# Systematic Search for Water Fountain Candidates using
the Databases of Circumstellar Maser Sources

Haichen Fan(范海辰) School of Physics and Astronomy, Sun Yat-sen University, Tang
Jia Wan, Zhuhai, 519082, P. R. China  Jun-ichi Nakashima(中岛淳一) School of
Physics and Astronomy, Sun Yat-sen University, Tang Jia Wan, Zhuhai, 519082,
P. R. China  CSST Science Center for the Guangdong-Hong Kong-Macau Greater Bay
Area,
Sun Yat-Sen University, 2 Duxue Road, Zhuhai 519082, Guangdong Province, P. R.
China  D. Engels Hamburger Sternwarte, Universität Hamburg, Gojenbergsweg 112,
D-21029 Hamburg, Germany  Yong Zhang (张泳) School of Physics and Astronomy, Sun
Yat-sen University, Tang Jia Wan, Zhuhai, 519082, P. R. China  CSST Science
Center for the Guangdong-Hong Kong-Macau Greater Bay Area,
Sun Yat-Sen University, 2 Duxue Road, Zhuhai 519082, Guangdong Province, P. R.
China  Jian-Jie Qiu (邱建杰) School of Physics and Astronomy, Sun Yat-sen
University, Tang Jia Wan, Zhuhai, 519082, P. R. China  Key Laboratory of
Modern Astronomy and Astrophysics (Nanjing University), Ministry of Education,
Nanjing 210093, Jiangsu Province, P. R. China  CSST Science Center for the
Guangdong-Hong Kong-Macau Greater Bay Area,
Sun Yat-Sen University, 2 Duxue Road, Zhuhai 519082, Guangdong Province, P. R.
China  Huan-Xue Feng (冯焕雪) School of Physics and Astronomy, Sun Yat-sen
University, Tang Jia Wan, Zhuhai, 519082, P. R. China  Jia-Yong Xie (谢嘉泳)
School of Physics and Astronomy, Sun Yat-sen University, Tang Jia Wan, Zhuhai,
519082, P. R. China  Hiroshi Imai (今井裕) Center for General Education,
Institute for Comprehensive Education,
Kagoshima University 1-21-30 Korimoto, Kagoshima 890-0065, Japan  Amanogawa
Galaxy Astronomy Research Center (AGARC), Graduate School of Science and
Engineering,
Kagoshima University 1-21-30 Korimoto, Kagoshima 890-0065, Japan  Chih-Hao
Hsia (夏志浩) The Laboratory for Space Research, Faculty of Science,
The University of Hong Kong, Cyberport 4, Hong Kong SAR, China

###### Abstract

Water fountains (WFs) are thought to be objects in the morphological evolution
of the circumstellar envelopes of low- and intermediate-mass evolved stars,
transitioning from spherically symmetric to asymmetric shapes. We used
databases of circumstellar 1612 MHz OH and 22.235 GHz H2O maser sources to
search for new WF candidates using the criterion of a larger velocity range of
the H2O maser emission compared to that of the OH maser emission. Thus, it is
in principle possible to identify WFs with an H2O velocity ranges smaller than
those for the previously known WFs. For the OH maser line, we analyzed
database entries of 8,474 observations from 2,195 sources, and 6,085
observations from 3,642 sources for H2O maser line. After a close examination
of the velocity ranges and line profiles, we identified 11 sources that meet
the criterion mentioned above. We examined the IRAS colors of the selected
sources and found that two of them (IRAS 19069+0916 and IRAS 19319+2214) are
in the color region for post-AGB stars. We find that the maser velocity
criterion can discover other astrophysically interesting objects than just
WFs. Such objects may include peculiar planetary nebulae with maser emissions
and stellar merger remnants.

##  1 Introduction

Low- and intermediate-mass stars with an initial mass of up to about 8
$\rm{M}_{\odot}$ undergo active mass loss during the Asymptotic Giant Branch
(AGB) phase, in which their circumstellar envelopes (CSE) gradually develop
with losing material consisting of cold molecular gas and dust to interstellar
space (see, e.g., Olofsson, 1996; Höfner & Olofsson, 2018). In the post-AGB
phase, active mass loss ceases and the inner edge of the CSE is detached from
the stellar photosphere. One of the notable phenomena that occurs during
AGB/post-AGB phases is the morphological change of the CSE, so as to deviate
from spherical symmetry and eventually to evolve into many different shapes
observed in planetary nebulae (PNe). The physical mechanism of this
morphological change is not yet fully understood (see, e.g., Balick & Frank,
2002). Here it must be emphasized that understanding how the AGB/post-AGB CSE
changes their shapes is about more than just knowing the dynamics of the CSE:
i.e., the structure and kinematics of the CSE are closely related to the
molecular chemical reactions that occurred in the CSE. In order to accurately
model these reactions, we need to have a clear understanding of the
distributions of temperature, gas/dust density, and internal/external
radiation field within CSEs throughout their evolution.

One of the key questions is when the CSE morphological changes begin (in other
words, when the CSEs begin to deviate from spherical symmetry). Many previous
studies have mentioned that water fountains (WFs) are the key to the problem.
(see, e.g., Imai, 2009, 2013). WFs are often thought to be low- and
intermediate-mass evolved stars that harbor a small-scale collimated molecular
jet at the center of their CSEs. The conditions traditionally required to
identify WFs are as follows: (1) the velocity range of the 22.235 GHz H2O
maser emission (i.e., twice the expansion velocity of the molecular jet)
exceeds 100 km s-1, (2) the collimated structure of the molecular jet is
confirmed by radio interferometric observations, and (3) the infrared
properties are not inconsistent with the status of low- and intermediate-mass
evolved stars in terms of color index and morphology. To date, 16 WFs are
known (see Appendix E), including several well-studied objects, such as IRAS
16342$-$3814 (Sahai et al., 1999), W43A (Imai et al., 2002), IRAS 19134$+$2131
(Imai et al., 2004), IRAS 18286$-$0959 (Yung et al., 2011) and IRAS
18043$-$2116 (Uscanga et al., 2023). Since the length and speed of the
molecular jet can be obtained from very long baseline Interferometry (VLBI)
observations, the kinematic timescale of the jet can be calculated. The
calculated timescales of WF jets have remarkably small values, ranging from a
few decades to hundreds of years (see, for example, Morris et al., 2003; Imai
et al., 2004; Vlemmings et al., 2014; Imai et al., 2020). However, because
maser emission is used as the probe, the derived parameters are a subject to
uncertainty. However, at least for W43A, of which the CO thermal line (CO
$J=2-1$) has been observed using the Atacama Large Millimeter/submillimeter
Array (ALMA), it has been confirmed that the kinematic timescale of the
molecular jet is definitely short ($\sim 65$ yr; Tafoya et al., 2020). In
addition, a previous VLBI observation in the 22 GHz H2O maser line (Vlemmings
et al., 2014) also show that magnetic fields are involved in jet formation. On
the base of the number of known WFs in the Milky Way and the assumed duration
of a WF jet, Khouri et al. (2021) recently proposed that WF sources are binary
stars which are just experiencing a common-envelope phase.

An important issue when considering the relationship between the evolution of
low- and intermediate-mass evolved stars and the formation of WFs is to
accurately determine the current evolutionary state of the known WFs.
Considering the observational properties of the known WFs, the evolution of
the CSE to deviate from spherical symmetry may begin earlier than the post-AGB
phase. For example, SiO masers were detected in some WFs (see, e.g., Nakashima
& Deguchi, 2003; Amada et al., 2022). Most SiO maser sources are found around
AGB stars or red supergiants, with several known exceptions for young stellar
objects (YSOs; see, e.g., Cho et al., 2016). Even if there exists a period of
emitting the SiO maser after the transition to the post-AGB phase, it is
likely to be short, and very close to the AGB phase. Conversely, some WFs have
been proposed as more evolved objects, such as PNe (Gómez et al., 2015b). If
the number of known WF samples is large, it would be possible to discuss the
evolutionary status of WFs from a statistical analysis, in terms of color
index, etc., for instance, using public infrared photometric data.

The WFs known so far were mostly found by their unusually wide velocity range
($>$ 100 km s-1) of H2O maser emission. However, the apparent velocity of the
WF jet is expected to vary with the inclination angle and the degree of jet
development, i.e. how much time has elapsed since the acceleration began
(Engels, 2002; Yung et al., 2013). Therefore, the sample bias caused by
selecting objects by a physically unjustified jet velocity criterion (i.e.,
$>$ 100 km s-1) would be another serious scientific problem.

In view of these two problems, our team used for the first time the comparison
of the H2O and OH maser velocity range systematically to find new WF
candidates (Yung et al., 2013). In the present study, with the aim of
searching for new WFs, we cross-checked H2O (22.235 GHz) and OH (1612 MHz)
maser sources with the databases of circumstellar maser sources and also
performed a large-scale systematic comparison of the velocity ranges of H2O
and OH maser lines. We then investigated the basic properties of the sources
satisfying the velocity conditions proposed in Yung et al. (2013) using maser
spectra from the literature and infrared archival data. Finally, we discuss
how effective such a comparison of maser line velocity ranges is as a method
of selecting WF candidates, and what other types of objects might be selected
under the same selection conditions. Because a detailed explanation of the
methodology is not provided in Yung et al. (2013), it is given in Section 2 in
this paper.

##  2 Methodology and Data Processing

###  2.1 Basic ideas for finding WF candidates using circumstellar maser
source databases

Firstly we briefly summarize the basic ideas behind our WF candidate search,
and in the following subsections we describe the more specific search process
in detail. The main purpose of the present work is to select WF “candidates”,
and therefore confirmatory interferometric observations are needed to conclude
that each of them is indeed a WF. The limitations of our search and the
properties of the selected objects are discussed in Section 4.

The working hypothesis here is that a CSE is fully spherical before a WF jet
is launched. This situation is illustrated in Figure 1. A spherically
expanding CSE of an evolved star with oxygen-rich chemistry often emits H2O
maser and OH maser lines (note:“oxygen-rich” means that the number of oxygen
atoms exceeds that of carbon atoms). Here we suppose a CSE in which both maser
lines are detected. Due to the different physical conditions required to
excite the maser lines, the location of the H2O maser emission region is
closer to the central star than the OH maser emission region (generally, the
H2O maser line is emitted near the central star within a few tens of stellar
radii, and the OH maser line is emitted from the region within a few hundred
stellar radii; see e.g, Richards et al., 2011; Orosz et al., 2017). In
addition, the formation of OH from H2O molecules by photodissociation due to
interstellar UV radiation also affect the spatial distribution of H2O and OH.
In the H2O maser region, the molecular gas is accelerated outwards by the
stellar radiation pressure on the dust component (stellar pulsation could also
be involved in the acceleration mechanism; see, e.g., Höfner & Olofsson,
2018). As a result, there is no fixed pattern to the H2O maser line profile,
and it can take on many different types of profiles. On the other hand, in the
OH maser region, the molecular gas has already reached the terminal velocity.
Thus, the line profiles of the OH maser emission of the spherically expanding
envelope are known to be in the form of double peaks corresponding to the
receding and approaching sides of the CSE (see, e.g., Habing, 1996).
Importantly, in the spherically expanding envelope, the radial velocity of the
H2O maser emission never exceeds the velocity range of the OH maser emission.
$\Delta V_{\rm OH}\geq\Delta V_{\rm H_{2}O}$ always holds for a spherically
expanding envelope.

Secondly, we consider the situation after the launch of a WF jet hosting H2O
maser emission. A schematic of this situation is illustrated in Figure 2.
Here, we assume that the expansion velocity of the WF jet is faster than that
of the spherically expanding envelope. Under this assumption, if the jet axis
is not nearly perpendicular to the line-of-sight direction, then the high-
velocity component of the H2O maser emission emitted by the WF jet is detected
outside the OH velocity range (namely, $\Delta V_{\rm OH}\leq\Delta V_{\rm
H_{2}O}$). This means that for CSEs in which both the H2O 22.235 GHz and the
OH 1612 GHz maser lines are detected, if $\Delta V_{\rm OH}\leq\Delta V_{\rm
H_{2}O}$ relation is satisfied, then such a maser source can be considered a
candidate for WF.

Figure 1: Schematic view of the spherically expanding CSE emitting the OH 1612
MHz and H2O 22.235 GHz maser lines. The right panel shows the structure of the
CSE, and the left panel shows the conceptual line profiles of the OH and H2O
maser lines. The emission regions of the OH and H2O masers within the CSE are
shown in the same color as the spectra (i.e. dark blue for OH and red for H2O
masers). $\Delta$VOH represents the velocity difference between the red-
shifted and blue-shifted peaks of the OH maser line. $\Delta$V${}_{\rm
H_{2}O}$ is the velocity range of the H2O maser emission. Figure 2: Schematic
of the CSE of a water fountain emitting the OH 1612 MHz and H2O 22.235 GHz
maser lines. It differs from Figure 1 in the presence of a WF jet at the
center of the spherically expanding envelope. The two-sided green arrow
denotes the WF jet that produces double-peaked spectral profile. Note that
there is still no conclusive information on the relative size of the WF jet to
the spherically expanding envelope, and there are uncertainties in the
representation. There also may exist low-velocity components, which may be
associated with the WF jet and/or relics of a spherical CSE, within the width
represented by $\Delta$VOH in the left panel. In this figure, it is assumed
that the velocity and inclination of the jet with respect to the observer’s
line of sight combine to produce a H2O velocity range that is larger than that
of the OH emission.

###  2.2 Data used in the analysis

Figure 3 shows the overall flow of the data analysis performed in the present
study (excluding the process of more detailed inspection discussed in Section
3.2). As a source of information on OH maser sources, we used the database of
Engels & Bunzel (2015, hereafter OH database) accessible from
CDS/VizieR111https://cdsarc.cds.unistra.fr/viz-bin/cat/J/A+A/582/A68. This
database contains the records of 13,655 observations toward 2,341 sources in
the OH 1612 MHz, 1665 MHz, and 1667 MHz lines. In the present analysis we used
only the 1612 MHz line data, which consist of 8,474 observations toward 6,068
sources (2,195 out of them are detected sources). For the H2O 22.235 GHz maser
line, we used an unpublished database (hereafter H2O database) compiled by the
same authors. The version 0.1 of the database used for our analysis contains
records of 6,085 observations towards 3,642 sources (890 of them are detected
sources). A machine-readable table listing the detected H2O maser sources and
the reference information of the H2O maser observations is appended to this
paper (see, Sections 2.3 and Appendix F). The H2O maser database contains data
from 57 papers. It should be noted that these searches are based on different
selection criteria (including blind surveys), have different sensitivity
limits, cover different areas of the sky, etc., and therefore the database
must be considered heterogeneous in nature. Maps with the distributions of the
OH and H2O maser sources in the sky are given in the appendix A (see, Figures
9 and 10).

In addition to the above two databases of circumstellar maser sources, three
infrared data archives were used: i.e., the AllWISE Data
Release222https://wise2.ipac.caltech.edu/docs/release/allwise/ of the Wide-
field Infrared Survey
Explorer333https://wise2.ipac.caltech.edu/docs/release/allsky/ (WISE, Wright
et al., 2010; Mainzer et al., 2011, W1 (3.4 $\mu$m), W2 (4.6 $\mu$m), W3 (12
$\mu$m) and W4 (22 $\mu$m)), the Two Micron All Sky
Survey444https://irsa.ipac.caltech.edu/Missions/2mass.html (2MASS, Skrutskie
et al., 2006, J (1.235 $\mu$m), H (1.662 $\mu$m) and Ks (2.159 $\mu$m) ), IRAS
catalog of Point Sources, Version 2.0555https://cdsarc.cds.unistra.fr/viz-
bin/cat/II/125 (IRAS PSC, 12, 25, 60 and 100 $\mu$m).

The AllWISE data were used for two purposes: (1) to know the exact location of
the maser sources by identifying the WISE counterparts, and (2) to remove
contamination, especially YSOs, according to the infrared morphology and color
indices. The 2MASS data were originally intended to be used for the same
purpose as the AllWISE data. However, since the observing wavelengths of 2MASS
are shorter than those of WISE, there exist many cases where counterparts of
maser sources are not found, so the 2MASS data were used only as an additional
aid to remove contamination by inspecting the near-infrared morphology. The
IRAS PSC was primarily used to compare the mid-infrared properties of the
selected sources with previous studies.

Figure 3: Flowchart of the source selection in this study.

###  2.3 Processing of H2O maser data and cross-checking with infrared sources

The database of circumstellar H2O masers contains the following information:
i.e., name of the source, coordinate values used in the observation, date of
the observation, the maximum and minimum velocities at which maser emission
was detected, the velocity at the emission peak, velocity integrated intensity
of the maser emission, root-mean-square noise level, and bibliographic
information of the observation. This database includes non-detection records,
and the number of objects for which H2O maser emission has been detected is
890 out of 3,642 objects666For these 890 sources, a machine-readable table
summarizing the coordinate information used in the radio observations, the
coordinate information of the WISE counterparts, the results of the infrared
image inspection, and bibliographic information is attached to this paper. See
Appendix F for details.. These 890 H2O maser sources include all 16 known WF
sources and 4 WF candidates found by Yung et al. (2013, 2014)777Yung et al.
(2013, 2014) actually found five WF candidates, one of which (IRAS 18455+0448)
was confirmed for WF status by Vlemmings et al. (2014).

The role of the H2O maser line in this analysis is to trace the high-velocity
component possibly caused by the WF jet, as mentioned in Section 2.1. The
maximum and minimum velocities for each observation are included in the
database. Since the maser line profile varies with time, when multiple
observations were available for a single object, the maximum and minimum
velocities were determined by comparing all observations. In the following
text we will refer to the maximum and minimum velocities of the H2O maser
emission as $V_{\rm max}$ and $V_{\rm min}$ respectively.

Here, it should be noted that the database is an early version that is still
under construction, and that the database may contain a significant number of
YSOs. Therefore, initially, for a subset of the above 890 objects with H2O
maser emission, we briefly examined WISE and 2MASS infrared images to see how
serious the YSO contamination is. Specifically, we checked 2MASS and WISE
images for the presence of extended nebulosity, and indeed we found a non-
negligible number of objects with a strong possibility of being YSOs (see,
Figure 4). Then, based on the result of this initial inspection, we further
inspected the WISE and 2MASS images by eye for all 890 objects and identified
128 YSO candidates with extended nebulosity. This eye inspection was done
using the Aladin Sky Atlas888http://aladin.cds.unistra.fr/aladin.gml provided
by CDS on the web. The data of these YSO candidates were excluded from further
analysis. The results of the inspection are summarized in the machine readable
table (see, Appendix F). It is also true that there are cases where it is
difficult to distinguish between YSOs and evolved stars based solely on the
eye inspection of IR images (Ramos-Larios et al., 2012). However, such cases
where the distinction is difficult are rare, and in most cases the nebulosity
of AGB/post-AGB seen in IR images can basically be distinguished from YSOs by
the shape and extent of the nebulosity.

Then additional 230 H2O maser sources were eliminated from the WF candidates
similarly by eye inspection. The eliminated sources encounter the following
two cases: (1) no clear infrared counterpart is found within about 1′ of the
H2O maser source, and (2) the location of the maser source is close to the
Galactic Center or the Galactic plane that are too crowded with infrared
sources for identifying the counterpart. These maser sources were excluded
from the analysis for two reasons. First, their property is not consistent
with that of the known WF sources that have bright infrared counterparts.
Second, if the source is in a crowded region, it is not possible to fully
constrain the infrared counterpart due to the lack of angular resolution of
radio observations.

We have noticed that among the known WFs there exists a case where the
position of the H2O maser source and that of the infrared counterparts given
in the infrared source catalog do not match. For example, IRAS 18139$-$1816 is
such a case. The position of the WISE and 2MASS infrared point source
occasionally differs by about 30′′ from the coordinates recorded in the
database of H2O maser sources. The misalignment is caused by the fact that
early H2O maser observations were made using IRAS positions, which can
sometimes have an uncertainty larger than 30′′ (see, Section 4.1). We have
excluded such cases to ensure the reliability of the analysis.

Then we identified infrared stellar counterparts (i.e., infrared point
sources) for the remaining H2O maser sources using the WISE or 2MASS catalogs
to obtain accurate coordinate information. We found that all of those H2O
maser sources have WISE counterparts in at least one of the 4 WISE bands, but
many sources were not identified in the 2MASS J, H or K images (see Figure 5).
Therefore, we cross-checked the WISE point sources with the maser sources.

To perform this cross-matching, we used the
X-Match999http://cdsxmatch.u-strasbg.fr/ service provided by CDS. When the
cross-matching radius was set to 20′′ and the H2O maser sources were matched
against the AllWISE catalog, about 10% of the H2O maser sources had two or
more WISE sources matched. In these cases, we checked the WISE infrared colors
to determine the correct counterpart. Mass-losing evolved stars are typically
surrounded by a cold dust envelope with an effective temperature of about
150–300 K. As a result, their spectral energy distribution (SED) has an excess
at mid-infrared wavelengths. Therefore, the AllWISE colors can be used to
exclude sources other than evolved stars. We determined the search radius in
X-Match to be 20′′ after much trial and error. According to our eye inspection
there exist 11 sources that have point source counterparts in the WISE images,
but for which no counterpart was found in the AllWISE catalog. These objects
were excluded from the subsequent analyses for reliability reasons.

We calculated color indices using a combination of the W1 and W4 fluxes and
the W2 and W4 fluxes of the WISE source to see if there exist a mid-infrared
excess in the possible counterpart. Since the flux values in the AllWISE
catalog are given with respect to Vega magnitudes, a color index value of 0
does not correspond to equal flux in the two bands. Therefore, we first
converted the WISE magnitudes to fluxes in unit of Jy as described in Wright
et al. (2010) and then calculated the color indices as follows:
[W1]-[W4]=$2.5\log(F_{\scriptscriptstyle\text{W4}}/F_{\scriptscriptstyle\text{W1}})$
or
[W2]-[W4]=$2.5\log(F_{\scriptscriptstyle\text{W4}}/F_{\scriptscriptstyle\text{W2}}$).
Using these color indices, we eliminated WISE point sources with no infrared
excess (i.e., [W1]-[W4]$<0$ and [W2]-[W4]$<0$) from the candidates for
counterparts. The WISE photometric values for some sources have a low quality
flag due to saturation, but the saturation is not a serious problem here since
the WISE photometric values were used to check for excess intensity in the W4
band over the W2 band (in evolved stars, the difference in intensity between
W2 and W4 is quite obvious). There exist a very small number of cases where
the counterparts could not be constrained to a single object based on the
color index alone (i.e., there are two red sources within a beam). In such a
case, priority was given to those with high intensity in the W4 band. When
there are two or more sources of the same color and intensity, the choice
would be difficult, but fortunately such a case did not exist in the present
samples. Through the above process, we finally uniquely identified WISE point
sources that are considered mass-losing evolved stars for 496 H2O maser
sources.

Figure 4: Examples of sources excluded by eye inspection of infrared images.
The top row shows composite color images created from WISE W1, W2, and W4
bands. The bottom row shows composite color images created from 2MASS J, H,
and K bands. For both WISE and 2MASS images, the wavelength bands are assigned
to the blue, green, and red channels in order from shortest to longest
wavelength when the composite color image is generated. Both image sizes are
600′′ square, and the north is up, east is left. The two sources on the left
were excluded as obvious YSOs. The two on the right were excluded because they
could not be constrained to a single infrared counterpart to the maser source
(these sources also appear to be YSOs). Figure 5: Examples of sources that
could be clearly identified as counterparts to the maser source. The details
of the infrared images are the same as in Figure 4. The source on the left
(IRAS 06469+0333) is an example of a clear counterpart found only in the WISE
images (see text for details). The location of the H2O maser source is around
the tip of the white arrows in the lower left panel (2MASS image), but no
corresponding red point source is clearly visible in this vicinity.

###  2.4 Processing of OH maser data and velocity comparison

The 1612 MHz OH maser database of Engels & Bunzel (2015) contains information
such as source name, coordinate values, bibliographic information of the
observations, velocity information of the line, line intensity, and profile
classification for each source. In this analysis we used the velocities of the
blue-shifted and red-shifted peaks of the double peak profile of the OH maser
line associated with a spherically expanding CSE. Hereafter, these velocities
are denoted as $V_{b}$ and $V_{r}$. However, at times only a single peak
(either on the receding or on the approaching side of the spherically
expanding envelope) is detected due to time variability of the maser emission
and/or inhomogeneous gas distribution. In such cases the peak velocity is
recorded in the $V_{b}$ column by default, unless the peak can be assigned to
the blue- or red-shifted part of the OH maser profile, using other independent
observations.

Similar to the H2O maser database, there are many cases where multiple
observations are registered for a single object. In such cases, to obtain the
most accurate expansion velocity of the spherically expanding envelope, we
chose the largest value of all previous observations for $V_{r}$ and the
smallest value of all previous observations for $V_{b}$ as the representative
value. Observations for which no values were given for the $V_{r}$ and $V_{b}$
columns (i.e. non-detection) were excluded from the analysis.

We then cross-matched the OH maser sources with the AllWISE point sources
using the X-Match service, in the same way as for the H2O maser sources. As
with the H2O maser data, the search radius was set to 20′′. Then, if an
AllWISE source corresponding to the H2O maser source, which is determined in
Section 2.3, was found within 20′′ of the OH maser source, it was assigned to
be the counterpart as the OH maser source. That is, the AllWISE source acts as
an interface to pair the OH maser source with the H2O maser source.

Using the above procedures, we uniquely matched 229 objects between OH maser ,
H2O maser , and AllWISE sources. For these objects we performed velocity
comparisons as described in Section 2.1. Specifically, for the case of a
double peak profile of the OH maser source, we compared the velocity range of
the H2O emission ($V_{\rm min}<V_{\rm lsr}<V_{\rm max}$) with that of the OH
maser emission ($V_{b}<V_{\rm lsr}<V_{r}$). We have supposed that the maximum
and minimum velocities of the H2O maser emission ($V_{\rm max}$ and $V_{\rm
min}$) need not both exceed the velocity range of the OH maser emission to
allow for the possibility of time variability in the intensity of the H2O
maser emission (or to deal with uncertainties caused by inhomogeneity of the
material distribution) for identification of a WF candidate.

We also consider the case where the OH maser has a single peak. In this case,
the velocity of the single peak is denoted as $V_{s}$ ($=V_{r}=V_{b}$). If
either the maximum or minimum velocity of the H2O maser emission (i.e.,
$V_{\rm max}$ or $V_{\rm min}$) is more than 60 km s-1 away from the peak
velocity of the OH maser emission (i.e., $V_{s}$), the object is also
considered as WF candidate. This means that we set an upper limit of 30 km s-1
for the expansion velocity of the spherically expanding envelope formed during
the AGB phase. It is generally known that the expansion velocity of the CSE of
AGB stars is usually between 10–20 km s-1, but we set 30 km s-1 here to ensure
a more reliable selection of WF candidates.

After comparing the velocities, 52 objects met the criteria. Of these 52,
seven were already known WFs and four were previously reported WF candidates
(see, Section 4.3 for known WFs in the sample). This means that 41 of the 52
objects are “potentially” new WF candidates. However, as mentioned in Section
2.1, we note again that we cannot immediately say that these 41 are WF
candidates as they are. The properties of these 41 objects will be carefully
discussed in more detail in following sections.

##  3 Results

###  3.1 Overview of the selected sources

Finally, we found 41 OH/H2O maser sources that satisfy the velocity conditions
described in Section 2.1. This means that 17.9% of the 229 sources meets the
velocity condition. The maser velocities, coordinate information (WISE
position), and infrared (IRAS) fluxes for the infrared counterparts of the 41
selected sources are summarized in Tables 1 and 2. The results of the velocity
comparison between the OH and H2O maser lines are shown in Figure 6. In the
past, H2O maser sources with velocities greater than 100 km s-1 were
identified as WFs. However, this search did not find any such “classical” WFs
that were slipped from previous research. (Some of the sources classified as
YSO candidates have H2O maser velocity ranges exceeding 100 km s-1. These
sources are summarized in Table 4 in Appendix D .) Half the velocity between
the two peaks of the OH maser emitted by the spherically expanding envelope
corresponds to the expansion velocity of the envelope. This expansion velocity
is typically about 10–20 km s-1, which corresponds to 20–40 km s-1 in the
velocity range shown in Figure 6. For some sources the velocity range of the
OH maser emission is clearly less than 10 km s-1. This can be explained by two
possible scenarios: (1) the expansion velocity of the spherically expanding
envelope is slow, or (2) the expansion velocity of the spherically expanding
envelope is within the normal range, but only one peak is detected, either on
the receding or on the approaching side of a spherically expanding envelope.
By examining the OH maser line profiles of individual sources, it is possible
to make some speculations as to which of the two situations applies. A
detailed discussion of maser profiles for individual sources is given in
Section 3.2 and Appendix B and C.

The histograms of the velocity ranges of the OH and H2O maser emission are
shown in Figure 7. Looking at the lower velocity range below 20 km s-1 (as
indicated by the blue vertical dotted line in Figure 7), we can see that there
are more samples in this range for H2O maser sources than for OH maser
sources. As explained in Section 2.1, this is due to the fact that the
material in the H2O maser emission region is still accelerating, resulting in
a slower overall expansion velocity compared to the material in the OH maser
emission region. The maximum expansion velocity of the OH maser among all
samples is 29.8 km s-1 (the object exhibiting this velocity is IRAS 19083+0851
showing the velocity range of 59.5 km s-1), which is a relatively large value
but still possibly within the range of the typical expansion velocity of an
AGB envelope (there are previous studies suggesting that IRAS 19083+0851 is
possibly an RSG; See, Section 3.2.7 for details). Notably, IRAS 19069+0916,
has a velocity range of H2O maser emission that exceeds this maximum velocity
range of OH maser emission. The properties of this object are discussed in
detail in Section 3.2.6.

In Figure 8 we have plotted the selected 41 sources on an IRAS two-color
diagram to briefly inspect their evolutionary state with respect to mid-
infrared colors. The sources are divided into two groups, “Group A” (red dots)
and “Group B” (blue dots), based on the maser profile checks described in
Section 3.2 (See also this section for definitions of Group A and B). The
areas surrounded by the black dashed lines in the figure indicate the well-
known classification of evolved stars by van der Veen & Habing (1988).
According to this classification, regions II and IIIa generally correspond to
the AGB phase, while regions IIIb, IV, V, and VIII represent the post-AGB and
later phases [it should be noted, however, that RSGs often have infrared
colors similar to those of AGB/post-AGB stars and cannot be distinguished from
AGB/post-AGB stars by their infrared colors alone; See, Jiménez-Esteban et al.
(2021)]. The 41 selected sources are distributed over a very wide range in
Figure 8, from the AGB phase to the post-AGB phase and beyond. It is worth
noting that quite a large fraction of the 41 selected sources are distributed
in regions II and IIIa, corresponding to the AGB phase. In the figure, the
black dots and triangles represent known WFs and WF candidates, respectively.
Most of the black data points are outside regions II and IIIa, but two
sources, IRAS 15193+3132 and IRAS 18056–1514, are within region IIIa. These
two sources were previously identified as WF candidates using methods similar
to those used in this study (Yung et al., 2013). The gray crosses in the
figure represent the sources for which the velocity range of the H2O maser
line did not exceed that of the OH maser line. Of the sources we considered in
Group A, IRAS 19319+2214 is in the IIIb region, where early post-AGB stages
predominate; IRAS 19069+0916 is located near the boundary between the IIIb and
IV regions, which is another source that shows early post-AGB stellar color.
Three other Group A sources (IRAS 18251–1048, IRAS 18588+0428, IRAS
19067+0811) are found in the color regions (IV, VIb, VIII), containing more
evolved AGB stars and early post-AGB stars. These properties will be discussed
in the next section.

Figure 6: Comparison of the velocity ranges of the H2O 22.235 GHz and OH 1612
MHz maser lines for 41 selected sources. The blue vertical lines represent the
velocities of the red-shifted and blue-shifted peaks of the OH maser line. The
red line represents the velocity range of the H2O maser emission. The sources
are arranged in ascending order with respect to right ascension. Figure 7:
Histograms of the velocity ranges of the H2O 22.235 GHz (upper panel) and OH
1612 MHz (lower panel) maser lines. Half of the velocity range corresponds to
the expansion velocity. The blue vertical dashed line represents a velocity
range of 20 km s-1 (corresponding to an expansion velocity of 10 km s-1; see
text). Figure 8: IRAS two-color diagram of 41 sources selected by the velocity
condition. This plot uses color indices defined as follows
$[12]-[25]=2.5\log\left(\frac{F_{25\mu\mathrm{m}}}{F_{12\mu\mathrm{m}}}\right)$
and
$[25]-[60]=2.5\log\left(\frac{F_{60\mu\mathrm{m}}}{F_{25\mu\mathrm{m}}}\right)$.
The red and blue dots represent the 41 objects selected by our velocity
comparison analysis. “Group A” represents sources for which the deviation of
the H2O maser velocity range from the OH maser velocity range is reliable;
“Group B” represents sources for which the velocity deviation is not
sufficiently reliable (see, Section 3.2 for details). The gray crosses
represent the sources that did not meet the velocity condition. The black
circles are known WFs, and the black triangles are WF candidates reported by
Yung et al. (2013, 2014). The black ”X” marks represent objects that were
included in the 41 selected sources but were clearly not AGB or post-AGB
stars. The data points with gray arrows are sources with an IRAS flux quality
flag (QF) of 1. The IRAS flux with a QF of 1 is practically an upper limit,
and therefore the true value of the IRAS color is assumed to be in the
direction of the arrow. The QF values for all sources are summarized in Table
3. The black line corresponds to the blackbody radiation curve, and the area
separated by the dotted line is the classification of evolved stars given by
van der Veen & Habing (1988). Table 1: Velocity information for 41 selected
sources

Source name |  $V_{b}$ (OH) |  $V_{r}$ (OH) | $\Delta V_{\rm OH}$ |  $V_{\rm min}$ (H2O) |  $V_{\rm max}$ (H2O) | $\Delta V_{\rm H_{2}O}$ | Reference for H2O
---|---|---|---|---|---|---|---
| (km s-1) | (km s-1) | (km s-1) | (km s-1) | (km s-1) | (km s-1) |
IRAS 04396+0647 | 15.5 | 20.5 | 5.0 | 12.1 | 14.4 | 2.3 | VAL01
IRAS 05027$-$2158 |  $-$15.9 |  $-$0.2 | 15.7 |  $-$30.5 |  $-$25.9 | 4.6 | TAK01
IRAS 05528+2010 |  $-$46.0 |  $-$39.0 | 7.0 |  $-$42.0 |  $-$36.0 | 6.0 | NEU21
IRAS 07445$-$2613 | 76.9 | 87.4 | 10.5 | 76.6 | 85.3 | 8.7 | VAL01; KIM13
IRAS 07446$-$3210 | 39.0 | 47.4 | 8.4 | 19.5 | 31.9 | 12.4 | TAK01; KIM10
IRAS 08357$-$1013 |  $-$3.9 | 39.4 | 43.3 | 65.0 | 66.6 | 1.6 | TAK01; TAK94
IRAS 13001+0527 | 23.5 | 24.0 | 0.5 | 8.4 | 25.0 | 16.6 | VAL01
IRAS 15060+0947 |  $-$16.3 | 1.8 | 18.1 |  $-$14.4 | 3.0 | 17.4 | VAL01; TAK01
IRAS 15255+1944 |  $-$1.4 | 14.0 | 15.4 |  $-$1.8 | 8.8 | 10.6 | VAL01; KIM10
IRAS 15408$-$5413 |  $-$105.2 |  $-$101.1 | 4.1 |  $-$85.7 |  $-$75.9 | 9.8 | GOM15A; WAL11
IRAS 17343+1052 |  $-$62.6 |  $-$53.1 | 9.5 |  $-$63.0 |  $-$53.0 | 10.0 | YUN13
IRAS 17392$-$3020 |  $-$24.9 | 18.5 | 43.4 |  $-$18.0 | 21.0 | 39.0 | DEA07
IRAS 17443$-$2949 |  $-$16.7 |  $-$4.8 | 11.9 |  $-$6.4 | 0.9 | 7.3 | GOM08; SUA07
IRAS 17540$-$1919 | 56.4 | 64.0 | 7.6 | 55.0 | 55.0 | 0.0 | CHO17
IRAS 17580$-$3111 | 1.0 | 29.2 | 28.2 | 0.4 | 21.3 | 20.9 | SUA07; GOM08
IRAS 18009$-$2019 | 3.0 | 26.0 | 23.0 |  $-$1.5 | 21.0 | 22.5 | VAL01
IRAS 18034$-$2441 |  $-$12.7 | 23.7 | 36.4 |  $-$13.2 |  $-$13.2 | 0.0 | CHO17
IRAS 18050$-$0518 |  $-$34.2 |  $-$20.3 | 13.9 |  $-$32.0 |  $-$20.0 | 12.0 | YUN13
IRAS 18052$-$2016 | 23.0 | 80.2 | 57.2 | 14.0 | 43.0 | 29.0 | DEA07
IRAS 18085+0752 |  $-$71.2 |  $-$63.5 | 7.7 |  $-$70.6 |  $-$67.9 | 2.7 | VAL01; YUN13; CHO17
IRAS 18251$-$1048 | 71.0 | 111.2 | 40.2 | 70.0 | 110.0 | 40.0 | YUN13
IRAS 18348$-$0526 | 13.2 | 41.7 | 28.5 | 25.8 | 44.7 | 18.9 | TAK01; TAK94
IRAS 18436+4334 |  $-$28.2 |  $-$13.8 | 14.4 |  $-$29.4 |  $-$20.2 | 9.2 | TAK01; KIM10
IRAS 18494$-$0130 | 65.8 | 91.7 | 25.9 | 65.0 | 72.0 | 7.0 | YUN13
IRAS 18588+0428 | 29.4 | 74.7 | 45.3 | 29.0 | 74.0 | 45.0 | DEA07
IRAS 19052+0922 | 26.0 | 62.3 | 36.3 | 5.6 | 12.4 | 6.8 | VAL01
IRAS 19067+0811 | 42.3 | 75.9 | 33.6 | 47.2 | 80.0 | 32.8 | VLE14
IRAS 19068+1127 |  $-$8.2 | 23.6 | 31.8 |  $-$13.0 | 22.0 | 35.0 | YUN13
IRAS 19069+0916 | 10.0 | 54.5 | 44.5 |  $-$32.2 | 47.0 | 79.2 | TAK01; VLE14
IRAS 19083+0851 | 16.3 | 75.8 | 59.5 | 7.3 | 60.7 | 53.4 | TAK01; KIM13
IRAS 19229+1708 | 14.6 | 18.7 | 4.1 | 24.3 | 53.0 | 28.7 | VAL01; YUN13
IRAS 19255+2123 |  $-$3.0 | 9.0 | 12.0 | 22.4 | 22.6 | 0.2 | GRE04
IRAS 19319+2214 | 12.3 | 29.1 | 16.8 | 8.8 | 33.6 | 24.8 | GOM15A
IRAS 19395+1949 | 25.0 | 28.2 | 3.2 | 28.0 | 56.0 | 28.0 | YUN13
IRAS 19422+3506 |  $-$65.5 |  $-$31.2 | 34.3 |  $-$68.5 |  $-$33.5 | 35.0 | VLE14
IRAS 19466+2751 | 29.6 | 57.3 | 27.7 |  $-$2.8 | 5.0 | 7.8 | VAL01
IRAS 20015+3019 | 5.5 | 7.3 | 1.8 | 16.7 | 20.1 | 3.4 | TAK01
IRAS 20440+0412 |  $-$54.8 |  $-$49.9 | 4.9 |  $-$52.0 |  $-$48.0 | 4.0 | YUN13
IRAS 20547+0247 |  $-$80.4 |  $-$75.3 | 5.1 |  $-$75.0 |  $-$75.0 | 0.0 | ZUC87
IRAS 22516+0838 |  $-$8.6 | 12.3 | 20.9 |  $-$7.8 | 14.0 | 21.8 | KIM10
IRAS 22556+5833 |  $-$67.0 |  $-$59.0 | 8.0 |  $-$58.3 |  $-$44.8 | 13.5 | VAL01; TAK01

Note. — For the OH maser observations, values from Engels & Bunzel (2015) are
used. See Appendix G for reference code.

Table 2: Infrared data for 41 selected sources

Source name | WISE position | Angdist | IRAS flux | QF
---|---|---|---|---
| R.A. | Dec. |  | $F_{\rm 12\mu m}$ | $F_{\rm 25\mu m}$ | $F_{\rm 60\mu m}$ |
| (J2000.0) | (arcsec) | (Jy) | (Jy) | (Jy) |
IRAS 04396+0647 | 70.5888064 | 6.8778639 | 1.32 | 49.3 | 30.2 | 5.7 | 333
IRAS 05027$-$2158 | 76.2118956 |  $-$21.9046758 | 0.36 | 157.0 | 82.1 | 11.8 | 333
IRAS 05528+2010 | 88.9547044 | 20.1751498 | 1.09 | 682.0 | 259.0 | 39.1 | 333
IRAS 07445$-$2613 | 116.6595203 |  $-$26.3432689 | 2.71 | 77.9 | 37.1 | 5.0 | 333
IRAS 07446$-$3210 | 116.6422915 |  $-$32.3045251 | 0.01 | 117.0 | 93.7 | 21.7 | 333
IRAS 08357$-$1013 | 129.5355866 |  $-$10.4042302 | 4.22 | 62.5 | 59.2 | 10.7 | 333
IRAS 13001+0527 | 195.6584898 | 5.1856231 | 9.27 | 462.0 | 226.0 | 39.3 | 333
IRAS 15060+0947 | 227.1072940 | 9.6051642 | 2.73 | 35.2 | 26.9 | 4.0 | 333
IRAS 15255+1944 | 231.9459816 | 19.5643970 | 7.85 | 235.0 | 150.0 | 18.4 | 333
IRAS 15408$-$5413 | 236.1658919 |  $-$54.3847874 | 0.25 | 167.0 | 351.0 | 111.0 | 333
IRAS 17343+1052 | 264.1852174 | 10.8518791 | 6.86 | 37.6 | 16.4 | 2.9 | 333
IRAS 17392$-$3020 | 265.6273152 |  $-$30.3686911 | 0.87 | 6.3 | 17.3 | 36.4 | 332
IRAS 17443$-$2949 | 266.8971436 |  $-$29.8481106 | 3.8 | 15.8 | 39.4 | 34.5 | 333
IRAS 17540$-$1919 | 269.2627897 |  $-$19.3377071 | 0.25 | 83.9 | 30.3 | 4.2 | 333
IRAS 17580$-$3111 | 270.3349569 |  $-$31.1889743 | 10.28 | 3.2 | 15.3 | 8.0 | 333
IRAS 18009$-$2019 | 270.9855195 |  $-$20.3167520 | 3.17 | 419.0 | 294.0 | 46.5 | 333
IRAS 18034$-$2441 | 271.6230128 |  $-$24.6799678 | 0.16 | 17.9 | 23.5 | 7.3 | 333
IRAS 18050$-$0518 | 271.9209666 |  $-$5.3053559 | 0.32 | 35.7 | 30.8 | 4.6 | 333
IRAS 18052$-$2016 | 272.0678773 |  $-$20.2697236 | 1.31 | 21.4 | 32.5 | 82.4 | 333
IRAS 18085+0752 | 272.7440259 | 7.8856846 | 6.41 | 21.7 | 13.1 | 1.9 | 333
IRAS 18251$-$1048 | 276.9846802 |  $-$10.7816116 | 4.21 | 6.7 | 13.3 | 16.2 | 331
IRAS 18348$-$0526 | 279.3854407 |  $-$5.3997526 | 0.09 | 360.0 | 634.0 | 463.0 | 333
IRAS 18436+4334 | 281.2922769 | 43.6354216 | 4.44 | 64.6 | 50.6 | 9.7 | 333
IRAS 18494$-$0130 | 283.0060527 |  $-$1.4461850 | 0.13 | 12.6 | 14.8 | 13.3 | 331
IRAS 18588+0428 | 285.3338523 | 4.5420510 | 1.07 | 10.8 | 13.1 | 20.1 | 332
IRAS 19052+0922 | 286.9128078 | 9.4543756 | 3.06 | 24.8 | 18.7 | 5.3 | 333
IRAS 19067+0811 | 287.2846276 | 8.2761164 | 0.22 | 24.6 | 72.1 | 38.9 | 333
IRAS 19068+1127 | 287.2986926 | 11.5466987 | 12.02 | 18.1 | 19.2 | 3.4 | 333
IRAS 19069+0916 | 287.3302736 | 9.3532127 | 2.4 | 4.8 | 10.0 | 7.2 | 331
IRAS 19083+0851 | 287.6972070 | 8.9396218 | 1.22 | 41.6 | 36.4 | 6.3 | 332
IRAS 19229+1708 | 291.3021044 | 17.2469940 | 5.71 | 46.1 | 35.6 | 7.6 | 333
IRAS 19255+2123 | 291.9334338 | 21.5009583 | 0.05 | 2.4 | 29.3 | 48.2 | 333
IRAS 19319+2214 | 293.5149397 | 22.3544324 | 3.28 | 3.2 | 6.6 | 3.2 | 333
IRAS 19356+0754 | 294.5049011 | 8.0257915 | 0.21 | 1.1 | 8.0 | 7.6 | 333
IRAS 19422+3506 | 296.0292220 | 35.2355954 | 0.14 | 203.0 | 172.0 | 28.4 | 333
IRAS 19466+2751 | 297.1552342 | 27.9956429 | 15.74 | 27.8 | 27.1 | 6.8 | 333
IRAS 20015+3019 | 300.9103669 | 30.4692632 | 0.05 | 162.0 | 115.0 | 25.5 | 333
IRAS 20440+0412 | 311.6382059 | 4.3934590 | 1.61 | 9.1 | 5.2 | 0.9 | 333
IRAS 20547+0247 | 314.3179038 | 2.9791841 | 1.14 | 45.5 | 33.8 | 10.0 | 333
IRAS 22516+0838 | 343.5474973 | 8.9019877 | 0.37 | 113.0 | 63.9 | 11.6 | 333
IRAS 22556+5833 | 344.4207333 | 58.8201558 | 7.17 | 189.0 | 170.0 | 24.8 | 333

Note. — QF: The quality flag for the IRAS flux density, where 3 means high
quality, 2 moderate quality, and 1 an upper limit. The numbers correspond to
12, 25, and 60 micron flux densities from left to right. Angdist: Angular
separation between the location of the maser sources in the database and that
of the AllWISE counterpart.

###  3.2 Characteristics of individual sources

At the beginning of Section 3.1, we mentioned that 17.9% of the sources for
which velocity comparisons were made met the velocity requirements as WF
candidates. This rate value seems too high for a percentage of WFs, and the
reliability of the maser data used and the contamination in the sample must be
carefully considered. For this purpose, here, we discuss the maser
characteristics of individual objects using available information from the
literature. Specifically, we consider the following perspectives:

  1. 1.

Whether the OH maser line profiles indeed show double-peaked profile typically
seen towards AGB stars . A considerable number of sources classified as such
profile sources have a small velocity gap between the two peaks (a few km s-1
or less). The line profiles are informative to consider whether such short-
spaced double peaks is really associated with, respectively, the receding and
approaching sides of a spherically expanding CSE. We also checked for any
velocity components that were slipped from the database.

  2. 2.

Whether there are any velocity components of the H2O maser line that we have
missed. In particular, the weak velocity component at velocities beyond the OH
maser velocity range was carefully inspected. We also checked whether the H2O
maser line profiles were similar to known WFs. The known WFs share the common
features of the H2O maser spectra as follows: (1) the velocity range is larger
than 100 km s-1, and (2) two or more peaks are always or often found. The
intensity peaks appear randomly at different radial velocities and there is no
fixed pattern: Sometimes they exhibit double peaks or multiple peaks. However,
the H2O maser spectra used for confirmation were obtained by a variety of
different instruments, and the S/N ratios achieved vary from observation to
observation. It should be added, therefore, that not all spectra meet the S/N
ratios necessary to distinguish the above conditions. For the reader’s
reference, the papers giving the spectra of H2O and OH maser lines of known
WFs are summarized in Appendix E.

  3. 3.

Comparison of SiO and OH maser velocities. Since SiO maser lines are expected
to be close to the systemic velocity of the source, the comparison with the OH
maser provides a guideline for considering whether the OH maser peak is
associated with the approaching or receding side. In addition, a large SiO
maser linewidth (e.g. $>$20 km s-1) is a hint for a possible RSG.

  4. 4.

Comparison of line profiles between the OH 1612 and 1665/1667 MHz lines. Since
the 1665 and 1667 MHz lines have pumping route different from that of the 1612
MHz line, the difference of the OH line profiles provides a clue to the nature
of CSE. We also checked the intensity ratio between the 1612 MHz and the
1665/1667 MHz lines for sources where all three lines are detected. Typically,
the 1612 MHz line is weaker than the 1665 and 1667 MHz lines in YSOs (see,
e.g., Habing, 1996). In fact, there exist no sources showing the weaker 1612
MHz intensity than the 1665/1667 MHz lines.

  5. 5.

We also checked the literature information available from SIMBAD to see if it
had been identified in objects other than AGB and post-AGB stars.

The details of the above confirmations are summarized for each source in the
subsections below and in Appendix B and C. The literature search with SIMBAD
database revealed that two of the 41 selected objects (IRAS 17443$-$2949 and
IRAS 19255+2123) are identified as planetary nebulae. The remaining 39 sources
are classified into “Group A” and “Group B” based on the results of the above
confirmation. Group A includes sources that meet either of the following
conditions: (1) it is fairly certain that the velocity range of the H2O maser
exceeds that of the OH maser, and the H2O maser line profile resembles known
WFs, or (2) even if condition (1) is not met, judging from the maser
characteristics found in the literature, it is unlikely to be an AGB or post-
AGB star with a normal spherically expanding circumstellar envelope.

Group B, on the other hand, includes objects where velocity range excess of
the H2O maser is uncertain and the difference in the maser line profile from
those of normal AGB or post-AGB stars is unclear.

To avoid making the main text too long, the following subsections describe
only the sources in Group A. Descriptions for the sources in Group B and those
identified as non-AGB/post-AGB stars are given in the Appendix B and C. For
Group A sources, we note in the description of each object whether each object
meets condition (1) or (2) above.

Table 3: Variable star names and source classification

Source name | Group | Region | Variable star name | JE21 | AAVSO | Note
---|---|---|---|---|---|---
IRAS 04396+0647 | B | IIIa | BZ Tau | 1 | — |
IRAS 05027$-$2158 | B | IIIa | T Lep | — | mira-like |
IRAS 05528+2010 | B | II | U Ori | — | mira-like |
IRAS 07445$-$2613 | B | II | SS Pup | — | mira-like |
IRAS 07446$-$3210 | B | IIIa | V353 Pup | — | — |
IRAS 08357$-$1013 | B | IIIa | NSV 17942 | — | — |
IRAS 13001+0527 | B | IIIa | RT Vir | — | mira-like |
IRAS 15060+0947 | B | IIIa | FV Boo | 1 | mira-like |
IRAS 15255+1944 | B | IIIa | WX Ser | — | mira-like |
IRAS 15408$-$5413 | B | IIIb | — | — | — |
IRAS 17343+1052 | B | II | V790 Oph | — | — |
IRAS 17392$-$3020 | B | VIII | — | — | — |
IRAS 17443$-$2949 | — | IV | — | — | — | Known PN
IRAS 17540$-$1919 | B | II | VV Sgr | — | — |
IRAS 17580$-$3111 | B | V | — | — | — |
IRAS 18009$-$2019 | B | IIIa | V4120 Sgr | — | mira-like |
IRAS 18034$-$2441 | B | IIIb | — | — | — |
IRAS 18050$-$0518 | B | IIIa | — | — | — |
IRAS 18052$-$2016 | B | VIII | — | — | — |
IRAS 18085+0752 | B | IIIa | XY Oph | 1 | — |
IRAS 18251$-$1048 | A | VIII | — | — | — |
IRAS 18348$-$0526 | B | IV | — | — | — |
IRAS 18436+4334 | B | IIIa | RW Lyr | — | mira-like |
IRAS 18494$-$0130 | B | VIb | — | — | — |
IRAS 18588+0428 | A | VIb | — | — | — |
IRAS 19052+0922 | A | IIIa | — | 1 | — |
IRAS 19067+0811 | A | IV | — | 1 | — |
IRAS 19068+1127 | A | IIIa | — | 1 | — |
IRAS 19069+0916 | A | IV | — | 1 | — |
IRAS 19083+0851 | A | IIIa | — | 4 | — |
IRAS 19229+1708 | A | IIIa | — | 4 | — | Gaia DR3 suggests LPV
IRAS 19255+2123 | — | V | — | — | — | Known PN
IRAS 19319+2214 | A | IIIb | — | 2 | — |
IRAS 19395+1949 | B | IIIa | — | 1 | — |
IRAS 19422+3506 | A | IIIa | — | 1 | — |
IRAS 19466+2751 | B | IIIa | — | — | — |
IRAS 20015+3019 | B | IIIa | V719 Cyg | 4 | — |
IRAS 20440+0412 | B | IIIa | BR Del | 1 | mira-like |
IRAS 20547+0247 | B | IIIa | U Equ | 2 | — |
IRAS 22516+0838 | A | IIIa | KZ Peg | 1 | — |
IRAS 22556+5833 | B | IIIa | — | — | — |

Note. — TE21: “Group” in Table 1 of Jimenez-Esteban et al. (2021). 1 = LPLAV,
2 = Post-AGBs, 3 = Unclassified, 4 = YSO; RSG; False OH detection; Irregular;
Optical low-amplitude regular variability.

AAVSO: Results of checking the AAVSO light curves. The souces that are
considered to show Mira-type pulsation are marked as ”mira-like”.

####  3.2.1 IRAS 18251–1048

The profile of the OH 1612 MHz line is typically double-peaked, with some fine
structure observed in the red-shifted component, as reported in previous
studies by te Lintel Hekkert et al. (1989) and Yung et al. (2014). The OH 1667
MHz line is weakly detected, with peaks located near 115 km s-1 and 20 km s-1,
as reported in Yung et al. (2014). The H2O maser line profile is double-
peaked, according to observations by Yung et al. (2013) and Cho et al. (2017),
and the emission is mostly within the velocity range of the OH maser line. The
information in the database suggests that the deviation of the H2O maser line
from the OH maser velocity range is 1.0 km s-1. The SiO $v=1$ and 2, $J=1-0$
lines have been detected near 90 km s-1, and the line profiles are roughly
single-peaked, as reported in previous studies by Nakashima & Deguchi (2003)
and Cho et al. (2017). Based on the comparison of the OH and H2O velocity
ranges, it is difficult to conclude that the H2O maser velocity deviations are
real. However, it is worth noting that the OH 1667 MHz line deviates from the
velocity range of the OH 1612 MHz line (see, Section 4.2). This source
satisfies condition (2).

####  3.2.2 IRAS 18588+0428

The profile of the OH maser line shows a typical double-peak profile (Lewis,
1994; Deacon et al., 2004). According to the spectra in the literature, the OH
maser emission appears to be detected even outside the redshifted peak, near
20 km s-1. The H2O maser spectrum shows many peaks, which are uniformly
distributed within the velocity range of the OH maser line (Deacon et al.,
2007). The SiO $v=1$ and 2, $J=1-0$ lines are detected near 55 km s-1, and the
line profiles are a typical single peak (Deguchi et al., 2004b; Yoon et al.,
2014). The deviation of the H2O maser line from the velocity range of the OH
maser is small (0.4 km s-1). However, it is worth noting that the H2O maser
profile with many peaks is similar to the known WF, IRAS 18286$-$0959 (Yung et
al., 2011), except for the H2O velocity range not largely exceeding that of
the OH maser line. This source satisfies condition (2).

####  3.2.3 IRAS 19052+0922

The OH maser shows a typical double-peak profile, as reported in the
literature such as Lewis (1994). For this source there are two H2O maser
observations in the literature. One by Valdettaro et al. (2001) and another by
Engels & Lewis (1996). Engels & Lewis (1996) observed a double peaked profile
at 37.8 and 54.8 km s-1, which fits well with the OH maser velocity. The H2O
maser profile in Valdettaro et al. (2001) has a single peak at 10 km s-1. The
10 km s-1 peak of Valdettaro et al. (2001) was apparently not present in the
spectra of Engels & Lewis (1996). The 10 km s-1 peak is clearly outside the
velocity range of the OH maser line (the peak velocity of the OH red-shifted
component is near 26 km s-1). The SiO $v=1$, $J=1-0$ line is detected, as
reported in Cho et al. (2017), with a typical single peak profile and peak
velocity near 40 km s-1. Although the H2O maser profile is single-peaked, it
should be considered a WF candidate according to our criteria. This source
satisfies condition (1).

####  3.2.4 IRAS 19067+0811

The OH maser profile shows a typical double-peak profile, as reported in
references (te Lintel Hekkert et al., 1989; Lewis, 1994; Wolak et al., 2012).
The H2O maser profile is also double-peaked, with some of the fine components
on the red-shifted side detected outside the velocity range of the OH maser
line, as reported in Vlemmings et al. (2014). The SiO $v=1$, 2 and 3, $J=1-0$
lines, SiO $v=1$, $J=2-1$ lines, and 29SiO $v=0$, $J=1-0$ and $2-1$ lines have
been detected, as reported in Nakashima & Deguchi (2007). All the SiO line
profiles are single-peaked with a peak velocity near 64 km s-1. The deviation
of the velocity of the H2O maser line from the OH maser velocity range is 4.1
km s-1 according to the database, but there may be even faster velocity
components when looking at the spectra in the literature. The H2O maser line
exhibits a double-peaked profile reminiscent of a bipolar flow. Vlemmings et
al. (2014) made mapping observations in H2O maser line with JVLA, which did
not show a narrowly collimated jet. However, the overall velocity field can be
interpreted as a loosely squeezed bipolar flow. This source satisfies
condition (2).

####  3.2.5 IRAS 19068+1127

The OH maser line shows a typical double-peak profile, as reported in the
literature such as Lewis (1994). The H2O maser profile has multiple peaks and
is distributed throughout the velocity range of the OH maser line, as reported
in the literature such as Yung et al. (2013). This source is similar to IRAS
18286–0959 (Yung et al., 2011), except that the velocity range of the H2O
maser line does not largely exceed the OH maser’s velocity range. The minimum
velocity of the H2O maser line deviates from the OH maser’s velocity range by
4.8 km s-1, according to the database. The H2O spectrum from Yung et al.
(2013) does not clearly show the velocity excess suggested by the database.
This source satisfies condition (2).

####  3.2.6 IRAS 19069+0916

The OH maser profile shows a typical double-peak profile (Eder et al., 1988;
Lewis, 1994). Three H2O maser spectra were found in the literature (Takaba et
al., 2001; Vlemmings et al., 2014; Kim et al., 2010), but the line profiles
show very large temporal variability. Looking at all the spectra in the
literature, emission lines appear over velocities from near $-$35 km s-1 to
near 45 km s-1. The velocity range of this H2O maser completely deviates from
that of the OH maser line. The $-$35 km s-1 component was detected only by
Takaba et al. (2001). Mapping observations in the H2O maser line have been
made by Vlemmings et al. (2014) using JVLA, but the $-$35 km s-1 component was
not detected. The SiO $v=1$ and 2, $J=1-0$ lines are detected near 30 km s-1,
and the line profiles are typical single peaks (Nakashima & Deguchi, 2003; Kim
et al., 2010). This source satisfies condition (1).

####  3.2.7 IRAS 19083+0851

The OH maser spectrum in Chengalur et al. (1993) show three peaks at around 17
km s-1, 37 km s-1, and 75 km s-1. Two H2O maser spectra were found in the
literature (Engels & Lewis, 1996; Kim et al., 2013), showing strong time
variation in its profile. The H2O maser profile is characterized by multiple
peaks over a wide velocity range, similar to the profile of the known WF, IRAS
18286$-$0959 (Yung et al., 2011). The minimum velocity of 7.3 km s-1 given in
the database was not confirmed by the H2O spectra in the literature. But, in
the spectrum given by Engels & Lewis (1996), a weak H2O maser emission is
detected around 27 km s-1, implying the existence of high-velocity components
on the blue-shifted side. The SiO $v=1$ and 2, $J=1-0$ lines have been
detected ( near 56 km s-1, Kim et al., 2013). The $v=1$ line shows five peaks,
and the $v=2$ line shows double peaks, with line widths as wide as 15–20 km
s-1 at zero-intensity levels. This differs from the single-peak profile seen
in a spherically expanding envelope of an AGB star. In any case, the
information in the database supports the deviation of the H2O maser velocity
from the range of the OH maser velocity. This source is suggested to be an RSG
because of the relatively large velocity range of the OH maser (Engels &
Lewis, 1996) and the small amplitude of the light curves in the near-infrared
bands (Jiménez-Esteban et al., 2021). However, similar properties may also be
found in early post-AGB stars. This source satisfies condition (1).

####  3.2.8 IRAS 19229+1708

The OH maser line has a typical double-peaked profile, with peak velocities at
14.6 km s-1 and 18.7 km s-1. However, a spectrum presented in Lewis et al.
(1990) shows a sharp peak near 15 km s-1, with a decreasing intensity on the
blue-shifted side and a slower decrease on the red-shifted side. This profile
is reminiscent of the blue-shifted component of the double-peak profile of the
OH maser line emitted from a typical spherically expanding envelope of an AGB
star. There are two H2O maser spectra in the literature (Valdettaro et al.,
2001; Yung et al., 2013) showing multiple peaks within a velocity range of
about 50 km s-1, with strong time variability. The most red-shifted peak is
detected near 51 km s-1. The 51 km s-1 component of the H2O maser line is 36.4
km s-1 away from the 14.6 km s-1 peak of the OH maser line. Assuming the
approaching side of the spherically expanding envelope is detected in the OH
maser line, and the 51 km s-1 component of the H2O maser line corresponds to
the receding side, the expansion velocity of the envelope is calculated to be
18.2 km s-1. This source satisfies condition (2).

####  3.2.9 IRAS 19319+2214

The OH maser profile is characterized by typical double peaks, but there also
exist wing components outside of the two peaks, as seen in the spectrum from
Lewis et al. (1990). The database indicates that the velocity range of the H2O
maser deviates from that of the OH maser on both the red- and blue-shifted
sides. However, the spectra in the literature only show deviation on the red-
shifted side. The H2O maser profile also shows multiple peaks, similar to the
known WF IRAS 18286$-$0959 (Yung et al., 2011). This source satisfies
condition (1).

####  3.2.10 IRAS 19422+3506

The OH maser profile is a double peak, which is typically emitted from a
spherically expanding envelope of an AGB star, as reported in Le Squeren et
al. (1992). The H2O maser line, however, is detected from a wide velocity
range that extends beyond the velocity range of the OH maser line, and its
profile varies significantly depending on the time of observation, as reported
in Vlemmings et al. (2014). The H2O maser line of this source has been mapped
using the JVLA (Vlemmings et al., 2014), and the resulting map clearly shows
structures that are not spherically symmetric, but it is not a collimated
bipolar flow. This source satisfies condition (1).

####  3.2.11 IRAS 22516+0838

The OH maser profile is basically double-peaked, but it is a rather complex
profile with multiple peaks on both the red-shifted and blue-shifted sides
(Eder et al., 1988; Wolak et al., 2012). The OH 1667 MHz line is also detected
at a velocity that deviates from the velocity range of the OH 1612 MHz line,
which suggests that there is fast molecular gas motion inside the OH 1612 MHz
line emission region (Silverglate et al., 1979). The H2O maser profile is
double-peaked and deviates from the OH maser velocity range by about 1.7 km
s-1 (Kim et al., 2010). The SiO $v=1$ and 2, $J=1-0$ lines have also been
detected (Nakashima & Deguchi, 2007). The SiO maser profile shows a strong
peak around the center in the OH maser velocity range, but also broad, weak
fast-wing components are seen in the spectra. Although the deviation of the
H2O maser velocity from the OH maser velocity range is small, the profiles of
the OH main lines and SiO maser line suggest that there are kinematic
components inside the envelope that deviate from the spherically expanding
motion. This source satisfies condition (2).

##  4 Discussions

###  4.1 Problems related to making maser source pairs for velocity comparison

In our analysis we compared the velocity ranges of the H2O and OH maser lines
for 229 sources. This looks like a relatively small number, considering that
there are 2,195 OH (1612 MHz) maser sources and 3,642 H2O (22.235 GHz) maser
sources in the databases. Only 10.4% of the total number of OH maser sources
and 6.2% of the total number of H2O maser sources could be paired for velocity
comparison. Of course, it is known that the molecular species of the detected
maser lines differ to some extent depending on the evolutionary stage of the
evolved star (i.e., Lewis’ chronology; see, e.g., Lewis, 1989; Habing, 1996).
However, such an explanation alone cannot account for the small overlap
between OH and H2O maser detections.

The low rate of matched pairs compared to the total number of maser sources is
due to several reasons. First, for many OH maser sources discovered by blind
surveys, follow-up observations to detect H2O masers are not available.
Second, the number of stars with OH masers that also exhibit H2O maser
emission is only about 50% (Engels & Lewis, 1996). Conversely, not all stars
with H2O maser emission have OH masers, especially stars with blue IRAS
[12]–[25] colors, where the OH masers are too weak, if present at all, to be
detected (Lewis & Engels, 1995). Finally, pairing the velocities of OH and H2O
maser sources requires reasonably accurate coordinates for a match with an
AllWISE and/or 2MASS counterpart. Most of the observations in the database
were made with single-dish telescopes, resulting in low angular resolution of
up to 1′ in some cases. In addition, many survey observations of maser sources
in the 1990s and 2000s used coordinates from the IRAS Point Source Catalog
(PSC), which are known to have uncertainties over 30′′ in some cases, as noted
in studies such as Deguchi et al. (1998).

Another point we should note is that in this study we have only selected maser
sources that have a clear near-infrared point source as their counterpart. In
the future, we will need to confirm the presence of WF jets with a high-
resolution radio interferometer, such as VLBI, to confirm whether our selected
sources have WF status. To do this, we need precise coordinate information (at
least within a few arcseconds). Therefore, as explained in Section 2.3, we
screened out those sources that had multiple possible counterparts within
20′′–30′′ and could not be narrowed down to a single counterpart. As the
angular resolution of OH and H2O maser observations improves in the future
through increased opportunities for radio interferometric observations, the
problems described here may disappear.

###  4.2 Velocity comparison OH and H2O maser lines and their profiles

The first point to consider is the classification of the line profiles given
in the OH maser source database. Due to limitations on the amount of
information that can be included as text in the database, it is obviously not
possible to provide extensive details on the line profiles.

As mentioned in Section 2.4, each OH maser source in the database is
classified as either double peaked (D) or single peaked (S). When we examined
the spectra in the literature, we found that sources classified as “S” have
simple, single peaks in fact. The “D” sources, on the other hand, exhibit a
wide variety of line profiles. As stated in Section 2.1, for the purpose of
finding WF candidates, we expect double peaks to be caused by the receding and
approaching sides of a spherically expanding envelope. However, upon further
examination of the literature, we found that a significant number of OH maser
sources classified as “D” do not fit this expectation.

When the velocity interval between the two peaks is as small as a few km s-1,
both peaks may be associated with either side of a spherically expanding
envelope, since such a small interval cannot be explained as the expansion
velocity of an AGB envelope.

The line profile is very helpful in assessing such a situation. In the case of
a typical double peak profile associated with the receding and approaching
sides of a spherically expanding envelope, the side closer to the systemic
velocity of each peak will have a gradual decrease in intensity, while the far
side will be steeply truncated. With a good signal-to-noise ratio, it may be
possible to distinguish between these cases by looking at the spectrum alone.
However, if the signal-to-noise ratio is not high enough, it may be difficult
to tell from the OH maser spectrum alone whether we are looking at a
spherically expanding envelope with a slow expansion velocity or the receding
or approaching side of a spherically expanding envelope. Even in this case, we
have found that a comparison between the OH and SiO maser lines is useful to
consider the situation. SiO maser emission lines are often detected near the
systemic velocity of the CSEs (see, e.g., Jewell et al., 1991). Thus, for
example, if both OH double peaks lie to one side of the SiO peak velocity, we
can assume that both are emitted from either the receding or the approaching
side of a spherically expanding envelope.

Another point to note here is that there may be a significant number of
velocity components in the spectra that are missing from the database records.
From our inspection of spectra in the literature, there are two cases where
they are missing. One is when the intensity of the emission lines is so weak
that they are missed. The other is a time variation of the maser line profile.
In fact, when examining the spectra, faint features are often found outside
the velocity range given in the database. We have also found cases where the
line profiles vary considerably with time (e.g. IRAS 19083+0851, IRAS
19229+1708, IRAS 19422+3506, etc.). For sources with large variations, there
is naturally a greater chance of missing high velocity components outside the
OH velocity range. The profiles of OH and H2O maser lines emitted from the
CSEs of evolved stars can vary on timescales of about a month (see, e.g.,
Etoka & Le Squeren, 2000; Engels, 2002; Etoka et al., 2003; Kim et al., 2016).
Therefore, to fully resolve this problem, it would be necessary to perform
monitoring observations for all sources with time intervals of several months,
leaving aside the question of whether this can actually be done with existing
instruments.

In addition, we note the potential usefulness of the OH maser main lines
(i.e., the OH 1665 MHz and 1667 MHz lines). Since the pumping route of the
main lines are different from the satellite line, the main lines and satellite
lines are thought to trace parts of the circumstellar envelope in different
physical conditions. Therefore, comparing the OH 1612 MHz and OH 1665/1667 MHz
lines provides similar information to comparing the OH 1612 MHz line and the
H2O maser line. For example, as discussed in Section 3.2.1, IRAS 18251-1048
has a velocity range for the 1667 MHz line that exceeds the velocity range for
the 1612 MHz line. The OH 1612 MHz and 1665/1667 MHz lines can usually be
observed simultaneously with a single instrument, which requires less
observing effort than comparing the OH and H2O maser lines. As more radio
interferometers become available to observe the OH maser lines in the future,
this comparison may be a useful tool for finding evolved stars that are out of
spherical symmetry.

The issues discussed here with the OH maser database show the limitations and
emphasize the need to evaluate the original maser spectra to improve the
velocity ranges obtained from the database. Due to the chosen database format,
which gives the velocities of the two strongest maser peaks (for masers
classified as ”D”), the full complexity of the spectra is not adequately
reproduced and the OH maser velocity range is sometimes underestimated.

###  4.3 Confirmation of the methodology using known WFs

We have tested what happens when known WFs are run through our screening
process to consider what conditions affect the completeness of the screening.
However, in this discussion we will not consider the classical WF condition
that the velocity range of the H2O maser emission exceeds 100 km s-1. Since
the velocity range of the H2O maser is always considered first, known WFs will
never be missed if this condition on the H2O velocity range is included. In
the following discussion, we will only consider the comparison of the H2O
maser velocity range with the OH maser, as applied in the present work. In
this case, the results of the screening showed that only 7 out of 16 known WFs
passed. The passed objects include IRAS 15544–5332, IRAS 16342–3814, IRAS
18043–2116, IRAS 18286–0959, IRAS 18450–0148, IRAS 18455+0448, and IRAS
18596+0315.

The 9 known WFs did not pass our screening for the following three reasons.
The first reason is that the corresponding OH maser source is not found (i.e.,
the 1612 MHz OH maser line has been observed in the past but not detected),
and three WFs fall into this case: i.e. IRAS 16552–3050, IRAS 17291–2147, and
IRAS 19134+2131.

The second reason is that the velocity range in the database does not meet the
WF requirement. IRAS 18460–0151 did not pass screening for this reason. As
mentioned in Section 4.2, the velocity range information in the database is
not always complete due to time variations in the profiles and other reasons.
It is clear that inspection of line profiles, as we did in Section 3.2, are
important to reduce this problem.

The third reason was that no infrared counterpart was found, or we could not
converge the number of infrared counterparts into one object. The following
five objects fit this reason: IRAS 15103–5754, IRAS 15445–5449, IRAS
18113–2503, IRAS 18139–1816, and IRAS 19190+1102. As discussed in Section 4.2,
this problem is primarily due to the listing of poor coordinates in the
databases, which were taken from the papers describing the discovery
observations.

The above results suggest that to improve the completeness of our screening
method, (1) the H2O and OH maser observations must be fully achieved for all
sources, (2) multiple observations must be made for a single source to deal
with temporal variations in maser intensity and profile, (3) radio
observations should provide accurate source positions for comparison with
infrared data.

###  4.4 Distribution of selected sources on the IRAS two-color diagram

Many previous studies have often assumed that the structure of the CSE begins
to deviate from spherical symmetry in the post-AGB phase. In Figure 8, the
IIIb region adjacent to the IIIa region (AGB region) appears to correspond to
the early post-AGB region. After careful checking of the spectra, as mentioned
in Section 3.2, one source (i.e., IRAS 19319+2214) remained in this region
that is certain to exhibit the velocity deviation of the H2O maser line. Given
the uncertainty of the IRAS 60 $\mu$m flux, we can assume that another source
IRAS 19069+0916 is also substantially in the IIIb region. No radio
interferometric observations have been made for IRAS 19319+2214 in the past.
We note that the profile of the H2O maser line of this source is similar to
that of known WFs (see, Section 3.2.9). The H2O maser line of IRAS 19069+0916
was observed with the VLA (Vlemmings et al., 2014). However, due to the time
variation of the line profile, the timing of the VLA observation was not
appropriate to confirm the WF status (see, Section 3.2.6 for details).

The II and IIIa regions in Figure 8 contain mainly (oxygen-rich) AGB stars. As
mentioned above, since previous studies have often assumed that the structure
of the CSE begins to deviate from spherical symmetry in the post-AGB phase,
the identification of WF candidates in these color regions must be carefully
considered. Our screening selected 27 sources in the II and IIIa regions.
After carefully inspecting the spectra, no sources remained in Region II that
showed a clear velocity deviation of the H2O maser line. In Region IIIa, on
the other hand, six sources remained with a clear deviation of the H2O maser
emission. The six sources are IRAS 19052+0922, IRAS 19068+1127, IRAS
19083+0851, IRAS 19229+1708, IRAS 19422+3506, and IRAS 22516+0838. We would
first need to verify if these six sources are AGB stars.

One of the observational characteristics of AGB stars is that many of them are
pulsating variables (i.e., Mira-type variables). To confirm whether the star
is pulsating, we checked the AAVSO light curves and the light curves of the
visible-near-infrared monitoring observations of the Arecibo samples by
Jiménez-Esteban et al. (2021): See, Table 3. First, the four sources IRAS
19052+0922, IRAS 19068+1127, IRAS 19422+3506, and IRAS 22516+0838 have been
confirmed as LPLAVs (long-period and large-amplitude variables) in Jiménez-
Esteban et al. (2021) and are considered AGB stars. For IRAS 19422+3506, a
mapping observation in the 22.235 GHz H2O maser line has been made with JVLA
(Vlemmings et al., 2014), and the non-spherically symmetric structure of the
envelope has been confirmed. IRAS 19422+3506 is a pulsating AGB star, but the
fact that the CSE has already deviated from spherical symmetry is curious. No
WF-like bipolar jet has been identified, but another interferometric
observation for confirmation would be desirable. IRAS 22516+0838 is also an
AGB star, but it is an interesting object with a similar H2O maser profile to
the known WFs.

On the other hand, IRAS 19083+0851 and IRAS 19229+1708 are not considered to
be AGB stars because no clear pulsation are observed. These two sources are
classified as Group 4 in Jiménez-Esteban et al. (2021), not post-AGB (Group
2), and the authors of this paper suggest that these two sources are RSGs.
However, one should be careful with the classification of these sources.
Jiménez-Esteban et al. (2021) proposes that IRAS 19083+0851 is an RSG because
it shows no pulsation and has a wide OH maser velocity range. However, these
properties are also seen in post-AGB stars. It has already been noted above
that post-AGB stars do not show pulsation, and a wide range of OH maser
velocities such as $>$60 km s-1 has indeed been detected in post-AGB stars
(Deacon et al., 2004). As for the other source, IRAS 19229+1708, the question
also remains whether it shows a pulsation or not, as mentioned in (Deacon et
al., 2004). In fact, it should be noted that in Gaia DR3, this source is
classified as an LPV (long period variable) with a period of 908 days.

With respect to the possibility of the presence of post-AGB stars in the IIIb
region, we would like to note that the infrared color of the WF may depend on
the direction of observation due to the non-spherically symmetric structure of
the WF. In general, the conventional WFs selected only for their H2O maser
velocity range above 100 km s-1 are observed from an angle relatively close to
the direction of the jet. In contrast, the selection method used in this study
may have selected WFs with a larger angle between the line-of-sight direction
and the jet direction than WFs selected by the conventional selection method.
For example, assuming that the cooler dust component ejected as a jet is in
the foreground, obscuring the relatively warmer component in the background,
this could possibly explain the presence of post-AGB stars in the IIIb region
and the redder infrared color of a conventional WF viewed from a direction
closer to the jet. It is beyond the scope of this study and not addressed
here, but the relationship between the structure of the dust envelope of the
WF, the direction of observation, and the infrared color needs to be confirmed
by radiative transfer calculations.

There are 3 Group A sources in regions IV, VIb, and VIII, corresponding to
more evolved sources (see the red dots in Figure 8; note that IRAS 19069+0916
is practically located in the IIIb region based on the uncertainty of the 60
$\mu$m flux.). Although these sources may have evolved to the planetary nebula
phase in terms of color, it should be noted that the OH maser profiles of
these four sources show the double peak profile typically seen in AGB/post-AGB
stars (see, Sections 3.2.1, 3.2.2 and 3.2.4). Since there is a large
uncertainty in determining the evolutionary stage by mid-infrared colors, the
possibility that these sources still have a spherically expanding AGB/post-AGB
envelope cannot be completely ruled out. However, note that there are
phenomena such that, once a source has evolved into a planetary nebula, it
temporarily undergoes AGB-like mass ejections again (Cohen et al., 2005). In
any case, these objects are unlikely to be very young WFs.

###  4.5 Possibility of finding by-products

Our initial goal in the present data analysis was to search for WF candidates.
However, a closer examination of the spectra for individual sources shows that
our selection is likely to select additional interesting sources other than
WFs. A good example is K3-35, shown on the right side of Figure 8. This source
is known to be a planetary nebula with maser emission (see, e.g., Miranda et
al., 2001), so its identity is different from that of the WF. However, in
common with the WF, the molecular envelope has a deviation from spherical
expansion.

We would also like to point out that some sources that meet the condition for
velocity screening emit SiO maser lines with a relatively wide linewidth
larger than $\sim$15–20 km s-1 (e.g. IRAS 04396+0647, IRAS 15408–5413, IRAS
19083+0851). SiO maser lines in AGB stars usually have a narrow linewidth,
typically a few km s-1 (see e.g., Jewell et al., 1991). SiO maser lines with a
linewidth larger than 10 km s-1 are often detected in red supergiants (RSGs;
see e.g., Deguchi et al., 2010a; Verheyen et al., 2012). RSGs are evolved
stars with a large initial mass, and the dynamical structure of their
circumstellar molecular envelopes is often more complicated than that of AGB
stars (see e.g., Kamiński et al., 2013). Although it is arguable whether it is
possible to determine RSG status from maser profiles alone, there is no doubt
that a non-negligible number of RSGs can be included in the sample for the
purpose of WF search.

Then IRAS 19312+1950, although excluded from the present study, also fulfills
the velocity condition of our screening (see, e.g., Nakashima et al., 2011).
The identity of this maser source remained unknown for a long time (see, e.g.,
Nakashima et al., 2015, 2016), but recent research has pointed to the
possibility of a red nova remnant (Qiu et al., 2023), which is an explosive
event that occurs when two stars merge. And importantly, it has been known
that a cold molecular envelope is formed after the red nova explosion (see
e.g., Kamiński et al., 2015, 2021). The red nova remnants surrounded by a cold
molecular envelope are often difficult to distinguish from evolved stars from
an observational point of view. However, it is easy to speculate that the
molecular envelope of red nova remnants exhibits motion that deviates from a
simple spherical expansion, since it is formed by the merger of two stars.
Indeed, a radio interferometric observation in molecular lines (Nakashima &
Deguchi, 2005) and a near-infrared imaging observation (Murakawa et al., 2007)
revealed the highly complex structure and motion of the envelope of IRAS
19312+1950.

Thus, as described above, many of the objects that may be selected for our
screening, even if they are not WFs, are interesting from the point of view of
stellar physics and stellar evolution. Therefore, the task of constructing a
large list of circumstellar maser sources, regardless of whether they are WFs
or not, where the velocity range of the H2O maser line exceeds that of the OH
maser line, would be a scientifically significant project.

##  5 Summary

In this study, databases of circumstellar OH (1612 MHz) and H2O (22.235 GHz)
maser sources were used to select WF candidates on the basis of a comparison
of the velocity ranges of OH and H2O maser emission. By comparing the velocity
ranges of the OH and H2O maser emission, it is in principle possible to detect
WFs with a smaller velocity range for the H2O maser emission than for the
previously known WFs. A cross-check between the 2,195 OH maser sources and the
3,642 H2O maser sources was performed to compare the velocity ranges of OH and
H2O maser emission. As part of the selection process, we examined the WISE and
2MASS infrared images and excluded sources with large, extended nebulae, as
they are most likely to be YSOs. We also excluded sources that could not be
clearly identified with their infrared counterparts due to uncertainties in
the coordinate information. In the first stage of selection, the OH and H2O
maser velocity ranges were compared for 229 sources, and 41 sources met the
velocity criterion. Then, after a detailed examination of maser line profiles
available in the literature, we concluded that the deviation of the H2O maser
velocity range was significant for 11 sources among the 41 initially selected
sources. The main results are as follows:

  * •

We examined the IRAS colors of the samples and found that two of the 11
sources with a confirmed H2O maser velocity deviation (IRAS 19069+0916 and
IRAS 19319+2214) are in the color region for post-AGB stars. The H2O maser
profile of these sources are similar to that of known WFs.

  * •

Of the 11 selected sources, six sources were located in the color region of
the AGB stars. For two of the six sources (IRAS 19422+3506 and IRAS
22516+0838), the H2O maser properties are different from those of typical AGB
stars, and confirmation observations with radio interferometry are desired.

  * •

Among the 11 selected sources, three sources (IRAS 18251-1048, IRAS 18588+0428
and IRAS 19067+0811) were identified in color regions corresponding to more
evolved objects. The OH maser profiles of these three sources showed the
double peak profile AGB/post-AGB stars. This fact suggests that the infrared
colors may depend on the orientation of the object.

  * •

We also confirmed the possibility that sources exhibiting the velocity
deviation of the H2O maser line could include astrophysically interesting
sources other than WFs. Such objects could include, for example, peculiar
planetary nebulae with maser emission and stellar merger remnants.

We acknowledge the science research grants from the China Manned Space Project
with No. CMS-CSST-2021-A03, No.CMS-CSST-2021-B01. JN acknowledges financial
support from the ‘One hundred top talent program of Sun Yat-sen University’
grant no. 71000-18841229. YZ are grateful for financial supports from the
National Science Foundation of China (NSFC, Grant No. 11973099) and the
science research grants from the China Manned Space Project (NO. CMS-
CSST-2021-A09 and CMS-CSST-2021-A10). JJQ thanks for support from the NSFC
(No. 12003080), the Guangdong Basic and Applied Basic Research Foundation (No.
2019A1515110588), and the Fundamental Research Funds for the Central
Universities (Sun Yat-sen University, No. 22qntd3101). CHH is supported by the
Hong Kong Research Grants Council for GRF research support under grants
17326116 and 17300417. C.-H. Hsia also thanks Q. A. Parker and HKU for the
provision of his research post.

##  Appendix A Distribution of circumstellar maser sources

Figures 9 and 10 show the galactic coordinate distributions of the 1612 MHz OH
and 22 GHz H2O maser sources analyzed in this study.

Figure 9: Distribution of 6068 OH maser sources (1612 MHz line) on the
galactic coordinate. The data are taken from Engels & Bunzel (2015). The red
filled circles and blue crosses represent detections (2195 sources) and non-
detections (3873 sources), respectively. Figure 10: Distribution of 3642 H2O
maser sources (22.235 GHz line) on the galactic coordinate. The red filled
circles and blue crosses represent detections (890 sources) and non-detections
(2752 sources), respectively.

##  Appendix B Maser characteristics of Group B sources

This section summarizes the characteristics of the maser profiles for each
individual source for the objects classified as “Group B”. See Section 3.2 for
Group B classification criteria.

###  B.1 IRAS 04396+0647

The OH maser line profile shows a double peak (Lewis, 1994), but the velocity
difference between the two peaks is somewhat narrower than is typical for AGB
stars. Additionally, there is a fast component outside the two peaks, which is
not typical of the spherically expanding envelopes of AGB stars. The H2O maser
emission is also detected outside the two peaks of the OH maser emission, but
the velocity difference between the two peaks is small, about 2.3 km s-1
(Valdettaro et al., 2001; Cho et al., 2017). Because of the presence of the
velocity component of the OH maser emission outside of the two peaks, the
velocity range of the H2O maser does not deviate completely from the velocity
range of the OH maser emission. The SiO $v=1$ and 2, $J=1-0$ lines have been
detected (Cho et al., 2017), with the $v=2$ line showing three peaks and a
relatively wider linewidth of about 15 km s-1. The $v=1$ line has a similar
linewidth. The radial velocity of the SiO maser lines is almost the same as
that of the OH maser line. These maser characteristics are reminiscent of red
supergiants, but not of AGB stars or WF candidates.

###  B.2 IRAS 05027–2158

The OH maser line is detected, but it is very weak, as reported by Rudnitski
et al. (2010). It is recorded in the database as having a double peak at
velocities of $-$15.9 km s-1 and $-$0.2 km s-1, but there are no supporting
spectra found in the literature. The OH 1667 MHz line was detected near $-$30
km s-1 by Rudnitski et al. (2010). The H2O maser line has been observed
multiple times (Takaba et al., 1994; Shintani et al., 2008), exhibiting a
highly time-variable profile. According to the information in the database,
the velocity difference between the two peaks is 4.6 km s-1, but in the
spectra found in the literature, it varies up to about 7 km s-1. The SiO $v=1$
and 2, $J=1-0$ lines have also been detected (Cho et al., 1996; Kim et al.,
2010). It is a single peak profile with a narrow line width and appears to be
a typical SiO maser profile for an AGB star. However, the SiO $v=2$, $J=2-1$
line shows a double peak profile with peaks at approximately $-$25 km s-1 and
$-$35 km s-1 (Haikala et al., 1994). The OH maser line profile is not clear
due to the insufficient signal-to-noise ratio.

###  B.3 IRAS 05528+2010

The profile of the OH maser line is not the typical double peak, but rather a
complex shape with multiple peaks, as reported by te Lintel Hekkert et al.
(1989). The H2O maser emission profile is a simple single peak, as reported by
Neufeld et al. (2021), and it is located just a few km s-1 outside the
velocity range of the OH maser emission. The SiO maser lines have also been
observed many times, as reported by, for example, Kim et al. (2010); the
linewidth of the SiO $v=1$ and 2, $J=1-0$ lines is up to 10 km s-1, which is
somewhat broad for a typical AGB star. Given that the H2O maser velocity range
is small (6.0 km s-1) and the deviation from the OH maser velocity range is
also not large, we cannot conclude that the envelope deviates from spherical
expansion.

###  B.4 IRAS 07445–2613

The profile of the OH maser line is the typical double peak of a spherically
expanding envelope, as reported by Slootmaker et al. (1985). The SiO $v=1$ and
2, $J=1-0$ maser lines also show a single-peak profile, which is typical for
AGB stars, as reported by Kim et al. (2013). The peak velocity of the SiO
maser is roughly equal to the average of the velocities of the OH double
peaks. The H2O maser’s velocity range deviates slightly from that of the OH
maser, with a deviation of only 0.3 km s-1, as reported by Valdettaro et al.
(2001); Kim et al. (2013). The velocity deviation is relatively small, but it
is still present. Overall, the evidence is weak to conclude that the envelope
deviates from spherical expansion.

###  B.5 IRAS 07446–3210

Although it is listed in the OH maser source database as having a double-peak
profile, we were unable to find any literature containing the spectra. The SiO
$v=1$ and 2, $J=1-0$ maser lines and the H2O maser line have been observed
(Kim et al., 2010), but there are no special characteristics other than the
slightly wide linewidth of the SiO maser (approximately 15 km s-1).

###  B.6 IRAS 08357–1013

According to the database, the OH maser line has a double-peaked profile with
peak velocities of $-$3.9 km s-1 and 39.4 km s-1. However, we couldn’t find
any literature that confirms these peak velocities through spectra. The
spectrum in Slootmaker et al. (1985) confirms the peak on the red-shifted
side, but the peak on the blue-shifted side is missing. The SiO maser lines
have a relatively wide line width, around 10 km s-1 for the $v=1$, $J=1-0$
line and 7–8 km s-1 for the $v=2$, $J=1-0$ line (Kim et al., 2010).

###  B.7 IRAS 13001+0527

The database shows that the OH maser line has two peaks at 23.5 km s-1 and
24.0 km s-1. However, the spectrum in Etoka et al. (2003) shows another peak
at 11.5 km s-1. Based on the line profile, the peaks at 23.5 km s-1 and 24.0
km s-1 are associated with the receding component of a spherically expanding
envelope. Therefore, it is likely that the velocities of approaching and
receding components for this source should be around 11.5 km s-1 and 24.0 km
s-1. The velocity range of the H2O maser line (Valdettaro et al., 2001) also
falls between these two velocities. The SiO maser lines show the profile of a
typical AGB star with a single peak (see, e.g., Gómez-Garrido et al., 2020).
Based on the overall properties of the maser emission lines, it is likely that
this source is a typical AGB star, not a WF.

###  B.8 IRAS 15060+0947

The OH maser profile is that of a typical spherical expanding envelope with
double peaks, but there is some fine structure on the blue-shifted side
(Lewis, 1994). The database shows that the minimum and maximum velocities of
the H2O maser are $-$14.4 km s-1 and 3.0 km s-1 respectively. However,
additional peaks are seen around $-$3 km s-1 (Yung et al., 2013) and 16 km s-1
(Kim et al., 2010). The 16 km s-1 feature is marginal, but if it is a true
detection, it can be concluded that the velocity range of the H2O maser
deviates from that of the OH maser. Thus, confirmation observation of the H2O
maser is desired. The SiO $v=2$, $J=1-0$ line is detected with a narrow peak
near $-$12 km s-1 (Kim et al., 2010).

###  B.9 IRAS 15255+1944

The profile of the OH maser emission line is a typical double peak (Wolak et
al., 2012). The profile of the H2O maser line shows multiple peaks (Kim et
al., 2010), slightly deviating from the velocity range of the OH maser
emission (it only exceeds 0.4 km s-1). The profiles of the SiO $v=1$ and 2,
$J=1-0$ maser lines (Kim et al., 2010) show a single peak which is also within
the velocity range of the OH maser emission. Although the velocity range of
the H2O maser differs slightly from that of the OH maser, the evidence is weak
to determine that the velocity range of the H2O maser line deviates from that
of the OH maser line.

###  B.10 IRAS 15408–5413

According to the information in the database, the profile of the OH maser line
is classified as double-peaked with peaks at $-$105.2 km s-1 and $-$101.1 km
s-1. However, another peak around $-$60 km s-1 can be seen in the spectrum in
te Lintel Hekkert et al. (1991). If the $-$60 km s-1 peak is the red-shifted
component of the spherically expanding envelope, the velocity range of the H2O
maser emission is within that of the OH maser emission. Considering the $-$60
km s-1 peak, the expansion velocity of the spherically expanding envelope
should be about 20 km s-1. The H2O maser spectra in the literature (Walsh et
al., 2011; Gómez et al., 2015a) agree with the velocity information in the
database. The SiO $v=1$, $J=2-1$ maser line was detected (Le Bertre & Nyman,
1990). The linewidth of the SiO maser line is somewhat wide (about 25 km s-1),
which is more RSG-like than AGB-like.

###  B.11 IRAS 17343+1052

The OH maser line profile is a typical double peak (Lewis, 1994). The SiO
$v=1$ and 2, $J=1-0$ maser lines are detected and the line profile is a
typical single peak (Cho et al., 2017). The profile of the H2O maser line
(Yung et al., 2013; Cho et al., 2017) has a shape that extends over the
velocity range of the OH maser emission. According to the information in the
database, the velocity range of the H2O maser emission exceeds that of the OH
maser emission by only about 0.4 km s-1. We have not confirmed the possibility
of high velocity components in the H2O maser spectra. The evidence is weak to
conclude that the velocity range of the H2O maser deviates from that of the OH
maser.

###  B.12 IRAS 17392–3020

According to the information in the database, the profile of the OH maser line
is double-peaked with peaks at $-$24.9 km s-1 and 18.5 km s-1. The OH maser
spectra in the literature (Deacon et al., 2004) show another peak at $-$10 km
s-1. However, the $-$10 km s-1 peak has an unusually narrow line width
(probably 1 channel), so it is possible that it is a spurious line. The
profile of the H2O maser line (Deacon et al., 2007) is a double peak with
peaks at similar velocities to the OH maser line. No reports of SiO maser
detected. We believe the evidence is weak to conclude that the velocity range
of the H2O maser deviates from that of the OH maser.

###  B.13 IRAS 17540-1919

The database shows that the OH maser line has two peaks at 56.4 km s-1 and
64.0 km s-1. Although these peaks can be seen in the spectrum in te Lintel
Hekkert et al. (1989), the red shifted peak is faint. Judging from the
spectrum given in Cho et al. (2017), the profile of the H2O maser line is
single-peaked, and the peak velocity is consistent with the information in the
database. The SiO $v=1$, $J=1-0$ line is detected (Cho et al., 2017); the
profile of the SiO maser line is a typical single peak of an AGB star, and the
peak velocity is about the same as that of the H2O maser. The H2O maser
emission deviates from the velocity range of the OH maser line, but only by a
small amount (1.4 km s-1).

###  B.14 IRAS 17580–3111

In the OH maser spectra presented in the literature (te Lintel Hekkert et al.,
1991; Hu et al., 1994) four peaks can be identified in the velocity range from
about 0 km s-1 to 30 km s-1. According to the information in the database, the
velocity range of the H2O maser line is from 0.4 km s-1 to 21.3 km s-1.
However, in the spectra presented in the literature (Suárez et al., 2007;
Gómez et al., 2008), a strong peak was observed only around 21 km s-1, but no
emission was observed around 0 km s-1. The spectrum in Suárez et al. (2007)
shows the possibility of a weak emission around 47 km s-1, but it is marginal.
There are no reports of SiO maser lines being detected.

###  B.15 IRAS 18009–2019

The profile of the OH maser line presented in the literature (te Lintel
Hekkert et al., 1989; Valdettaro et al., 2001) is a typical double peak. The
velocity range also matches the information in the database. The velocity
range of the H2O maser line given in the database is $-$1.5 km s-1 to 21.0 km
s-1, but the spectra found in the literature (Takaba et al., 1994; Kim et al.,
2010) did not confirm the blue-shifted end (i.e. $-$1.5 km s-1). The OH maser
velocity range given in the database is from 3.0 km s-1 to 26.0 km s-1, but
the H2O maser spectra in the literature do not clearly show any velocity
excess. The SiO $v=1$ and 2, $J=1-0$ lines are detected, and the line profile
is a typical single peak with a peak velocity near 10 km s-1 (Balister et al.,
1977; Cho et al., 1996; Kim et al., 2010).

###  B.16 IRAS 18034–2441

Spectra confirming the OH maser line profile cannot be found in the
literature. According to the spectrum of the H2O maser line in Cho et al.
(2017), the line profile is a single peak with a peak velocity near $-$13 km
s-1, similar to the velocity in the database. No velocity components other
than the $-$13 km s-1 peak can be identified. The SiO $v=1$ and 2, $J=1-0$
lines have been observed twice with an interval of more than 15 years, but the
two profiles differ (Deguchi et al., 2004b; Cho et al., 2017). The spectrum
reported by (Cho et al., 2017) shows a double peak profile with peaks at about
2 km s-1 and 8 km s-1, while that reported by (Deguchi et al., 2004b) shows a
single peak profile with a peak at about 8 km s-1. Considering that the peak
velocities of the SiO maser lines are close to the systemic velocity of the
source, the feature of the H2O maser line presumably traces the approaching
side of a spherically expanding envelope. The deviation of the H2O maser line
from the velocity range of the OH maser line is small (0.5 km s-1).

###  B.17 IRAS 18050–0518

We could not find any literature with spectra of the OH maser line. The H2O
maser spectrum has a peak near $-$28 km s-1 and shows a relatively broad
emission from about $-$32 km s-1 to $-$21 km s-1 (Yung et al., 2013). The SiO
$v=1$ and 2, $J=1-0$ lines are detected, both peaking at $-$32 km s-1 (Deguchi
et al., 2012). The profiles of the SiO maser lines are a typical single peak
of AGB stars. According to the information in the database, the deviation of
the H2O maser emission line from the velocity range of the OH maser line is
small (0.3 km s-1), and the profile of the H2O maser line does not indicate
the possible presence of a high-velocity component.

###  B.18 IRAS 18052–2016

The OH maser line profile, as reported in various studies such as (te Lintel
Hekkert et al., 1989; Blommaert et al., 1994; Deacon et al., 2004; Wolak et
al., 2012), displays the typical double peak of a spherically expanding
envelope. According to the information in the database, the peak velocity of
the blue-shifted component of the OH maser line is 23.0 km s-1. Additionally,
the H2O maser profile exhibits a weak extended wing component near this OH
peak velocity, as reported in (Deacon et al., 2007; Yoon et al., 2014).
However, it is unclear from the spectra whether the H2O maser emission exists
outside the velocity range of the OH maser line. We note that the OH maser
spectrum in Deacon et al. (2004) reveals that the OH maser emission of another
evolved star (b298) is also detected in the same spectrum, and the velocity of
the red-shifted peak of b298 is very close to the blue-shifted peak of IRAS
18052–2016 (with a velocity separation of about 5 km s-1). Therefore, it is
possible that the velocity excess of H2O may be due to the emission of b298.
The SiO $v=1$ and 2, $J=1-0$ and 29SiO $v=1$, $J=1-0$ lines have been detected
according to the reference (Yoon et al., 2014). These lines exhibit a typical
single-peak profile of an AGB star. The peak velocity is around 53 km s-1,
which is roughly the average of the two peaks of the OH maser line. Due to the
lack of sensitivity in the observations, the velocity of the wing component of
the H2O maser is not clearly known, making it difficult to completely rule out
the possibility of the velocity deviation of H2O maser emission.

###  B.19 IRAS 18085+0752

The profile of the OH maser line is a typical double peak (Lewis, 1994). The
H2O maser line has been observed twice. The spectrum in (Cho et al., 2017)
shows a simple single peak profile with a peak near $-$68 km s-1, while the
spectrum in (Yung et al., 2013) shows a peak near $-$64 km s-1 in addition to
the $-$68 km s-1 peak. No other possible velocity components can be seen in
the H2O spectrum. The SiO $v=1$ and 2, $J=1-0$ lines are observed and only the
$v=2$ line is detected (Deguchi et al., 2010b). The fact that only the $v=2$
line is detected suggests that this is an AGB star relatively close to the
post-AGB phase (Nakashima & Deguchi, 2007). The deviation of the H2O maser
line from the OH maser velocity range is small (0.6 km s-1).

###  B.20 IRAS 18348–0526

The OH maser line profile is a typical double peak (te Lintel Hekkert et al.,
1989; Wolak et al., 2012). The OH 1665 MHz and 1667 MHz lines are also
detected and show double peaked line profiles with peak velocities nearly the
same as the OH 1612 MHz line (te Lintel Hekkert et al., 1989). In the H2O
maser spectra, two peaks can be seen at about 16 km s-1 and 36 km s-1 (Takaba
et al., 1994; Nakashima & Deguchi, 2007). The information in the database
indicates that the maximum velocity of the H2O maser exceeds the redshifted
peak of the OH maser line by 3.0 km s-1, but no spectra were available in the
literature to confirm this deviation. The SiO $v=0$, 1, 2 and 3, $J=1-0$ lines
and the 29SiO $v=0$, $J=1-0$ lines are detected (Nakashima & Deguchi, 2007).
The $v=0$ lines are thought to be thermal lines, but the line widths are not
much different from the other SiO maser lines.

###  B.21 IRAS 18436+4334

The OH 1612 MHz maser line is a typical double peak (te Lintel Hekkert et al.,
1989; Valdettaro et al., 2001). The OH 1665 MHz and 1667 MHz lines are also
detected, but the line profiles are different from the OH 1612 MHz line, with
velocity components of these lines closer to the system velocity. The H2O
maser spectrum shows a strong peak at about $-$22 km s-1 (Valdettaro et al.,
2001; Kim et al., 2010), but no fast components that would deviate from the
velocity range of the OH maser line can be identified. The SiO $v=1$ and 2,
$J=1-0$ lines and the SiO $v=1$, $J=2-1$ and $3-2$ lines are detected (Cho et
al., 1998, 2009; Kim et al., 2010). The line profiles are typical single peaks
for all lines. The deviation of the H2O maser line from the velocity range of
the OH maser line is small (0.8 km s-1).

###  B.22 IRAS 18494–0130

The profile of the OH maser line is a typical double peak (te Lintel Hekkert
et al., 1989). The H2O maser line is strongly detected at velocities close to
the blue-shifted peak of the OH maser line (Yung et al., 2013). The emission
can also be weakly detected in the red-shifted peak of the OH maser line,
although this is marginal. The SiO $v=1$ and 2, $J=1-0$ line is detected near
78 km s-1, and the line profile is a typical single peak (Jewell et al., 1991;
Deguchi et al., 2004b). According to the information in the database, the
deviation of the H2O maser line from the velocity range of the OH maser line
is only 0.8 km s-1.

###  B.23 IRAS 19395+1949

The OH maser profile is double peaked in the database classification. However,
we can speculate that both peaks are emissions associated with the approaching
side of a spherically expanding envelope, according to the maser spectra in
the literature (Lewis et al., 1990; Yung et al., 2013; Cho et al., 2017). The
reason is that the velocity interval between the two peaks ($\sim$3.2 km s-1)
is too narrow for an AGB envelope (Lewis et al., 1990), and both peaks are on
the blue-shifted side of the SiO maser line, which has a peak at about 44 km
s-1 (Cho et al., 2017). According to the information in the database, the
velocity range of the H2O maser line is 28.0 km s-1 to 56.0 km s-1, which
agrees with the spectrum in the literature (Yung et al., 2013). The red-
shifted end of the velocity range of the H2O maser is 56 km s-1, which is only
31 km s-1 away from the blue-shifted end of the velocity range (25 km s-1) of
the OH maser line. This difference in velocity can be converted to an
expansion velocity of 15.5 km s-1, which is within the range of expansion
velocities of normal AGB envelopes.

###  B.24 IRAS 19466+2751

The OH maser spectrum from the literature shows four marginal peaks (Lewis,
1994), and their velocities are not inconsistent with the information in the
database. However, it is difficult to tell from the spectra of Lewis (1994)
whether the OH maser detection is true because of the high noise level. The
detection of the H2O maser is clearly confirmed by the spectra in the
literature (Valdettaro et al., 2001), and the velocity is consistent with the
information in the database. If we believe the velocity range of the OH maser
line given in the database, the velocity range of the H2O maser emission would
deviate from that of the OH maser emission. However, the signal-to-noise ratio
of the OH maser spectrum is not good.

###  B.25 IRAS 20015+3019

The OH maser profile is classified as double peaked in the database. In the
spectra given in the literature, peaks are found at corresponding velocities
(Chengalur et al., 1993). However, the narrow velocity spacing between the two
peaks (1.8 km s-1) makes it unlikely that they are emitted from the receding
and approaching sides of a spherically expanding envelope. Presumably, both
were emitted from either the receding or the approaching side. According to
the spectra in the literature (Engels & Lewis, 1996), the H2O maser line was
detected at velocities consistent with the information in the database. The
difference between the velocity at the blue-shifted feature of the OH maser
line (5.5 km s-1) and the velocity at the red-shifted end of the H2O maser
(20.1 km s-1) is 14.6 km s-1, which is not above the expansion velocity of the
normal AGB envelope.

###  B.26 IRAS 20440+0412

The database information classifies the OH maser profile of this source as
double peaked, and the peak velocities are consistent with the spectra found
in the literature (te Lintel Hekkert et al., 1989; Lewis et al., 1990).
However, the velocities of the two peaks are close together at $-$54.8 km s-1
and $-$49.9 km s-1, making it unlikely that they are peaks associated with the
receding and approaching sides of a spherically expanding envelope. It is
speculated that both peaks are associated with either the receding or the
approaching side. In the literature spectrum (Yung et al., 2013; Cho et al.,
2017), the H2O maser profile is a single peak with a peak velocity near $-$50
km s-1. The SiO $v=1$ and 2, $J=1-0$ lines are detected (Cho et al., 2017).
The SiO maser profile is a single peak, similar to the H2O maser line, with a
peak around 50 km s-1. This is a relatively rare case where maser lines from
three molecules, OH, H2O, and SiO, are detected at approximately the same
velocity. However, the H2O maser profile shows no signs of a jet, such as
high-velocity components or multiple peaks.

###  B.27 IRAS 20547+0247

The OH maser profile in the literature spectrum is double-peaked (Sivagnanam
et al., 1990; Chengalur et al., 1993). The velocities of the peaks seen in the
spectrum also match the information in the database. However, the velocity
interval between the peaks is narrow, less than 5 km s-1, and both are
probably associated with either the receding or approaching side of the
spherically expanding envelope. In the OH 1667 MHz spectrum (Sivagnanam et
al., 1990), a weak emission feature appears to be detected around $-$63 km
s-1. If the detection of the $-$63 km s-1 component of the OH 1667 MHz line is
true, then the two peaks of the OH 1612 MHz line would be associated with the
approaching side of the envelope. The H2O maser profile is a simple single
peak with a peak velocity around $-$75 km s-1, as given in the database
(Zuckerman & Lo, 1987).

###  B.28 IRAS 22556+5833

According to the information in the database, the OH maser profile is
classified as double-peaked with peak velocities at $-$67.0 km s-1 and $-$59.0
km s-1. However, another peak is also seen near $-$45 km s-1 in the spectrum
of Ashimbaeva et al. (2017). In addition, strong temporal variations in
intensity are seen on a timescale of a few years. The profile of the H2O maser
line has a complex shape with multiple peaks over a wide velocity range (about
$-$62 km s-1 to $-$40 km s-1 or so, Takaba et al., 1994; Wouterloot et al.,
1993; Kim et al., 2010). The SiO $v=1$ and 2, $J=1-0$ lines were detected (Kim
et al., 2010). The SiO maser profile is a single peak. Given the presence of
the $-$45 km s-1 component, the velocity range of the H2O maser does not
deviate from that of the OH maser.

##  Appendix C Maser characteristics of other type sources

In this section, we summarize the maser properties of two of the 41 sources
selected in the process described in Section 2 that have been identified as
non-AGB and post-AGB star types.

###  C.1 IRAS 17443–2949

This source is classified as a planetary nebula in the SIMBAD object type
(see, e.g., Kohoutek, 1994). The information in the database shows that the
profile of the OH maser line is double-peaked with peak velocities at $-$16.7
km s-1 and $-$4.8 km s-1. Although the double-peak profile can be recognized
in the spectra (Gómez et al., 2008), the red-shifted side peak is weak. Most
velocity components of the H2O maser emission (Suárez et al., 2007; Gómez et
al., 2008) fall within the velocity range of the OH maser emission. However,
the H2O maser line shows an isolated velocity component in the vicinity of 1
km s-1. This velocity component deviates from the velocity range of the OH
maser line by 5.71 km s-1. There are no reports of SiO maser detection.

###  C.2 IRAS 19255+2123

This object is classified as a young planetary nebula (K 3-35, see, e.g., Volk
& Cohen, 1990; Miranda et al., 2001). According to information in the
database, the OH maser profile is classified as double-peaked, with peak
velocities of $-$3.0 km s-1 and 9.0 km s-1. In the spectrum given in Engels et
al. (1985), the OH maser line appears to be weakly detected in the vicinity of
20 km s-1 as well. The H2O maser spectrum shows about three peaks at
velocities around 20 km s-1 to 25 km s-1 (Engels et al., 1985). The velocities
at which the H2O maser line was detected are clearly outside the velocity
range of the OH maser line. No SiO maser emission lines were detected. A VLBI
observation in the H2O maser line has been made by Miranda et al. (2001). The
H2O maser line has been detected from a bipolar structure according to the
VLBI map of Miranda et al. (2001). The location where the H2O maser line was
detected is as far as 5000 AU from the central star.

##  Appendix D List of YSO candidates with fast outflows

Some of the sources classified as YSO candidates mentioned in Section 2.3 have
H2O maser velocity ranges exceeding 100 km s-1. For the purpose of avoiding
confusion with WF, these sources are summarized in Table 4.

Table 4: YSO candidates with fast outflows  Source name | Position |  $V_{\rm min}$ (H2O) |  $V_{\rm max}$ (H2O) | $\Delta V_{\rm H_{2}O}$
---|---|---|---|---
| R.A. | Dec. |  |  |
| (J2000.0) | (km s-1) | (km s-1) | (km s-1)
G21.80$-$0.12 | 18 31 21.3 |  $-$09 56 56 |  $-$53.1 | 164.7 | 217.8
NGC 7129 FIR | 21 43 01.3 | +66 03 37 |  $-$71.5 | 46.0 | 117.5
IRAS 21519+5613 | 21 53 39.2 | +56 27 45 |  $-$100.5 |  $-$0.5 | 100.0
IRAS 20081+3122 | 20 10 09.2 | +31 31 34 |  $-$59.0 | 63.5 | 122.5
IRAS 15452$-$5459 | 15 49 11.3 |  $-$55 08 51 |  $-$83.0 | 22.0 | 105.0
IRAS 19078+0901 | 19 10 15.6 | +09 06 08 |  $-$272.5 | 268.5 | 541.0
IRAS 19213+1424 | 19 23 40.0 | +14 31 06 |  $-$30.0 | 170.0 | 200.0
Orion$-$A | 05 35 14.5 |  $-$05 22 21 |  $-$119.0 | 67.5 | 186.5
IRAS 02232+6138 | 02 27 04.1 | +61 52 22 |  $-$117.0 | 6.0 | 123.0
NGC 6334I | 17 20 53.0 |  $-$35 46 57 |  $-$71.0 | 50.0 | 121.0
IRAS 02219+6152 | 02 25 40.5 | +62 05 52 |  $-$100.5 | 18.0 | 118.5
IRAS 18162$-$2048 | 18 19 11.9 |  $-$20 47 34 |  $-$93.5 | 19.0 | 112.5
Orion$-$S | 05 35 12.5 |  $-$05 24 11 |  $-$57.0 | 54.5 | 111.5
NGC 7538 IRS 11 | 23 13 44.8 | +61 26 53 |  $-$153.5 |  $-$46.5 | 107.0
IRAS 22543+6145 | 22 56 17.9 | +62 01 46 |  $-$67.0 | 39.5 | 106.5

##  Appendix E Reference of 22.235 GHz H2O and 1612 MHz OH maser observations
for known WFs

Table 5 summarizes the papers with 22.235 GHz H2O maser and 1612 MHz OH maser
spectra of known WFs. See Appendix G for the reference codes given in Table 5.

Table 5: Maser spectra references for known WFs.  Source name | 22.235 GHz H2O lines | 1612 MHz OH lines
---|---|---
IRAS 15103$-$5754 | SUA09 | LIN96
IRAS 15445$-$5449 | DEA07, TAF14 | DEA04, SEV02
IRAS 15544$-$5332 | DEA07, GOM17 | DEA04, SEV02
IRAS 16342$-$3814 | CLA09, LIK92 | LIN96, WOL12
IRAS 16552$-$3050 | SUA07, SUA08 |
IRAS 17291$-$2147 | GOM15A, GOM17 |
IRAS 18043$-$2116 | DEA07, TAF14, USC23 | DEA04
IRAS 18113$-$2503 | GOM15A, YOO14 |
IRAS 18139$-$1816 | BOB07, ENG02 | AMI11, BOB05
IRAS 18286$-$0959 | TAF14, YUN14 | CHE17
IRAS 18450$-$0148 | DEA07, LIK92 | DEA04, WOL12
IRAS 18455+0448 | YUN13, YUN14 | LEW01, CHE93
IRAS 18460$-$0151 | IMA13 | IMA13, WOL12
IRAS 18596+0315 | DEA07, ENG02, GOM15A, GOM17 | AMI11, DEA04, GOM17
IRAS 19134+2131 | GOM15A, LIK92, YUN14 |
IRAS 19190+1102 | DAY10, GOM15A | LIK89

##  Appendix F Machine-readable catalog of H2O maser sources

A catalog of the 890 H2O maser sources used in Section 2.3 is attached to this
paper as an electronic table. The electronic table consists of 11 columns
(see, Table 6). The 11 columns, in order from the beginning, give the serial
number, the source name, the right ascension used in the radio observations,
the declination used in the radio observations, the right ascension of the
WISE counterpart, the declination of the WISE counterpart, the minimum
velocity (low veloc), the maximum velocity (high veloc), the velocity range
(veloc width), the image inspection results (image-inspec), and literature
information are given. The equinox of right ascension and declination is
J2000.0. The unit for radial velocity is “km s-1”. If multiple observations
have been made for the same source, the maximum and minimum velocities are
calculated from all previous observations (see Section 2.3). The values of
right ascension and declination used for the radio observations were obtained
from the references marked with ”*”. Information on the references
corresponding to the literature codes is given in Appendix G. The results of
the infrared image inspection fall into three main categories. In the“image-
insepc” column, “point source” indicates that a point source with an evolved
star-like infrared color is found in the vicinity of the radio observation
location, “YSO” indicates that an extended infrared nebulosity is seen, and
“crowded region” indicates that many stars are crowded together and no single
counterpart can be identified. Additionally, “!” indicates that no
corresponding infrared objects were found in the vicinity of the location
where the radio observation was made.

Table 6: Descriptive table of H2O maser source catalog  Column | Unit/Format | description
---|---|---
index |  | Row number of the table.
source name |  | Name of the observed source.
R.A.(J2000) | hh mm ss.s | Right ascension of the observed source. (The rows are sorted by the right ascensions.)
DEC(J2000) | dd mm ss.s | Declination of the observed source.
WISE R.A. | hh mm ss.s | Right ascension of the WISE counterpart.
WISE DEC | dd mm ss.s | Declination of the WISE counterpart.
low veloc | km s-1 | Minimum H2O maser velocity of the source among all the observations.
high veloc | km s-1 | Maximum H2O maser velocity of the source among all the observations.
veloc width | km s-1 | Velocity difference between the maximum velocity and minimum velocity of the source.
image-inspec |  | Image inspection result. ”!” means no corresponding WISE object was found in the vicinity of the source.
reference |  | References of the observations. The observed positions are cited from
|  | the reference with ”*” sign.

##  Appendix G Reference codes used in Tables 1, 5 and machine-readable table

The correspondence between the literature and the reference codes given in the
last column of Table 1 is as follows. AMI10:Amiri et al. (2010), AMI11:Amiri
et al. (2011), BOB05:Boboltz & Marvel (2005), BOB07:Boboltz & Marvel (2007),
CHE17:Chen et al. (2017), CHE93:Chengalur et al. (1993), CHO10:Cho & Kim
(2010), CHO12:Cho & Kim (2012), CHO14:Cho et al. (2014), CHO17:Cho et al.
(2017), CLA09:Claussen et al. (2009), COH06:Cohen et al. (2006), DAY10:Day et
al. (2010). DEA07:Deacon et al. (2007), DEG04C:Deguchi et al. (2004a),
DEG07A:Deguchi et al. (2007), DEG10A:Deguchi et al. (2010a), DEG10B:Deguchi et
al. (2010b), DEG89:Deguchi et al. (1989), ENG02:Engels (2002), ETO17:Etoka et
al. (2017), FOK12:Fok et al. (2012), GOM08:Gómez et al. (2008), GOM15A:Gómez
et al. (2015a), GOM17:Gómez et al. (2017), GOM90:Gómez et al. (1990),
GOM94:Gómez et al. (1994), GRE04:de Gregorio-Monsalvo et al. (2004),
GVA14:Gvaramadze et al. (2014), IMA13:Imai et al. (2013), KIM10:Kim et al.
(2010), KIM13:Kim et al. (2013), LEP76:Lepine et al. (1976), LEP77:Lepine &
Paes de Barros (1977), LEW01:Lewis et al. (2001), LIK89:Likkel (1989).
LIN96:te Lintel Hekkert & Chapman (1996), NAK00:Nakashima & Deguchi (2000),
NAK03B:Nakashima & Deguchi (2003), NAK06:Nakashima & Deguchi (2006),
NAK07:Nakashima & Deguchi (2007), NAK16:Nakagawa et al. (2016), NAKA93:Nakada
et al. (1993), NEU17:Neufeld et al. (2017), NEU21:Neufeld et al. (2021),
OHN13:Ohnaka et al. (2013), PAS06:Pashchenko et al. (2006), SEV02:Sevenster
(2002), SJO02:Sjouwerman et al. (2002), SUA07:Suárez et al. (2007), SUA08:
Suárez et al. (2008), SUA09:Suárez et al. (2009), SUN07:Sunada et al. (2007),
TAF14:Tafoya et al. (2014), TAK01:Takaba et al. (2001), TAK94:Takaba et al.
(1994), USC23:Uscanga et al. (2023) VAL01:Valdettaro et al. (2001),
VLE14:Vlemmings et al. (2014), WAL11:Walsh et al. (2011), WAL14:Walsh et al.
(2014), YOO14:Yoon et al. (2014), YUN13:Yung et al. (2013), YUN14:Yung et al.
(2014), ZUC87:Zuckerman & Lo (1987).

## References

  * Amada et al. (2022) Amada, K., Imai, H., Hamae, Y., et al. 2022, AJ, 163, 85, doi: 10.3847/1538-3881/ac3fb6
  * Amiri et al. (2010) Amiri, N., Vlemmings, W., & van Langevelde, H. J. 2010, A&A, 509, A26, doi: 10.1051/0004-6361/200913194
  * Amiri et al. (2011) —. 2011, A&A, 532, A149, doi: 10.1051/0004-6361/201116649
  * Ashimbaeva et al. (2017) Ashimbaeva, N. T., Colom, P., Krasnov, V. V., et al. 2017, Astronomy Reports, 61, 16, doi: 10.1134/S1063772916120027
  * Balick & Frank (2002) Balick, B., & Frank, A. 2002, ARA&A, 40, 439, doi: 10.1146/annurev.astro.40.060401.093849
  * Balister et al. (1977) Balister, M., Batchelor, R. A., Haynes, R. F., et al. 1977, MNRAS, 180, 415, doi: 10.1093/mnras/180.3.415
  * Blommaert et al. (1994) Blommaert, J. A. D. L., van Langevelde, H. J., & Michiels, W. F. P. 1994, A&A, 287, 479
  * Boboltz & Marvel (2005) Boboltz, D. A., & Marvel, K. B. 2005, ApJ, 627, L45, doi: 10.1086/431931
  * Boboltz & Marvel (2007) —. 2007, ApJ, 665, 680, doi: 10.1086/519067
  * Chen et al. (2017) Chen, X., Shen, Z.-Q., Li, X.-Q., et al. 2017, MNRAS, 468, 3602, doi: 10.1093/mnras/stx622
  * Chengalur et al. (1993) Chengalur, J. N., Lewis, B. M., Eder, J., & Terzian, Y. 1993, ApJS, 89, 189, doi: 10.1086/191843
  * Cho et al. (2017) Cho, C.-Y., Cho, S.-H., Kim, S., et al. 2017, ApJS, 232, 13, doi: 10.3847/1538-4365/aa85e6
  * Cho et al. (2009) Cho, S.-H., Chung, H.-S., Kim, H.-G., et al. 2009, ApJS, 181, 421, doi: 10.1088/0067-0049/181/2/421
  * Cho et al. (1998) Cho, S.-H., Chung, H.-S., Kim, H.-R., et al. 1998, ApJS, 115, 277, doi: 10.1086/313083
  * Cho et al. (1996) Cho, S. H., Kaifu, N., & Ukita, N. 1996, A&AS, 115, 117
  * Cho & Kim (2010) Cho, S.-H., & Kim, J. 2010, ApJ, 719, 126, doi: 10.1088/0004-637X/719/1/126
  * Cho & Kim (2012) —. 2012, AJ, 144, 129, doi: 10.1088/0004-6256/144/5/129
  * Cho et al. (2014) Cho, S.-H., Kim, J., & Yun, Y. 2014, Journal of Korean Astronomical Society, 47, 293, doi: 10.5303/JKAS.2014.47.6.293
  * Cho et al. (2016) Cho, S.-H., Yun, Y., Kim, J., et al. 2016, ApJ, 826, 157, doi: 10.3847/0004-637X/826/2/157
  * Claussen et al. (2009) Claussen, M. J., Sahai, R., & Morris, M. R. 2009, ApJ, 691, 219, doi: 10.1088/0004-637X/691/1/219
  * Cohen et al. (2006) Cohen, M., Chapman, J. M., Deacon, R. M., et al. 2006, MNRAS, 369, 189, doi: 10.1111/j.1365-2966.2006.10279.x
  * Cohen et al. (2005) Cohen, M., Parker, Q. A., & Chapman, J. 2005, MNRAS, 357, 1189, doi: 10.1111/j.1365-2966.2005.08684.x
  * Day et al. (2010) Day, F. M., Pihlström, Y. M., Claussen, M. J., & Sahai, R. 2010, ApJ, 713, 986, doi: 10.1088/0004-637X/713/2/986
  * de Gregorio-Monsalvo et al. (2004) de Gregorio-Monsalvo, I., Gómez, Y., Anglada, G., et al. 2004, ApJ, 601, 921, doi: 10.1086/380757
  * Deacon et al. (2004) Deacon, R. M., Chapman, J. M., & Green, A. J. 2004, ApJS, 155, 595, doi: 10.1086/425329
  * Deacon et al. (2007) Deacon, R. M., Chapman, J. M., Green, A. J., & Sevenster, M. N. 2007, ApJ, 658, 1096, doi: 10.1086/511383
  * Deguchi et al. (1998) Deguchi, S., Matsumoto, S., & Wood, P. R. 1998, PASJ, 50, 597, doi: 10.1093/pasj/50.6.597
  * Deguchi et al. (1989) Deguchi, S., Nakada, Y., & Forster, J. R. 1989, MNRAS, 239, 825, doi: 10.1093/mnras/239.3.825
  * Deguchi et al. (2007) Deguchi, S., Nakashima, J.-i., Kwok, S., & Koning, N. 2007, ApJ, 664, 1130, doi: 10.1086/519154
  * Deguchi et al. (2004a) Deguchi, S., Nakashima, J.-I., & Takano, S. 2004a, PASJ, 56, 1083, doi: 10.1093/pasj/56.6.1083
  * Deguchi et al. (2010a) Deguchi, S., Nakashima, J.-I., Zhang, Y., et al. 2010a, PASJ, 62, 391, doi: 10.1093/pasj/62.2.391
  * Deguchi et al. (2012) Deguchi, S., Sakamoto, T., & Hasegawa, T. 2012, PASJ, 64, 4, doi: 10.1093/pasj/64.1.4
  * Deguchi et al. (2010b) Deguchi, S., Shimoikura, T., & Koike, K. 2010b, PASJ, 62, 525, doi: 10.1093/pasj/62.3.525
  * Deguchi et al. (2004b) Deguchi, S., Fujii, T., Glass, I. S., et al. 2004b, PASJ, 56, 765, doi: 10.1093/pasj/56.5.765
  * Eder et al. (1988) Eder, J., Lewis, B. M., & Terzian, Y. 1988, ApJS, 66, 183, doi: 10.1086/191252
  * Engels (2002) Engels, D. 2002, A&A, 388, 252, doi: 10.1051/0004-6361:20020483
  * Engels & Bunzel (2015) Engels, D., & Bunzel, F. 2015, A&A, 582, A68, doi: 10.1051/0004-6361/201322589
  * Engels & Lewis (1996) Engels, D., & Lewis, B. M. 1996, A&AS, 116, 117
  * Engels et al. (1985) Engels, D., Schmid-Burgk, J., Walmsley, C. M., & Winnberg, A. 1985, A&A, 148, 344
  * Etoka et al. (2017) Etoka, S., Gérard, E., Richards, A. M. S., et al. 2017, MNRAS, 468, 1703, doi: 10.1093/mnras/stx448
  * Etoka & Le Squeren (2000) Etoka, S., & Le Squeren, A. M. 2000, A&AS, 146, 179, doi: 10.1051/aas:2000268
  * Etoka et al. (2003) Etoka, S., Le Squeren, A. M., & Gerard, E. 2003, A&A, 403, L51, doi: 10.1051/0004-6361:20030578
  * Fok et al. (2012) Fok, T. K. T., Nakashima, J.-i., Yung, B. H. K., Hsia, C.-H., & Deguchi, S. 2012, ApJ, 760, 65, doi: 10.1088/0004-637X/760/1/65
  * Gómez et al. (2015a) Gómez, J. F., Rizzo, J. R., Suárez, O., et al. 2015a, A&A, 578, A119, doi: 10.1051/0004-6361/201526009
  * Gómez et al. (2008) Gómez, J. F., Suárez, O., Gómez, Y., et al. 2008, AJ, 135, 2074, doi: 10.1088/0004-6256/135/6/2074
  * Gómez et al. (2017) Gómez, J. F., Suárez, O., Rizzo, J. R., et al. 2017, MNRAS, 468, 2081, doi: 10.1093/mnras/stx650
  * Gómez et al. (2015b) Gómez, J. F., Suárez, O., Bendjoya, P., et al. 2015b, ApJ, 799, 186, doi: 10.1088/0004-637X/799/2/186
  * Gómez et al. (1990) Gómez, Y., Moran, J. M., & Rodríguez, L. F. 1990, Rev. Mexicana Astron. Astrofis., 20, 55
  * Gómez et al. (1994) Gómez, Y., Rodríguez, L. F., Contreras, M. E., & Moran, J. M. 1994, Rev. Mexicana Astron. Astrofis., 28, 97
  * Gómez-Garrido et al. (2020) Gómez-Garrido, M., Bujarrabal, V., Alcolea, J., et al. 2020, A&A, 642, A213, doi: 10.1051/0004-6361/202037499
  * Gvaramadze et al. (2014) Gvaramadze, V. V., Menten, K. M., Kniazev, A. Y., et al. 2014, MNRAS, 437, 843, doi: 10.1093/mnras/stt1943
  * Habing (1996) Habing, H. J. 1996, A&A Rev., 7, 97, doi: 10.1007/PL00013287
  * Haikala et al. (1994) Haikala, L. K., Nyman, L. A., & Forsstroem, V. 1994, A&AS, 103, 107
  * Höfner & Olofsson (2018) Höfner, S., & Olofsson, H. 2018, A&A Rev., 26, 1, doi: 10.1007/s00159-017-0106-5
  * Hu et al. (1994) Hu, J. Y., te Lintel Hekkert, P., Slijkhuis, F., et al. 1994, A&AS, 103, 301
  * Imai (2009) Imai, H. 2009, in AGB Stars and Related Phenomena, ed. T. Ueta, N. Matsunaga, & Y. Ita, 62