerful stellar winds forms the environment of extended
magnetospheres with magnetically confined winds shocks. This phenomenon leads
to the variable emission in UV, visible, and IR spectra. In some extreme cases
the enhanced signal is observable in X-rays (Nazé et al. 2014). Strongly
magnetized atmospheres of Bp stars are also a suitable environment for certain
types of radiation in radio waves.

Radio emission of early type stars was discovered in 1980s, and in the
subsequent years the number of identified sources grew very fast. In a paper
by Leone et al. (1994), the authors summarized the occurrence of radio sources
among different types of CP stars, specifically, 31% for helium rich stars,
26% for helium weak stars, 23% for silicon stars, and zero — for cooler Ap
stars. The radio emission in these cases usually varies with a rotational
period of star, and its amplitude and intensity depend on the wavelength.
According to the 3D model proposed by Trigilio et al. (2004), gyro-synchrotron
radio continuum emission arises from the interaction of stellar winds with a
magnetic field. Measured radio flux density is determined by the geometry of a
magnetosphere and its physical properties. Radio emission at higher
frequencies comes from the layers located closer to the stellar surface. Thus,
the authors emphasise the importance of multifrequency observations. Such
example of a work is a study of hot B star HR 7355 by Leto et al. (2017).
Observing this star with VLA at frequencies 6, 10, 15, 22, 33, and 44 GHz, the
authors found variable non-polarized radio flux order of 10–25 mJy with an
amplitude decreasing at higher frequencies.

In addition to the incoherent radio emission described above, some magnetic CP
stars exhibit the radio signal of different nature. This type of emission is
characterised by the sharp circularly polarized signals appearing in a narrow
range of rotational phases, with fluxes which are almost one order higher than
in the case of gyro-synchrotron mechanism. The mechanism of excitation here is
similar to Auroral Radio Emission (ARE) observed on some magnetized planets. A
broadband manifestation of ARE is observable in all domains, including X-rays.
In radio, ARE appears through the mechanisms of Electron Cyclotron Maser
Emission (ECME). A model of this process in hot stars was given by Melrose &
Dulk (1982) and Leto et al. (2016). Unlike the incoherent gyro-synchrotron
emission which intensity resembles the curve of a magnetic field, sharp
impulses of ECME appear only when the star is observed from its magnetic
equator.

For the first time, ECME was detected in a famous CP ’pulsar’ star CU Vir by
Trigilio et al. (2000). Even the number of known sources of ECME among CP
stars had been growing in the consequent years, by the end of September 2020,
only seven stars showed variable coherent radio emission in observations.
These stars are: CU Vir (Trigilio et al. 2000), HD 133880 (Chandra et al.
2015; Das et al. 2018), HD 142990 (Lenc et al. 2018; Das et al. 2019a), HD
142301 (Leto et al. 2019), HD 35298 (Das et al. 2019b), $\rho$ Oph A and C
(Leto et al. 2020a, b). All items in this list are hot magnetic stars with the
effective temperature 12–18 kK and rotational periods shorter than 2 days. For
those rare stars where both types of radio emission are observed
simultaneously, one can possible to probe the structure of circumstellar
envelopes. In a paper by (Leto et al. 2020a) the authors presented the model
of radio emission of the star $\rho$ Oph C. To demonstrate it, we reproduce
here Fig. 4 of the original paper (Fig. 13).

Figure 13: Radio map of a star $\rho$ Oph C with indicated sources of radio
emission and corresponding frequencies. Estimated size of the middle-
magnetosphere is 15 $R_{*}$. Originally published in Leto et al. (2020a),
reproduced with permission of the authors.

In the Main Sequence, magnetic CP stars probably are the only a group of
objects where multiwavelength observations open a way to the exploration of
entire structure of the star, from the core to the nearest surroundings. With
the introduction of a new radio telescope, NARIT can extend its study of
chemically peculiar stars in a new dimension. In addition to the photometric
and spectroscopic observations, one become possible to monitor the physical
conditions in magnetic CP stars in radiowaves.

We propose the program for the Thai National Radio Telescope aimed at the
searches and monitoring of radio emission of selected CP stars in L- and
K-bands. Through this research, we plan to constrain the configuration of the
magnetic field derived from photometric and spectropolarimetric observations,
to probe the presence of magnetospheres which are not recognizable in optical
spectroscopy, and to explore their properties in the regions where the radio
emission originates.

In Table 10 we selected mostly bright (and, thus, close) late B and early A
candidate stars with notable peculiarities of Si and He-wk type (Renson &
Manfroid 2009) with available curves of a magnetic field and TESS photometry.
The photometric light curves give the accurate values of rotational period
(column 8 of the table) and their shape indicates the level of homogeneity of
the surface. Rotational periods all of these stars are below or around two
days, excepting HD 40312 where $P_{\mathrm{rot}}$ is 3.618 days. The reason
why this star was included is the Lorenz force, detected for this star by
Shulyak et al. (2007). The longitudinal magnetic field of all listed stars
changes its polarity and has a well determined curve, sometimes, like in the
case of HD 34736, of a non-sinusoidal shape (Fig. 14). Extrema of the
longitudinal field are presented in column 7 of the Table.

Figure 14: A curve of the longitudinal magnetic field of HD 34736 phased with a rotational period $P_{\mathrm{rot}}=1.\\!\\!^{\mathrm{d}}2799$. Table 10: List of chemically peculiar stars, possible sources of radio emission. Star ID | $\alpha_{\mathrm{J2000}}$ | $\delta_{\mathrm{J2000}}$ | Sp | $m_{\mathrm{V}}$ |  $\pi$, mas |  $\langle B_{l}\rangle$, kG |  $P_{\mathrm{rot}}$, d
---|---|---|---|---|---|---|---
HD 11503 | $01^{\mathrm{h}}53^{\mathrm{m}}31.\\!\\!^{\mathrm{s}}8$ | $+19^{\circ}17^{\prime}\,37.\\!^{\prime\prime}9$ | A1 SiCrSr | 4.5 | 19.98 | $-1/+0.5$ | 1.6083
HD 12447 | $02\phantom{{}^{\mathrm{h}}}02\phantom{{}^{\mathrm{m}}}02.8$ | $-02\phantom{{}^{\circ}}45\phantom{{}^{\prime}\,}49.6$ | A2 SiSrCr | 4.1 | 19.79 | $-0.5/+0.4$ | 1.4326
HD 12767 | $02\phantom{{}^{\mathrm{h}}}04\phantom{{}^{\mathrm{m}}}29.4$ | $-29\phantom{{}^{\circ}}17\phantom{{}^{\prime}\,}48.5$ | A0 Si | 4.7 | 8.79 | $-0.25/+0.25$ | 1.8897
HD 19832 | $03\phantom{{}^{\mathrm{h}}}12\phantom{{}^{\mathrm{m}}}14.2$ | $+27\phantom{{}^{\circ}}15\phantom{{}^{\prime}\,}25.1$ | B8 Si | 5.8 | 8.04 | $-0.3/+0.4$ | 0.728
HD 22470 | $03\phantom{{}^{\mathrm{h}}}36\phantom{{}^{\mathrm{m}}}17.4$ | $-17\phantom{{}^{\circ}}28\phantom{{}^{\prime}\,}01.5$ | B9 Si | 5.2 | 7.65 | $-1.6/+1.9$ | 1.9298
HD 30466 | $04\phantom{{}^{\mathrm{h}}}49\phantom{{}^{\mathrm{m}}}16.0$ | $+29\phantom{{}^{\circ}}34\phantom{{}^{\prime}\,}16.8$ | A0 Si | 7.3 | 6.51 | $-0.25/+2$ | 1.4062
HD 34736 | $05\phantom{{}^{\mathrm{h}}}19\phantom{{}^{\mathrm{m}}}22.2$ | $-07\phantom{{}^{\circ}}20\phantom{{}^{\prime}\,}50.0$ | B9 Si | 7.8 | 2.75 | $-5/+5$ | 1.2799
HD 36526 | $05\phantom{{}^{\mathrm{h}}}32\phantom{{}^{\mathrm{m}}}13.1$ | $-01\phantom{{}^{\circ}}36\phantom{{}^{\prime}\,}01.7$ | B8 Si/He-wk | 8.3 | 2.44 | $-1/+4$ | 1.5417
HD 37017 | $05\phantom{{}^{\mathrm{h}}}35\phantom{{}^{\mathrm{m}}}21.9$ | $-04\phantom{{}^{\circ}}29\phantom{{}^{\prime}\,}39.0$ | B2 He | 6.6 | 0.45 | $-2/0$ | 0.9012
HD 40312 | $05\phantom{{}^{\mathrm{h}}}59\phantom{{}^{\mathrm{m}}}43.3$ | $+37\phantom{{}^{\circ}}12\phantom{{}^{\prime}\,}45.3$ | A0 Si | 2.6 | 19.7 | $-0.3/0.3$ | 3.6181
HD 45583 | $06\phantom{{}^{\mathrm{h}}}28\phantom{{}^{\mathrm{m}}}10.7$ | $-04\phantom{{}^{\circ}}53\phantom{{}^{\prime}\,}56.5$ | B9 Si | 7.9 | 3.06 | $-3/+4$ | 1.1770
HD 63347 | $07\phantom{{}^{\mathrm{h}}}54\phantom{{}^{\mathrm{m}}}20.4$ | $+71\phantom{{}^{\circ}}04\phantom{{}^{\prime}\,}45.2$ | B8 SrCrEu | 7.3 | 5.66 | $-0.8/+1$ | 1.7495
HD 108945 | $12\phantom{{}^{\mathrm{h}}}31\phantom{{}^{\mathrm{m}}}00.6$ | $+24\phantom{{}^{\circ}}34\phantom{{}^{\prime}\,}00.2$ | A3 SrCr | 5.4 | 11.99 | $-0.3/+0.25$ | 2.0513
HD 170000 | $18\phantom{{}^{\mathrm{h}}}20\phantom{{}^{\mathrm{m}}}45.4$ | $+71\phantom{{}^{\circ}}20\phantom{{}^{\prime}\,}16.1$ | A0 Si | 4.2 | 10.77 | $-0.3/+0.5$ | 1.7165

Given the information about the magnetic field and stellar rotation from
spectroscopy and spectropolarimetry, a provisional observational strategy
might look as follows. To detect the coherent radio emission from candidate
stars, the observations must cover the time around the moments when the
longitudinal field changes the sign and carry out in L-band. Initially,
observations with TNRT will be made to focus on monitoring of flux density
without polarization studies. These monitoring should be continuously in whole
a day covering the periods of each target source. Sources in which the radio
emissions are detected with TNRT will be followed-up by another receiver at
phase 1, in K-band, and a forthcoming polarization mode on TNRT to study their
magnetic field activities. Unlike ECME, incoherent emission reach its maximum
when the star exposes its magnetic poles. In terms of the model of oblique
rotator, this corresponds to the maxima of the curve of the longitudinal field
$\langle B_{\mathrm{z}}\rangle$. The signal of incoherent gyro-synchrotron
emission is expected in both, L and K, radio bands. Because of the
corresponding moments of time can be predicted from known ephemerides, the
telescope time can be used in more efficient way.

##  7 Geodesy: Geodetic VLBI Observation in K-band

Led by Nattaporn Thoonsaengngam and Koichiro Sugiyama

Geodesy is the Earth science that has objectives of accurate measuring to
determine Earth’s geometric shape and its change with times. That includes the
orientation of the Earth in the spaces and the gravitational field. To obtain
expected results nowadays, sources external to the Earth become references of
the measurements, artificial satellites, retroreflectors on the Moons, and
also quasars, these techniques are so-called the space geodesy.

In space geodesy, the extragalactic or quasars are measured with the VLBI
technique to obtain baseline length between two (or more) radio telescopes
that operate the measurement with millimetre-accuracy. The results can be used
to determine the station coordinates that are contributed to the international
terrestrial reference frame (ITRF). VLBI is the only one of the space geodetic
techniques referenced to the kinematically non-rotating frame and the only one
that provides all Earth Orientation Parameters (EOPs), i.e., polar motion,
universal time, and celestial pole offsets (Heinkelmann 2008). This lead to
global objectives; obtaining precise positioning, monitoring plate tectonics,
etc. In addition, VLBI technique achieves residual delays of a few mm or
better, which are essential orders of accuracy to measure parallaxes for
astronomical sources at beyond 10 kpc distances. This kind of geodetic VLBI
observations has been organized as “The International VLBI Service for Geodesy
and Astrometry” (IVS)161616See IVS website at https://ivscc.gsfc.nasa.gov/.

Since the 1970s, radio telescopes with the legacy frequencies, S- and X-band
receivers have been using to perform geodetic VLBI observations. However, many
reference sources at the S and X frequencies natively have extended structures
which can lead to significant errors in VLBI observations. So that, geodetic
VLBI at higher frequencies are considered for reduction of sources core-shift.
This will allow construction of the more accurate celestial reference frame
(CRF), which will be a benefit to the higher accuracy of geodetic VLBI
observables. The CRF at higher frequencies is also advantageous for tying the
radio reference frames to optical reference frames such as GAIA (Mignard et
al. 2016). In 2006, four stations in Australia performed the 1.5-hour long
VLBI fringe test experiment at 22 GHz in K-band (Petrov et al. 2009). Results
from the test encouraged the scientists to run a full-scale 22-GHz for
geodetic VLBI. K-band experiments, afterwards also showed the good results
since the ionospheric contribution at the K-band is one order of magnitude
smaller than at the X-band. The ionospheric contribution can be negligible
during the period of low solar activity. These characteristics make results
from K-band comparable or even better than S- and X-band observations.

The TNRT will contain the K-band receiver thus, the radio telescope could
participate in geodetic VLBI observations at this band. Initial station
coordinates for the VLBI observing sessions are determined from the relation
of positions between the telescope reference point and a co-location GNSS base
station on site. The GNSS station can provide the coordinates with an accuracy
up to centimetre level. By performing the geodetic VLBI, higher accuracy up to
millimetre level of station coordinates can be obtained. Results of the
geodetic VLBI will be contributed to the International Terrestrial Reference
Frames (ITRF) to produce the EOPs and other geodetic components, such as
velocities of station coordinates which are important for studies of tectonic
motions. The precise station coordinates, also benefit usages of the TNRT not
only for astrometry or geodesy but astronomy, especially observations at
higher frequencies. Nonetheless, TNRT will benefit astrometry by enhancing the
K-band CRF. In particular, its location strongly contributes to create a nice
composition of baselines with Hobart 26-m (Australia) and HartRAO 26-m (South-
Africa) for achieving more effective geodetic VLBI observations at K-band in
the southern hemisphere following goals of the next generation beyond ICRF-3
(Jacobs et al. 2014), where the number of sources is still a factor 2 less
especially in he far-south (declination $\leq$ $-30^{\circ}$) and while the
number of sessions per source are roughly the same in the north and the south,
the average number of observations per source is also factor of 2 less in the
far-south, although the number of stations and their specifications have been
improved and upgraded mainly at S/X-bands in Southern hemisphere (Charlot et
al. 2020; de Witt et al. 2021).

## Acknowledgements

This work is based on the national flagship project entitled Radio Astronomy
Network and Geodesy for Development (RANGD) by NARIT. We wish to thank all of
the members of the International Technical Advisory Committee (ITAC) and the
International Scientific Advisory Committee (ISAC) for their supportive and
fruitful suggestions and advise, in which ITAC consists of Hideyuki Kobayashi
(Chair, NAOJ), Busaba H. Kramer (Secretariat, MPIfR/NARIT), Do-Young Byun
(KASI), Francisco P. Colomer (JIVE), Michael Garrett (JBCA), Yashwant Gupta
(NCRA), Mareki Honma (NAOJ), Jinling Li (SHAO), Young Chol Minh (KASI),
Zhiqiang Shen (SHAO), Tasso Tzioumis (CASS), Pablo de Vicente (IGN), and
Gundolf Wieching (MPIfR), while ISAC consists of Michael Bode (Chair, BIUST),
Busaba H. Kramer (Secretariat, MPIfR/NARIT), Hideyuki Kobayashi (NAOJ), and
Michael Kramer (MPIfR), respectively.

We would like to thank all of engineers and technicians in Radio Astronomy
Operation Center (RAOC) for their sincere contributions to construct and
develop 40-m TNRT: Apichat Leckngam, Kamorn Bandudej, Nonwarit
Borvornsareepirom, Attapon Bunwong, Teep Chairin, Nattawit Chanwedchasart,
Pathit Chatuphot, Awut Duangkhiaw, Nattapong Duangrit, Nutdanai Hantankul,
Prachayapan Jiraya, Kitipoom Kanjana, Pitak Kempet, Natthaphong Kruekoch,
Peerapol Meekhun, Nakornping Namkham, Thanadon Paksin, Attasit Phakam, Natee
Pongteerarat, Anya Poonnawatt, Nikom Prasert, Sothaya Prathumsub, Songklod
Punyawarin, Haseng Sani, Abraham Sanenga, Spiro Sarris, Nonnadda Silamai, Dan
Singwong, and Lalida Tantiparimongkol.

All the editors and authors should like to our thanks to the the East Asian
VLBI Network (EAVN) for their providing an observation data towards M87 at
K-band in 2017 as a model for the simulation shown in figure 9 of section 4.1.
The EAVN is operated under cooperative agreement by National Astronomical
Observatory of Japan (NAOJ), Korea Astronomy and Space Science Institute
(KASI), Shanghai Astronomical Observatory (SHAO), Xinjiang Astronomical
Observatory (XAO), Yunnan Astronomical Observatory (YNAO), and National
Geographic Information Institute (NGII), with the operational support by
Ibaraki University (for the operation of Hitachi 32-m and Takahagi 32-m radio
telescopes), Yamaguchi University (for the operation of Yamaguchi 32-m radio
telescope), and Kagoshima University (for the operation of VERA Iriki
antenna). We also thank Toshio Terasawa (The University of Tokyo) for his
fruitful advice and discussion in section 2. The simulations in figure 9 were
generated by using the Astronomical Image Processing System (AIPS: van Moorsel
et al. 1996; Greisen 2003) and the Difmap software (Shepherd 1997).

Simultaneously, we would like to thank for a forthcoming collaboration to the
Australia Long Baseline Array (LBA), which is part of the Australia Telescope
National Facility, which is funded by the Australian Government for operation
as a National Facility managed by the Commonwealth Scientific and Industrial
Research Organisation (CSIRO), and to the European VLBI Network as well, which
is a joint facility of independent European, African, Asian, and North
American radio astronomy institutes.

K.S. is supported by the grant of PIIF Heiwa Nakajima Foundation in 2019 and
MEXT/JSPS KAKENHI (or Grant-in-Aid for Scientific Research) Grant Numbers
JP19K03921. K.S., P.J., and K.A. are supported by the research grants of NARIT
(Public Organization) under MHESI (Ministry of Higher Education, Science,
Research and Innovation) Grant Numbers 75249, 75253, 76820, 30137, 30141,
30146, 30152, 30156, 88313, 88314, 88319, 88320, 88321, 88322, 88324, 168577,
168593, and 168595.

##  Appendix A UV-coverages and Specifications Expected in World-wide VLBI
Collaborations

In this appendix, we summarize results of simulations for uv-coverages in the
cases that TNRT will collaborate with the LBA and the EVN in L-band and K-band
towards astrnomical sources in the Northern and the Southern hemispheres, as
shown in figures 15 and 16, respectively, while the simulation results for the
collaboration with the EAVN are shown in figure 3 of section 1.2. Besides,
specifications achieved in these VLBI collaborations: baseline lengths and
baseline sensitivities as 5$\sigma$, are calculated for the cases of
EAVN$+$TNRT (table 11 in K-band), LBA$+$TNRT (tables 12, 13 in L-/K-bands),
and EVN$+$TNRT (tables 14, 15 in L-/K-bands), respectively. These simulationas
and calculations were achieved on the basis of station coordinates, array
configurations, and specifications for each telescope in each frequency band
released via web-pages of each array and their status reports 171717EAVN
Status Report, see tables 1 and 4 there:
https://radio.kasi.re.kr/status_report/status_report.php$?$site=eavn,
181818LBA, see “Available Telescopes” and “Telescope Parameters”:
https://www.atnf.csiro.au/vlbi/documentation/, 191919EVN Capabilities, see
“EVN Status Table” as well: https://www.evlbi.org/capabilities and/or the
station coordinates registered in the SCHED program202020SCHED:
http://www.aoc.nrao.edu/software/sched/ organized for VLBI scheduling
generally led by Craig Walker at National Radio Astronomy Observatory. These
uv simulations here were generated by using the SCHED program as well. For the
Warkworth 30-m radio telescope in New Zealand, its station coordinates were
referred from Woodburn et al. (2015), and we assumed that a system noise
temperature and an aperture efficiency of this telescope in K-band was 200 K
and 30%, respectively, on the basis of information its receiver was in room
temperature. The specifications of TNRT in L-/K-bands are summarized in table
2 of section 1.1.

Figure 15: Simulated results of uv-coverages for LBA$+$TNRT in the case of
observing sources with declination of $-$29 (toward the Galactic Center) and
$+$10 degree in L-band (top) and K-band (bottom), respectively, with an
elevation limit of 10 degree. The unit of horizontal and vertical axes is
kilometer. Contributions of TNRT are highlighted by a color of red.

Figure 16: Simulated results uv-coverages for EVN$+$TNRT in the case of
observing sources with declination of $+$60 and $+$20 degree in L-band (top)
and K-band (bottom), respectively, with an elevation limit of 10 degree. The
unit of horizontal and vertical axes is kilometer. Contributions of TNRT are
highlighted by a color of red.

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