# Deep optical study of the mixed-morphology supernova remnant G 132.7$+$1.3
(HB3)

P. Boumis,${}^{1\href https://orcid.org/0000-0002-4260-940X{}}$ A.
Chiotellis,${}^{1\href https://orcid.org/0000-0002-5535-2922{}}$ V. Fragkou,2
S.Akras,${}^{1\href https://orcid.org/0000-0003-1351-7204{}}$, S. Derlopa,1 M.
Kopsacheili,${}^{3,4\href https://orcid.org/0000-0002-3563-819X{}}$ I.
Leonidaki,${}^{3,4\href https://orcid.org/0000-0002-3563-819X{}}$ J.
Alikakos,1 E. V. Palaiologou3,4 E. Harvey,${}^{5\href
https://orcid.org/0000-0002-3014-3665{}}$ D. Souropanis${}^{1,6\href
https://orcid.org/0000-0002-6548-5489{}}$
1Institute for Astronomy, Astrophysics, Space Applications and Remote Sensing,
National Observatory of Athens, 15236 Penteli, Greece
2Instituto de Astronomía, Universidad Nacional Autónoma de México, Ensenada
22800, Baja California, Mexico
3Institute of Astrophysics, Foundation for Research and Technology-Hellas,
71110 Heraklion, Crete, Greece
4Department of Physics, University of Crete, GR-71003, Heraklion, Crete,
Greece
5Astrophysics Research Institute, Liverpool John Moores University, IC2,
Liverpool Science Park, Liverpool, L3 5RF, UK
6Department of Physics, National and Kapodistrian University of Athens,
Panepistimiopolis, 15784 Zografos, Greece
<EMAIL_ADDRESS>
(Accepted 2022 February 08. Received 2022 February 02; in original form 2021
July 08)

###### Abstract

We present optical ccd images of the large supernova remnant (SNR)
G132.7$+$1.3 (HB3) covering its full extent for the first time, in the
emission lines of H$\alpha+$[N ii], [S ii] and [O iii], where new and known
filamentary and diffuse structures are detected. These observations are
supplemented by new low-resolution long-slit spectra and higher-resolution
images in the same emission lines. Both the flux-calibrated images and spectra
confirm that the optical emission originates from shock-heated gas since the
[S ii]/H$\alpha$ $>$ 0.4. Our findings are also consistent with the recently
developed emission line ratio diagnostics for distinguishing SNRs from Hii
regions. A multi-wavelength comparison among our optical data and relevant
observations in radio, X-rays, $\gamma$-rays and CO bands, provided additional
evidence on the interaction of HB3 with the surrounding clouds and clarified
the borders of the SNR and the adjacent cloud. We discuss the supernova (SN)
properties and evolution that led to the current observables of HB3 and we
show that the remnant has most likely passed at the pressure driven snowplow
phase. The estimated SN energy was found to be $\left(3.7\pm 1.5\right)\times
10^{51}$ erg and the current SNR age $\left(5.1\pm 2.1\right)\times 10^{4}$
yrs. We present an alternative scenario according to which the SNR evolved in
the wind bubble cavity excavated by the progenitor star and currently is
interacting with its density walls. We show that the overall mixed morphology
properties of HB3 can be explained if the SN resulted by a Wolf-Rayet
progenitor star with mass $\sim 34\rm~{}M_{\rm\odot}$.

###### keywords:

ISM: general – ISM: supernova remnants – ISM: individual objects: G
132.7$+$1.3

††pubyear: 2022††pagerange: Deep optical study of the mixed-morphology
supernova remnant G 132.7$+$1.3 (HB3)–References

##  1 Introduction

Supernova remnants (SNRs) are the aftermath of supernova (SN) explosions, that
result from the interaction of the supersonically moving stellar ejecta with
the ambient medium. The resultant nebula is an excellent celestial lab from
which we can infer the mechanisms responsible for the explosive death of
certain stars, the shock wave physics and the local ambient medium properties.
In addition, SNRs are considered one of the most efficient cosmic-ray
accelerators while they substantially contribute on the chemical and dynamical
evolution of their host galaxies. For these reasons several targeted,
multiwavelength observational campaigns have been conducted aiming to decipher
the formation and evolution processes of these nebulae. While most SNRs have
been detected in radio wavelengths (see Dubner & Giacani 2015; Green 2017),
optical observation of SNRs allow us to inspect the SNR properties such as the
filamentary morphology, the chemical abundances and thermodynamical properties
of the shocked gas, the remnant’s expansion velocities and the density of the
ambient medium (e.g Boumis et al. 2008; Boumis et al. 2009; Alikakos et al.
2012; Stupar & Parker 2012; Sabin et al. 2013; Stupar et al. 2018; How et al.
2018; Fesen et al. 2020, 2021).

The SNRs in our Galaxy can be very extended structures depending on their
distance, their evolutionary stage, the characteristics of their parent stars
and the properties of the ISM that encloses them (see e.g. Asvarov 2014).
G132.7$+$1.3 (also known as HB3 radio source) is among the largest Galactic
SNRs observed so far 111Only recently optical SNRs were found with a few
degrees sizes at high Galactic latitudes; (Fesen et al., 2020, 2021) with an
angular size of $90\arcmin\times 123\arcmin$ in 408 MHz radio maps (Caswell,
1967) which implies a physical size of 52 $\times$ 72 $d_{2}$ pc, where
$d_{2}$ is the distance of the remnant in units of 2 kpc. The distance of HB3
has been suggested to be within the range of $\sim 1.6-2.4$ kpc (Routledge et
al., 1991; Fesen et al., 1995; Xu et al., 2006) and up to date it has been
determined by the distance of the adjacent Hii region/molecular cloud complex
W3.

HB3, was primarily detected in radio wavelengths by Brown & Hazard (1953)
during a 158.5 MHz survey for radio sources in the Milky Way. Further radio
observations (e.g. Williams et al. 1966; Caswell 1967) revealed the SNR nature
of this radio source distinguishing it from the adjacent Hii region W3. In the
radio band the remnant reveals a shell like morphology while its radio
spectral index has been determined to be between – 0.64 and – 0.56 (Landecker
et al., 1987; Fesen et al., 1995; Reich et al., 2003; Tian & Leahy, 2005;
Green, 2007; Shi et al., 2008). One pulsar discovered close to the SNR’s
boundaries seems not to be a part of it due to its large characteristic age
(Lorimer et al., 1998) and six masers detected in the line-of-sight of HB3 are
most probably associated with the adjacent W3 complex and not with the SNR
(Koralesky et al., 1998).

Optical emission associated with HB3 was first found in the western portion of
the remnant where a few filamentary structures were discovered (van den Bergh
et al., 1973; D’Odorico & Sabbadin, 1977). Low dispersion spectra of these
filaments exposed the shock heated origin of this emission and revealed a
filament’s density of $\leq$ 100 cm-3 and velocity of 35 – 40 $\rm
km~{}s^{-1}$ (D’Odorico & Sabbadin, 1977; Lozinskaya & Sitnik, 1980,
respectively). Fesen & Gull (1983) and Fesen et al. (1995) presented new $\rm
H\alpha$ images of HB3 covering the western two-thirds of the remnant. The
optical morphology of the remnant was found to be strongly spatially
correlated with the radio band. Optical follow-up spectroscopy on the western
limb filaments was also performed by Fesen et al. (1995), indicating a shock
velocity of $\leq 100~{}\rm km~{}s^{-1}$, an electron density of $\leq
150~{}\rm cm^{-3}$ and an E(B-V) = $0.71\pm 0.04$.

The large size of HB3 in conjunction with the strong radio to optical
correlation and the multi-shell appearance of the remnant, indicate a SNR of
advanced age that most likely has passed beyond the adiabatic phases of
evolution. However, the remnant is characterised by a centrally peaked thermal
X-ray emission (Venkatesan et al., 1984; Leahy et al., 1985; Rho & Petre,
1998; Lazendic & Slane, 2006). Such an evidence cannot be explained by the
canonical theory for SNR evolution, as in mature SNRs as HB3, the reverse
shock is expected to have been ceased long ago. Thus, the remnant’s interior
is expected to be characterised by a cold and low density gas. For this
reason, Rho & Petre (1998) classified HB3 as a mixed morphology (or thermal
composite) SNR (MMSNR). X-ray spectroscopy revealed that the central emitting
gas is close to ionization equilibrium characterised by a single temperature
(Lazendic & Slane, 2006). The element abundances of the X-ray emitting plasma
are still ambiguous. Lazendic & Slane (2006) suggested that either it is
characterised by enhanced abundances of O, Ne, Mg or it has marginally
enhanced abundances of Mg and under-abundant Fe.

Extended $\gamma$-ray emission is also discovered in the vicinity of the
remnant (Tsygankov et al., 2016; Katagiri et al., 2016). This emission has
been correlated to HB3 and it has been attributed to the decay of $\pi$0
produced by the interaction of hadrons -accelerated at the remnant’s shock
waves- with the surrounding interstellar gas (Katagiri et al., 2016).

Regarding the environment of the remnant, HB3 has been found to evolve within
and a rather complex medium where its eastern part is adjacent to the Hii
region / molecular complex W3 (e.g. Digel et al., 1996). Routledge et al.
(1991) presented radio observations on the 21 cm Hi emission covering the
whole region of the remnants and $\rm{}^{12}CO$ spectra in the southern
portion of HB3. A shell shaped structure superimposed on HB3 was found
indicating an atomic neutral gas acceleration by the SNR blast wave. In
addition, a bright $\rm{}^{12}CO$ emission near $-$43 km s-1 was detected
overlapping with the radio emission of HB3 and they interpreted this evidence
as a SNR interaction with molecular gas, in essence with W3. This evidence was
verified by multimeter observations of CO lines, conducted by Zhou et al.
(2016) which revealed a substantial amount of molecular gas around $-$43 km
s-1, being morphologically and dynamically correlated with HB3. Finally,
shocked molecular hydrogen ($\rm H_{2}$) and broad CO was recently detected in
the near/mid infrared and millimeter bands, respectively, whose morphological
and kinematic properties provide an additional evidence on the interaction of
the SNR with the surrounding clouds (Rho et al., 2021).

The origin and age of HB3 are still not well-determined. Lazendic & Slane
(2006) assuming that the SNR is in the adiabatic phase and adopting an
evaporating clouds model for MMSNRs, estimated the SN explosion energy to be
$7-20\times 10^{50}$ erg and the age of the SNR between 25$-$30 kyrs. On the
other hand, Zhou et al. (2016) considered that the remnant of HB3 has passed
into the radiative phase and came to the conclusion that it is approximately
$21\pm 2$ kyrs old and the explosion energy $E=\left(1.6\pm 0.9\right)\times
10^{51}$ erg. Regarding the progenitor mass, Zhou et al. (2016) estimated a
lower limit of 28 $\rm M_{\odot}$ by applying the linear relationship between
the size of the wind-blown and the progenitor mass suggested by Chen et al.
(2013).

In this paper, we explore the optical emission of the SNR G132.7$+$1.3 in its
full extent. In particular, H$\alpha+$[N ii], [S ii] and [O iii] emission line
images are presented, which show, for the first time, the complete structure
of this large SNR in the optical band. Moreover, we present optical higher
resolution imaging and long-slit spectroscopy at multiple areas of the SNR.
HB3’s optical emission is also compared with the radio, X–ray, $\gamma-$ray
and molecular emission. In Section 2, we describe our imaging and
spectroscopic observations, with the results and multi-wavelength comparison
presented in Section 3. We discuss the SNR’s properties, origin and evolution
in Section 4, while in Section 5, we summarize the conclusions acquired from
this work.

##  2 Observations

A brief summary and log of our observations are given in Table 1. In the
following sections, we describe the observations of HB3 in more detail.

###  2.1 Imaging

Table 1: Imaging and Spectral log.  WIDE- FIELD IMAGING
---
Filter | $\lambda_{\rm c}$ | $\Delta\lambda$ | Exposure time | Date
| (Å) | (Å) | (sec) |  |
H$\alpha+$[N ii] 6548, 6584 Å | 6570 | 75 | 2400 (7)222Numbers in parentheses represent the number of individual frames.[2]333Numbers in brackets represent the number of different fields in each filter. | 2009 Aug 19-20,22,25 | 2010 Aug 29-30
H$\alpha+$[N ii] 6548, 6584 Å | 6582 | 80 | 600 (22) [2] | 2019 Jul 27-29, Sep 30 | 2019 Oct 1-2
[O iii] 5007 Å | 5010 | 28 | 2400 (1)[1] | 2009 Aug 23 |
[O iii] 5007 Å | 5010 | 28 | 900 (25) [2] | 2018 Aug 7-8, 11, 17 | 2019 Sep 5, 15
[S ii] 6716, 6731 Å | 6720 | 27 | 2400 (5)[2] | 2009 Aug 20-21 | 2010 Aug 29-30
[S ii] 6716, 6731 Å | 6720 | 32 | 600 (17) [2] | 2019 Jul 27-29, Sep 30 | 2019 Oct 1-2
Continuum blue | 5470 | 230 | 180 [1] | 2009 Aug 23 |
Johnson V | 5380 | 980 | 20 (20) [2] | 2019 Jul 27-29, Sep 30 | 2019 Oct 1-2
Continuum red | 6096 | 134 | 180 [2] | 2009 Aug 19-22,25 | 2010 Aug 29-30
SDSS-r | 6214 | 1290 | 20 (20) [2] | 2019 Jul 27-29, Sep 30 | 2019 Oct 1-2
HIGHER-RESOLUTION IMAGING
Filter | $\lambda_{\rm c}$ | $\Delta\lambda$ | Exposure time | Area | Date
| (Å) | (Å) | (sec) |  |
H$\alpha$$+$[N ii] 6584 Å | 6578 | 40 | 1800 | S1, S2, S3, S4, S5, S6 | 2017 Oct 18-19
[S ii] 6716, 6731 Å | 6727 | 40 | 1800 | S1, S2, S3, S4, S5, S6 | 2017 Oct 18-19
[O iii] 5007 Å | 5011 | 30 | 1800 | S1, S2, S5, S6 | 2017 Oct 18-19
[O iii] 5007 Å | 5010 | 28 | 600 (11) | S1, S2, S3 | 2021 July 13-15
SPECTROSCOPY444The exposure time for all spectra is 3600 s.
Position | Slit centers | Offset555Spatial offset from the slit center in arcsec: N($=$North), S($=$South). | Aperture length666Aperture lengths for each area in arcsec. | Date
| $\alpha$ | $\delta$ |  |  |
| (h m s) | (° ′ ″) | (arcsec) | (arcsec) |
Slit 1a | 02 11 26.4 | 62 44 42.1 | 96 S | 86 | 2009 Aug 22
Slit 1b | 02 11 26.4 | 62 44 42.1 | 190 S | 29 | 2009 Aug 22
Slit 2a | 02 13 16.9 | 62 56 47.9 | 138 S | 41 | 2010 Sep 10
Slit 2b | 02 13 16.9 | 62 56 47.9 | 89 S | 10 | 2010 Sep 10
Slit 2c | 02 13 16.9 | 62 56 47.9 | 36 S | 14 | 2010 Sep 10
Slit 3a | 02 13 33.6 | 62 48 55.9 | 26 N | 37 | 2010 Sep 09
Slit 3b | 02 13 33.6 | 62 48 55.9 | 95 N | 27 | 2010 Sep 09
Slit 4a | 02 12 40.7 | 62 21 02.6 | 170 N | 63 | 2009 Aug 20
Slit 5a | 02 22 22.7 | 62 50 10.7 | 118 S | 80 | 2009 Aug 22
Slit 5b | 02 22 22.7 | 62 50 10.7 | 41 N | 117 | 2009 Aug 22
Slit 5c | 02 22 21.7 | 62 50 52.2 | 157 S | 96 | 2010 Sep 08
Slit 5d | 02 22 21.7 | 62 50 52.2 | 13 S | 61 | 2010 Sep 08
Slit 6a | 02 23 38.6 | 62 44 57.7 | 190 S | 29 | 2010 Sep 03
Slit 6b | 02 23 38.6 | 62 44 57.7 | 18 N | 33 | 2010 Sep 03
Slit 6c | 02 23 38.0 | 62 44 53.6 | 187 S | 23 | 2010 Sep 07
Slit 6d | 02 23 38.0 | 62 44 53.6 | 24 N | 31 | 2010 Sep 07
Slit 7a | 02 19 54.0 | 62 23 57.6 | 12 S | 80 | 2009 Aug 23
Slit 7b | 02 19 54.0 | 62 23 57.6 | 103 N | 8 | 2009 Aug 23
Slit 8a | 02 21 15.8 | 62 23 58.0 | 128 S | 43 | 2010 Sep 05
Slit 8b | 02 21 15.8 | 62 23 58.0 | 32 S | 18 | 2010 Sep 05

Figure 1: The continuum–subtracted mosaic of G132.7$+$1.3 in the light of
H$\alpha+$[N ii], [S ii] and [O iii] emission. The image centre is at R.A.
02h18m00s; Dec +62$\degr$45$\arcmin$00$\arcsec$ and its scale is
$\sim$3$\arcsec$ pixel-1.

Figure 2: The continuum–subtracted image of G132.7$+$1.3 in [S ii] emission.
The blue lines indicate the positions of the slits, while the cyan and magenta
rectangles indicate the position of the higher resolution images shown in Fig.
4. Shadings run linearly from 0 to 39$\times$$10^{-17}$ erg s-1 cm-2 arcsec-2
(the maximum measured [S ii] flux on the SNR is 27$\times$$10^{-17}$ erg s-1
cm-2 arcsec-2). Figure 3: The continuum–subtracted RGB image of G132.7$+$1.3,
combining the H$\alpha+$[N ii] (red), [S ii] (green) and [O iii] (blue)
emission line images. The yellow colour produced in regions where both
H$\alpha+$[N ii] and [S ii] exist.

####  2.1.1 Wide-field imaging

The wide-field imagery of HB3 was obtained with the 0.3 m Schmidt-Cassegrain
(f/3.2) telescope at Skinakas Observatory, Crete, Greece in two different
runs. The first was on the nights from 19th to 23rd and 25th of August 2009,
as well as, on the 29th and 30th of August 2010. The camera used, was a
$2048\times 2048$ Andor DW436 CCD, which has a pixel size of 13.5 $\mu$m that
results in a $101\arcmin\times 101\arcmin$ field of view and an image scale of
2″.96 pixel-1. In this run, the remnant was observed with the H$\alpha+$[N
ii], [S ii] and [O iii] filters with exposure times of 2400 s each and of 180
s for the continuum blue and red filters. Multiple exposures were taken in the
first two filters, resulting to a total observing time of 16800 s for the
H$\alpha+$[N ii] and 12000 s for the [S ii]. The second run was on the nights
7th, 8th, 11th and 17th of August and on the 5th and 15th of September 2018,
as well as, on the 27th – 29th of July, 30th of September and 1st – 2nd
October 2019. The camera was the same with the first run, resulting to the
same FOV and scaling. The filters used in this run were a new set of
H$\alpha+$[N ii] and [S ii], the same [O iii], all with sets of exposure times
of 600 s and the continuum SDSS-r and V with exposure time of 20 s (see Table
1 for details). As in the first run, multiple exposures were taken for all
filters, resulting to a total observing time of 13200 s for the H$\alpha+$[N
ii], 10200 s for the [S ii], 15000 s for the [O iii] and 400 s for the
continuum filters. During all the observations, the “seeing” varied between
0″.8 and 1″.2.

For the data reduction the IRAF and Montage packages were used and all frames
were bias-subtracted and flat-field corrected using a series of twilight flat-
fields, and they were all airmass extinction corrected. For the absolute flux
calibration the spectrophotometric standard stars HR 7596, HR 7950, HR 8634,
HR 9087, and HR 0718 (Hamuy et al., 1992, 1994) were used. The Hubble Space
Telescope (HST) Guide Star Catalogue (Lasker et al., 1999) was used for the
calculation of the astrometric solution for all data frames and all the
equatorial coordinates quoted in this work refer to epoch 2000. The
H$\alpha+$[N ii], [S ii] and [O iii] continuum subtracted mosaic images are
shown in Fig. 1, while in Fig. 2, we present the flux-calibrated [S ii] image.
In the latter image the blue lines indicate the positions of the slits, and
the cyan and magenta rectangles indicate the position of the higher resolution
images shown in Fig. 4. For a better comparison between the three emission
line images presented in Fig. 1, we created a colour RGB image shown in Fig.
3. In this figure, it is clear that most of the H$\alpha+$[N ii] filaments
coincide with the [S ii] ones, while there is a number of [O iii] filaments
which appear bright in this emission line and do not show any low ionization
emission.

####  2.1.2 Higher-resolution imaging

Follow-up, higher resolution images where obtained with the 2.3m (f/8)
Aristarchos telescope at Helmos Observatory in Peloponnese, Greece, on 18th
and 19th of October 2017 and with the 1.3m (f/7.64) Ritchey-Chretien telescope
at Skinakas Observatory on 13th to 15th of July 2021. These images cover the
regions of interest (hereafter areas S1 to S6) where spectroscopic
observations were performed (centered at slit positions 1-6). In particular,
S1–S6 were observed with the H$\alpha+$[N ii] and [S ii] filters with
Aristarchos telescope (cyan rectangles in Fig. 2), while areas S1, S2, S5, S6
and S1, S2, S3 were observed with the [O iii] filter with the Aristarchos
(cyan rectangles in Fig. 2) and Skinakas (magenta rectangles in Fig. 2)
telescopes, respectively.
The detector at Aristarchos telescope was a 2048$\times$2048, 13.5 $\mu$m
pixels CCD, with a field of view of 5′.5$\times$5′.5 (0″.16 pixel-1 in
2$\times$2 binning), while at Skinakas Observatory a 2048$\times$2048, 13.5
$\mu$m pixels Andor CCD was used which has a 9′.5$\times$9′.5 field of view
and an image scale of 0″.28 pixel-1. All information about the used
H$\alpha+$[N ii], [S ii] and [O iii] filters as well as the exposure times of
all higher-resolution images are shown in Table 1. The data reduction was
carried out using standard IRAF routine packages for the bias subtraction and
flat-field correction, while the same catalogue of stars was used for the
astrometric calculations. Seeing conditions were between 1″ to 1″.5.
The higher-resolution images are shown in Figs. 4, with two of them (S1 and
S6) being also presented in RGB colour images (Fig. 5). Details of all imaging
observations are given in Table 1.

Figure 4: Higher resolution images of G132.7$+$1.3 in H$\alpha+$[N ii], [S
ii]and [O iii] emission of the regions where slit positions 1-6 were taken
(see Fig. 2). The North-East orientation is the same for H$\alpha+$[N ii] and
[S ii] images and it can be seen on the top left image. The [O iii] images
have a larger FOV (Sect. 2.1.2), so the position of the H$\alpha+$[N ii], [S
ii] images are indicated with a dash rectangle

. The faint ring which appears to the north-east of all bright stars is a
ghost due to the filter.

Figure 4: Continued.

Figure 4: Continued.

Figure 5: Higher resolution RGB image of Areas S1 and S6, combining the
H$\alpha+$[N ii] (red), [S ii] (green) and [O iii] (blue) emission line
images. The yellow colour produced in regions where both H$\alpha+$[N ii] and
[S ii] exist. North is to the top, East to the left.

###  2.2 Spectroscopy

Low-dispersion, long-slit spectra where obtained with the 1.3 m Ritchey-
Cretien telescope (f/7.7), at Skinakas Observatory, on August 20, 22-23, 2009
and September 3, 5, 7-10, 2010. For these observations, the 1302 line mm-1
grating was employed with the $2000\times 800$ (13 $\mu$m) SITe CCD covering
the wavelength range from 4750 to 6815Å, while the exposure time for each data
frame was 3600 s. The spectral resolution being $\sim$8 Å and $\sim$11 Å for
the red and blue part of the spectra, respectively. The slit has a width of
7.7 arcsec and a length of 7.9 arcmin and it was oriented in the north-south
direction for all the positions. The coordinates of each slit are given in
Table 1. For the data reduction the IRAF package was used, while the
spectrophotometric standards stars HR 0718, HR 7596, HR 7950, HR 8634, and HR
9087 were used for the flux calibration of the spectra (Hamuy et al., 1992,
1994).

Figure 6: Long–slit low resolution spectra from different positions of the
observed area (see Table 1). The small shift in the wavelength axis is due to
the different observing season (slightly different position of the grating).
Figure 7: The correlation between H$\alpha+$[N ii] image and bf(a) the radio
emission at 4850 MHz (blue line) and 330 MHz (red line). The blue contours
scale linearly from 5$\times 10^{-3}$ Jy/beam to 0.03 Jy/beam, with step 0.05
Jy/beam and the red from 7$\times 10^{-3}$ Jy/beam to 4.5$\times 10^{-2}$
Jy/beam, with step 0.05 Jy/beam. (b) the X-ray emission of ROSAT PSPC (yellow
line, Lazendic & Slane 2006). The yellow contours scale linearly from 1$\times
10^{-5}$ counts arcmin-2 s-1 to 3$\times 10^{-4}$ counts arcmin-2 s-1, with
step 2.8$\times 10^{-5}$ counts arcmin-2 s-1. (c) the molecular emission of
12CO ($J=1-0$), 13CO ($J=1-0$) and C18O ($J=1-0$) in the velocity range of
$-$60 to $-$30 km s-1 (cyan line, Zhou et al. 2016). The cyan contours scale
linearly from 0 to 13.5 K with step 5 K. (d) the $\gamma-$ray emission of
background-subtracted FERMI LAT in the $1-10$ GeV (blue line, Katagiri et al.
2016).The blue contours scale linearly from 0 to 22 counts pixel-1 with step 5
counts pixel-1.

##  3 Results

Table 2: Typically measured fluxes over the brightest filaments with median
values over a 40″$\times$ 40″ box. Fluxes are in units of $10^{-17}$ erg s-1
cm-2 arcsec-2.

| Area S1 | Area S2 | Area S3 | Area S4 | Area S5 | Area S6 | Area S7 | Area S8 | Area A777Position of center of Area A: $\alpha$= 02h 23m 32s, $\delta$= + 62° 21′ 37″ .
---|---|---|---|---|---|---|---|---|---
H$\alpha+$[N ii] | 188 | 63 | 112 | 111 | 134 | 114 | 96 | 146 | 257
[S ii] | 54 |  $<$ 15 | 34 | 42 | 39 | 44 | 17 | 30 | 38
[O iii]8883$\sigma$ upper limit. |  |  |  |  |  $<$ 16 |  $<$ 13 |  $<$ 16 |  $<$ 16 |  $<$ 20
[S ii]/H$\alpha$ | 0.58 | 0.47 | 0.60 | 0.74 | 0.58 | 0.78 | 0.35 | 0.41 | 0.30

###  3.1 The optical emission line images

The optical emission of HB3 is detected for the first time in its full extent.
Both bright and faint, diffuse and filamentary structures spread along the
entire remnant are revealed, while the Hii region W3 is also noted at the
southeast region of the remnant. The [S ii] emission of HB3, though fainter,
displays a similar morphology to that of the H$\alpha+$[N ii] emission-line
image, tracing all the filamentary structures. The [O iii] emission line,
displays a different filamentary morphology compared to the former line
images, indicating different shock velocities and physical conditions.

H$\alpha+$[N ii], [S ii] and [O iii] fluxes are calculated from areas with
40x40 arcsec2 dimensions over the brightest filaments where the low resolution
spectra were taken (see Table 2). The errors of the integrated fluxes are
estimated between 10 and 20 percent. All selected slit areas but one (area S7)
and area A (where a known Hii region exists) exhibit [S ii]/H$\alpha>$0.4, as
expected for shock-heated regions.

The most interesting regions lie in the west, south-west and east directions,
where bright large filamentary and faint diffuse structures are present (Figs.
1, 4, between $\alpha\simeq$ 02h19m30s, $\delta\simeq$ 62°15′; $\alpha\simeq$
02h24m, $\delta\simeq$ 63°30′ and $\alpha\simeq$ 02h11m10s, $\delta\simeq$
62°23′; $\alpha\simeq$ 02h14m30s, $\delta\simeq$ 63°13′), which all are very
well correlated with the radio emission. The bright H$\alpha+$[N ii] filaments
cover most of the emission found at radio wavelengths. In contrast to the
above results, diffuse emission was mainly detected in the south, centre and
north with only a few small filamentary structures are present (i.e. those at
$\alpha\simeq$ 02h18m40s, $\delta\simeq$ 63°09′10″and $\alpha\simeq$
02h19m15s, $\delta\simeq$ 63°09′20″ ; $\sim$2′ long). We also detected [S ii]
emission where most of the H$\alpha+$[N ii] emission was found with
filamentary bright structures and diffuse emission in the same regions (Figs.
1, 2, 4). Both the H$\alpha+$[N ii] and [S ii] higher–resolution images (Fig.
4 show clearly filamentary and diffuse emission in all slit positions 1–6,
which have similar structures. Furthermore, very thin (less than 2-3″ wide),
as well as wider ($\sim$10″) filaments are present, while the shock front
region can be clearly seen (i.e. in S1 and S4 areas). Note that the detected
[O iii] emission (Figs. 1, 4) appears more filamentary and less diffuse than
in the H$\alpha+$[N ii] and [S ii] images. In Table 2 upper limit fluxes are
listed for the [O iii] $\lambda$5007 line. Significant differences between the
H$\alpha+$[N ii] and [O iii] images are present for many of the filaments. In
particular, most of the bright H$\alpha+$[N ii] and [S ii] complex filamentary
structures to the east, south–east are not present to the [O iii] image, while
there are other bright filaments only to the [O iii] emission. However, some
thin filaments (of the order of 8-10′ long), coincide with the H$\alpha+$[N
ii] and [S ii] emission (see Fig. 1). These are at slit positions 5 and 6 and
to the north at $\alpha\simeq$ 02h22m55s, $\delta\simeq$ 63°25′25″ and
$\alpha\simeq$ 02h22m07s, $\delta\simeq$ 63°30′00″. Further to the east, there
is a number of bright [O iii] filaments which are co-spatial with the
H$\alpha+$[N ii] structures, but still it is clear that the bright [O iii]
filaments do not coincide with those in the H$\alpha+$[N ii] and [S ii]
emission. Interestingly, there are many thin [O iii] filamentary structures to
the central region of HB3 (i.e. at $\alpha\simeq$ 02h20m16s, $\delta\simeq$
62°59′00″, $\alpha\simeq$ 02h19m50s, $\delta\simeq$ 62°53′00″ and
$\alpha\simeq$ 02h14m55s, $\delta\simeq$ 62°49′30″), which do not show any
emission in H$\alpha+$[N ii] and [S ii]. Taking into account all the emission
line images presented in Fig. 1, and the RGB image presented in Fig. 3, it
seems that HB3 traces an axi-symmetric bi-lobal structure with respect to its
centre. This is an interesting approach and should be investigated further in
the future when kinematical data will be available. In Fig. 5, we present an
RGB image of slit positions 1 and 6 regions, where the H$\alpha+$[N ii] (red),
[S ii] (green) and [O iii] (blue) higher–resolution images are combined. In
the S1 image, the H$\alpha+$[N ii] emission of the small thin filaments (i.e.
at $\alpha\simeq$ 02h11m15s, $\delta\simeq$ 62°43′00″) in front of the large
ones (i.e. at $\alpha\simeq$ 02h11m30s, $\delta\simeq$ 62°44′00″) coincide
with that of the [O iii] and lie immediately ahead of the diffuse [O iii]
emission (similar to SNR G 107.0$+$9.0; Fesen et al. 2020). Unfortunately,
there is no [O iii] emission where the bright H$\alpha+$[N ii] filaments are,
so we cannot have a complete and clear view of the whole region.

###  3.2 The optical spectra

Deep low-resolution long-slit spectra were obtained for the brightest
filaments of the remnant. The exact position of the silts is illustrated in
Fig. 2 and the coordinates of their centres are listed in Table 1. The
relative line fluxes as well as the signal-to-noise ratios for multiple
apertures free of background stars are computed and listed in Table 3. The
absolute H$\alpha$ fluxes are also determined for all the apertures and vary
from 5 to 192 $\times$ $10^{-17}$ erg s-1 cm-2 arcsec-2. The lengths of these
apertures and their offsets from the centre of the slits are also given in
Table 1. For the extraction of the background emission, the areas were chosen
north or south of the centre of the slits as being free of field stars and
depending on the location of the filaments. Fig. 6 presents as representative
examples the 1D spectra of six apertures.

Our slit 1 covers the western filament, which has also been studied by
D’Odorico & Sabbadin (1977) and Fesen et al. (1995), and it shows a good
agreement despite the different position and orientation. In particular, we
determine an [S ii]/H$\alpha$ ratio equals to 0.8 and 0.56 for the apertures
1a and 1b, respectively, very close to the values of 0.72 reported by Fesen et
al. (1995) and 0.62 from D’Odorico & Sabbadin (1977). The [S ii]
$\lambda\lambda$ 6716/6731 line ratio is calculated 1.32$\pm$0.06 and
1.39$\pm$0.05, close to the value of 1.5$\pm$0.2 found by Fesen et al. (1995),
and even closer to the value of 1.44$\pm$0.01 from D’Odorico & Sabbadin
(1977). A consistency between the studies is also found in the interstellar
reddening (0.71$\pm$0.04; Fesen et al. 1995 and a median of 0.76 in this
study). The [Oi] $\lambda$6300 line, a diagnostic for shock heated gas in SNRs
(Kopsacheili et al., 2020), is also detected in our spectrum of slit 1.
However, it has to be noted that the [Ni] $\lambda$5200 line which is reported
by Fesen et al. (1995) is not detected in our spectrum, while the [O iii]
$\lambda$5007/H$\beta$ ratio is three times higher in our data. Fig. 5
displays the RGB images of the western (area S1) and eastern (area-S6)
filaments. At the former (S1), there are interestingly two main filaments
bright in H$\alpha+$[N ii] and [O iii] $\lambda$5007 with an offset of a few
arcsec. The second displays some filaments in H$\alpha+$[N ii] emission. This
explains the difference in the [O iii] $\lambda$5007/H$\beta$ ratio between
our result and the one from Fesen et al. (1995) as their slit do not cover the
[O iii] dominated filament resulting in lower [O iii] $\lambda$5007/H$\beta$
ratio.

The common criterion for shock heated gas in SNRs, the [S ii]/H$\alpha$ line
ratio, is higher than 0.4 for all the filaments observed, while the [N
ii]/H$\alpha$ ratio varies from 0.5 to 1 (see Table 3). Both ratios verify the
SNR nature of HB3 (see also Fesen & Gull 1983; Fesen et al. 1995). All the
filaments are distributed in the locus of typical shock-heated SNRs well
separated from the photo-ionized nebulae.

Note that all slits/apertures but two (2a and 2b) exhibit a weak or even not
detectable [O iii] 5007Å line. The log([O iii]/H$\beta$) ratio is lower than
0.5 which indicates shocks with complete recombination lines (Cox & Raymond
1985; Hartigan et al. 1987) and velocities between 80-100 km s-1 (Raymond et
al. 1988; Osterbrock & Ferland 2006). The absence of the [O iii] 5007Å
emission line in some filaments implies the presence of shock with even lower
velocities ($<$80 km s-1), and this result is consistent with the observed [O
i]/H$\alpha$ ratios. On the other hand, high [O iii]/H$\beta$ ratio ($>$0.8)
is found in slit 2 (apertures a and b, see Table 3). Such a strong [O iii]
line indicates shock velocities higher than 100 km s-1. Furthermore, having a
look at Figs. 1 – 3, it can be clearly seen that there are many filaments with
strong [O iii] emission, which are not covered by the available slits, where
even higher shock velocities ($>$120 km s-1) are expected to be found.

The electron density of the filaments is lower than 240 cm-3 obtained using
the [S ii] diagnostic lines and the temden STSDAS/IRAF routine. Due to the
absence of the typical [O iii] $\lambda$4363 and [N ii] $\lambda$5755
diagnostic lines (based on Osterbrock & Ferland 2006), we are not able to
compute the electron temperature of the gas and a value of 104 K is considered
for all the filaments.

Finally, for all filaments where the H$\beta$ line is detected, the
interstellar extinction coefficient c(H$\beta$) is computed and varies from
0.62 ($\pm$ 0.22) up to 1.14 ($\pm$ 0.36) with a mean value of 0.9 or an Av
from 1.31 ($\pm$ 0.48) to 2.32 ($\pm$ 0.76) , as presented in Table 3.

Table 3: Relative line fluxes. The emission line ratios [S ii]/H$\alpha$ , F(6716)/F(6731) and [N ii]/H$\alpha$ are calculated using either the values corrected for interstellar extinction (when H$\beta$ line is available) or the values uncorrected for interstellar extinction.The errors of the emission line ratios, c(H$\beta$) and EB-V are calculated through standard error propagation. | Slit 1a | Slit 1b | Slit 2a | Slit 2b | Slit 2c
---|---|---|---|---|---
Line (Å) | F999Observed fluxes normalized to F(H$\alpha$)=100 and uncorrected for interstellar extinction. | I101010Observed fluxes normalized to F(H$\alpha$)=100 and corrected for interstellar extinction. | S/N111111Numbers represent the signal to noise ratio of the quoted fluxes. | F | I | S/N | F | I | S/N | F | I | S/N | F | I | S/N
H$\beta$ 4861 | 20 | 35 | 6 | 18 | 35 | 9 | 14 | 35 | 3 | 17 | 35 | 3 | $-$ | $-$ | $-$
[O iii] 4959 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | 37 | 88 | 5 | 45 | 87 | 5 | 17 | $-$ | 3
[O iii] 5007 | 6 | 11 | 3 | 3 | 6 | 4 | 127 | 281 | 13 | 152 | 288 | 16 | 63 | $-$ | 4
[N ii] 5755 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | 3 | 5 | 3 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$
[O i] 6300 | 18 | 20 | 9 | 21 | 24 | 14 | 13 | 15 | 4 | 6 | 7 | 4 | 14 | $-$ | 3
[O i] 6364 | $-$ | $-$ | $-$ | 7 | 8 | 6 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$
[N ii] 6548 | 16 | 17 | 10 | 11 | 11 | 10 | 25 | 26 | 7 | 30 | 32 | 5 | 30 | $-$ | 4
H$\alpha$ 6563 | 100 | 100 | 62 | 100 | 100 | 92 | 100 | 100 | 26 | 100 | 100 | 16 | 100 | $-$ | 14
[N ii] 6584 | 49 | 49 | 29 | 33 | 33 | 32 | 79 | 79 | 18 | 85 | 85 | 13 | 86 | $-$ | 13
[S ii] 6716 | 46 | 46 | 29 | 33 | 33 | 32 | 75 | 74 | 18 | 64 | 64 | 13 | 65 | $-$ | 13
[S ii] 6731 | 35 | 35 | 21 | 24 | 23 | 26 | 63 | 62 | 15 | 56 | 55 | 10 | 43 | $-$ | 7
Absolute H$\alpha$ flux121212In units of $10^{-17}$ erg s-1 cm-2 arcsec-2. | 191.5 | 122.8 | 26.2 | 4.9 | 5.1
[S ii]/H$\alpha$ | 0.80 $\pm$ 0.05 | 0.56 $\pm$ 0.03 | 1.34 $\pm$ 0.12 | 1.19 $\pm$ 0.18 | 1.08 $\pm$ 0.20
F(6716)/F(6731) | 1.32 $\pm$ 0.06 | 1.39 $\pm$ 0.05 | 1.21 $\pm$ 0.08 | 1.16 $\pm$ 0.12 | 1.50 $\pm$ 0.21
[N ii]/H$\alpha$ | 0.65 $\pm$ 0.07 | 0.44 $\pm$ 0.05 | 1.07 $\pm$ 0.17 | 1.16 $\pm$ 0.25 | 1.15 $\pm$ 0.31
c(H$\beta$)131313The logarithmic extinction is derived by c = 1/0.348$\times$log((H$\alpha$/H$\beta$)obs/2.85). | 0.69 $\pm$ 0.14 | 0.86 $\pm$ 0.06 | 1.14 $\pm$ 0.41 | 0.88 $\pm$ 0.43 | $-$
EB-V141414The interstellar reddening was measured from the relation EB-V $\approx$ 0.77c (Osterbrock & Ferland 2006) | 0.53 $\pm$ 0.11 | 0.66 $\pm$ 0.05 | 0.88 $\pm$ 0.30 | 0.68 $\pm$ 0.33 | $-$
| Slit 3a | Slit 3b | Slit 4a | Slit 5a | Slit 5b
Line (Å) | F | I | S/N | F | I | S/N | F | I | S/N | F | I | S/N | F | I | S/N
H$\beta$ 4861 | 21 | 35 | 4 | 17 | 35 | 5 | - | - | - | 16 | 35 | 6 | 15 | 35 | 6
[O iii] 4959 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | 5 | $-$ | 2 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$
[O iii] 5007 | 32 | 50 | 9 | 20 | 38 | 6 | 3 | $-$ | 3 | 4 | 8 | 3 | 1 | 2 | 2
[O i] 6300 | 14 | 15 | 6 | 10 | 11 | 5 | 30 | $-$ | 8 | 16 | 19 | 7 | 16 | 18 | 4
[O i] 6364 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | 10 | $-$ | 3 | 5 | 6 | 3 | $-$ | $-$ | $-$
[N ii] 6548 | 20 | 21 | 7 | 16 | 17 | 7 | 18 | $-$ | 5 | 14 | 15 | 7 | 13 | 14 | 6
H$\alpha$ 6563 | 100 | 100 | 29 | 100 | 100 | 31 | 100 | $-$ | 28 | 100 | 100 | 38 | 100 | 100 | 35
[N ii] 6584 | 63 | 63 | 18 | 56 | 56 | 19 | 62 | $-$ | 16 | 49 | 49 | 19 | 41 | 41 | 16
[S ii] 6716 | 70 | 69 | 21 | 59 | 59 | 19 | 89 | $-$ | 23 | 49 | 49 | 21 | 36 | 35 | 14
[S ii] 6731 | 48 | 47 | 15 | 42 | 42 | 13 | 64 | $-$ | 17 | 35 | 34 | 16 | 24 | 24 | 12
Absolute H$\alpha$ flux | 33.8 | 30.3 | 23.0 | 153.6 | 168.2
[S ii]/H$\alpha$ | 1.17 $\pm$ 0.11 | 1.00 $\pm$ 0.09 | 1.53 $\pm$ 0.13 | 0.83 $\pm$ 0.07 | 0.59 $\pm$ 0.07
F(6716)/F(6731) | 1.47 $\pm$ 0.10 | 1.42 $\pm$ 0.09 | 1.38 $\pm$ 0.08 | 1.43 $\pm$ 0.09 | 1.47 $\pm$ 0.12
[N ii]/H$\alpha$ | 0.84 $\pm$ 0.12 | 0.73 $\pm$ 0.11 | 0.80 $\pm$ 0.16 | 0.64 $\pm$ 0.10 | 0.55 $\pm$ 0.10
c(H$\beta$) | 0.62 $\pm$ 0.24 | 0.87 $\pm$ 0.21 | $-$ | 0.98 $\pm$ 0.13 | 1.10 $\pm$ 0.12
EB-V | 0.48 $\pm$ 0.19 | 0.67 $\pm$ 0.17 | $-$ | 0.75 $\pm$ 0.10 | 0.85 $\pm$ 0.09
Table 4: continued | Slit 5c | Slit 5d | Slit 6a | Slit 6b | Slit 6c
---|---|---|---|---|---
Line (Å) | F151515Observed fluxes normalized to F(H$\alpha$)=100 and uncorrected for interstellar extinction. | I161616Observed fluxes normalized to F(H$\alpha$)=100 and corrected for interstellar extinction. | S/N171717Numbers represent the signal to noise ratio of the quoted fluxes. | F | I | S/N | F | I | S/N | F | I | S/N | F | I | S/N
H$\beta$ 4861 | 17 | 35 | 6 | 19 | 35 | 6 | 16 | 35 | 5 | $-$ | $-$ | $-$ | 18 | 35 | 5
[O iii] 4959 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | 7 | 15 | 3 | 11 | $-$ | 3 | $-$ | $-$ | $-$
[O iii] 5007 | $-$ | $-$ | $-$ | 7 | 12 | 3 | 18 | 36 | 4 | 37 | $-$ | 6 | 14 | 25 | 5
[N ii] 5755 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$
[O i] 6300 | 15 | 17 | 7 | 7 | 7 | 7 | 21 | 23 | 7 | 25 | $-$ | 7 | 14 | 15 | 7
[O i] 6364 | 5 | 5 | 3 | 2 | 2 | 3 | 7 | 8 | 3 | 8 | $-$ | 2 | 5 | 5 | 3
[N ii] 6548 | 15 | 16 | 7 | 12 | 12 | 5 | 15 | 15 | 6 | 28 | $-$ | 7 | 11 | 11 | 5
H$\alpha$ 6563 | 100 | 100 | 39 | 100 | 100 | 35 | 100 | 100 | 28 | 100 | $-$ | 27 | 100 | 100 | 33
[N ii] 6584 | 50 | 50 | 20 | 41 | 41 | 14 | 45 | 45 | 16 | 65 | $-$ | 19 | 38 | 38 | 14
[S ii] 6716 | 48 | 48 | 21 | 37 | 37 | 13 | 54 | 53 | 17 | 58 | $-$ | 17 | 38 | 38 | 16
[S ii] 6731 | 33 | 33 | 15 | 25 | 25 | 10 | 37 | 37 | 15 | 36 | $-$ | 13 | 31 | 31 | 14
Absolute H$\alpha$ flux181818In units of $10^{-17}$ erg s-1 cm-2 arcsec-2. | 176.8 | 91.1 | 32.4 | 26.5 | 25.8
[S ii]/H$\alpha$ | 0.81 $\pm$ 0.07 | 0.62 $\pm$ 0.08 | 0.88 $\pm$ 0.08 | 0.94 $\pm$ 0.10 | 0.68 $\pm$ 0.07
F(6716)/F(6731) | 1.44 $\pm$ 0.10 | 1.47 $\pm$ 0.15 | 1.46 $\pm$ 0.10 | 1.58 $\pm$ 0.12 | 1.25 $\pm$ 0.09
[N ii]/H$\alpha$ | 0.65 $\pm$ 0.10 | 0.53 $\pm$ 0.11 | 0.60 $\pm$ 0.11 | 0.93 $\pm$ 0.14 | 0.49 $\pm$ 0.10
c(H$\beta$)191919The logarithmic extinction is derived by c = 1/0.348$\times$log((H$\alpha$/H$\beta$)obs/2.85). | 0.89 $\pm$ 0.16 | 0.78 $\pm$ 0.13 | 0.98 $\pm$ 0.23 | $-$ | 0.83 $\pm$ 0.13
EB-V202020The interstellar reddening was measured from the relation EB-V $\approx$ 0.77c (Osterbrock & Ferland 2006) | 0.68 $\pm$ 0.12 | 0.60 $\pm$ 0.10 | 0.75 $\pm$ 0.18 | $-$ | 0.64 $\pm$ 0.09
| Slit 6d | Slit 7a | Slit 7b | Slit 8a | Slit 8b
Line (Å) | F | I | S/N | F | I | S/N | F | I | S/N | F | I | S/N | F | I | S/N
H$\beta$ 4861 | 15 | 35 | 4 | 15 | 35 | 4 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$
[O iii] 5007 | 33 | 73 | 8 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$
[O i] 6300 | 13 | 15 | 3 | 11 | 12 | 5 | 14 | $-$ | 5 | 21 | $-$ | 4 | 6 | $-$ | 4
[O i] 6364 | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | $-$ | 12 | $-$ | 2 | $-$ | $-$ | $-$
[N ii] 6548 | 16 | 17 | 6 | 19 | 20 | 8 | 41 | $-$ | 5 | $-$ | $-$ | $-$ | 19 | $-$ | 3
H$\alpha$ 6563 | 100 | 100 | 27 | 100 | 100 | 35 | 100 | $-$ | 19 | 100 | $-$ | 19 | 100 | $-$ | 15
[N ii] 6584 | 48 | 48 | 16 | 58 | 57 | 22 | 60 | $-$ | 14 | 32 | $-$ | 4 | 45 | $-$ | 7
[S ii] 6716 | 37 | 37 | 13 | 41 | 41 | 18 | 54 | $-$ | 13 | 44 | $-$ | 12 | 32 | $-$ | 6
[S ii] 6731 | 31 | 30 | 11 | 26 | 26 | 15 | 32 | $-$ | 10 | 32 | $-$ | 9 | 24 | $-$ | 4
Absolute H$\alpha$ flux | 26.1 | 85.9 | 4.1 | 24.3 | 8.8
[S ii]/H$\alpha$ | 0.67 $\pm$ 0.08 | 0.67 $\pm$ 0.07 | 0.86 $\pm$ 0.12 | 0.76 $\pm$ 0.10 | 0.56 $\pm$ 0.16
F(6716)/F(6731) | 1.22 $\pm$ 0.11 | 1.58 $\pm$ 0.11 | 1.68 $\pm$ 0.17 | 1.35 $\pm$ 0.15 | 1.35 $\pm$ 0.31
[N ii]/H$\alpha$ | 0.65 $\pm$ 0.12 | 0.78 $\pm$ 0.11 | 1.01 $\pm$ 0.22 | 0.32 $\pm$ 0.08 | 0.64 $\pm$ 0.22
c(H$\beta$) | 1.10 $\pm$ 0.22 | 1.10 $\pm$ 0.25 | $-$ | $-$ | $-$
EB-V | 0.85 $\pm$ 0.17 | 0.85 $\pm$ 0.19 | $-$ | $-$ | $-$
Figure 8: 2D diagnostic curves of Kopsacheili et al. (2020) together with the
results extracted by our spectroscopic data (blue points). The blue region is
the locus that indicates shock excited gas emission.

###  3.3 Comparison with other wavelengths

G132.7+01.3 has been studied in $\gamma$-rays, X-rays, radio, infrared and mm
wavelengths (e.g. Routledge et al., 1991; Fesen et al., 1995; Lazendic &
Slane, 2006; Zhou et al., 2016; Rho et al., 2021, see also Introduction). Here
we present the main results and conclusions regarding the properties and
evolution of HB3 as extracted by the comparison of its emission properties in
different electromagnetic bands. In Fig. 7 the contours of the emission that
the remnant reveals in the radio, X-rays, CO lines and $\gamma$-rays are
superposed on the H$\alpha+$[N ii] image of HB3 from our observations.

The optical morphology of the remnant displays a strong correlation with its
radio emission at both large and small spatial scales, signifying that they
have the same origin (Fig. 7a). In particular, the outer bright H$\alpha+$[N
ii] filaments at the western portion of the remnant follow both the geometry
and brightness of the radio emission that declines gradually, moving from the
west to the northwest region of HB3. This advocates that the remnant
encounters a denser ambient medium in the west part than in the northwest.
Intriguingly, both in radio and optical the remnant reveals two antisymmetric
local protrusions in west and east giving the morphological impression of two
lobes in the overall shape of the remnant. In the southern portion of HB3 even
though the optical and radio display a similar morphology, the two bands are
not in good correlation . In particular, the optical emission in the south of
HB3 displays a relatively faint and diffuse emission, while the radio emission
in this portion is equally bright as compared to the west region of HB3. This
could be attributed either to dust absorption affecting the optical photons
and/or to a high local blast wave velocity that enhances the non-thermal
synchrotron emission at this region. Finally, the eastern region of HB3 the
radio emission of the remnant roughly follows the bright filament in
H$\alpha+$[N ii] marking the border between the SNR and the ambient Hii/MC
complex.

Regarding the inner, thermal X-ray emission of the remnant (Fig. 7b) -a
property that classifies HB3 as a MMSNR (Rho & Petre, 1998; Lazendic & Slane,
2006)\- it is lying close to the center of the optical image of HB3 being
centralized at about
$\left(\alpha,\delta\right)=\left(2^{h}~{}18^{m},62^{o}~{}40^{\prime}\right)$.
The remnant’s X-ray brightness declines moving from the center to the outer
portions of HB3, following initially a roughly spherical distribution but at
the same time it displays two large anti-asymmetric wings towards the
northeast and south that cover the whole optical image of the remnant in these
regions. The overall X-ray morphology of HB3 is shifted towards the eastern
region of the remnant where it intersects with the Hii/MC complex, something
that advocates that the mixed morphology properties of HB3 are related to the
interaction of the remnant with the surrounding cloud. Such a hypothesis is
aligned to evaporating clouds model (White & Long, 1991) and the reflected
shock model (Chen et al., 2008) for the formation of MMSNRs as both demand the
interaction of the SNR with a surrounding dense cloud (see also Sect. 4.3).

The molecular gas in the vicinity of HB3 as detected by millimeter
observations of CO lines (Zhou et al., 2016), coincides with the bright
H$\alpha+$[N ii] and [S ii] southeast region of the SNR, lying about
$\left(\alpha,\delta\right)=\left(2^{h}~{}18^{m}~{}30^{s},62^{o}~{}20^{\prime}\right)$.
Given that the optical emissivity of the shocked gas goes as $\rho^{2}$ where
$\rho$ its mass density, it is inferred that in this region the SNR blast wave
is propagating in a dense environment. This result offers an additional
evidence towards the conclusion that HB3 is currently interacting with the W3
Hii region/MC complex (see also Zhou et al., 2016; Routledge et al., 1991, for
an additional argumentation).

Finally, the $\gamma$-ray emission discovered in the vicinity of HB3
(Tsygankov et al., 2016; Katagiri et al., 2016, ; see also Fig. 7d) is
spatially correlated with the bright ${}^{12}CO$ (J = 1-0) emission and it is
adjacent to the southeast region of the remnant. Such a result favors for a
hadronic origin of the $\gamma$-ray photons resulted by interactions between
particles accelerated in the SNR and the molecular cloud.

All above evidence advocate that HB3 is indeed expanding in the vicinity of
the W3 Hii region/MC complex and currently is interacting with it at its
eastern portion. This interaction seems to substantially affect the morphology
of the remnant as well as its emission properties all over the electromagnetic
spectrum. The borders between the SNR and the surrounding cloud are well
justified by the radio observations and the bright H$\alpha+$[N ii] filaments
at the eastern region of the remnant.

##  4 Discussion

###  4.1 Properties of the optical filaments

Having an angular diameter of $\sim$90′$\times$120′, HB3 is among the largest
Galactic SNRs known. Furthermore, assuming a distance of $\sim$2 kpc, this
means a linear diameter of 52 $\times$72 pc, indicative of an evolved SNR
being in the later Sedov–Taylor or in the radiative phase. The composite
H$\alpha+$[N ii] image provides the best description of the morphology of
G132.7+01.3 in its full extent making possible the identification of all of
its detailed structures. Emission from the [S ii] $\lambda\lambda$ 6716,6731
doublet lines and the [O iii] $\lambda$5007 line has also been detected
matching in shape and position the H$\alpha+$[N ii] emission, though being
more diffuse and less filamentary. Furthermore, the prominent [O iii] emission
indicates the existence of fast shocks travelling through the ionized gas all
around the remnant. (Cox & Raymond 1985).

The [S ii]/H$\alpha$ emission line ratio measured from the optical images of
the Area A (Table 2) does not exceed the threshold (0.4) of supernova
remnants. We thus argue that this particular area does not belong to
G132.7+01.3, but it is associated with the Hii region. Note that the area A is
also the brightest in the H$\alpha$ emission line (see Table 2). All other
areas display [S ii]/H$\alpha>$ 0.4, except area S7 for which a ratio of 0.35
is determined. Note that the width of the H$\alpha$+[N ii] filter used for the
computation of H$\alpha$ emission is large enough to transmit the [N ii] 6584
Å line. This extra emission is likely responsible for the lower [S
ii]/H$\alpha$ ratio measured for the area S7.

Regarding our spectroscopic observations of eight filaments distributed over
the entire remnant, the [S ii]/H$\alpha$ ratios are clearly higher than 0.4.
The distribution of our spectra in the H$\alpha$/[S ii] –H$\alpha$/[Nii] space
is consistent with the locus of shock-heated gas in SNRs and well separated
from the UV-dominated PNe and Hii regions (e.g. Sabbadin et al., 1977; Meaburn
et al., 2010; Leonidaki et al., 2013; Sabin et al., 2013; Akras et al., 2020).
The strong [O i] $\lambda$6300 Å emission in most of the filaments provides an
extra confirmation of their shock-heated origin.

In Fig. 8, we present the newly developed diagnostic diagrams between shock-
heated and UV photo-ionized regions for different combinations of emission
line ratios (Kopsacheili et al., 2020). Our observed [S ii] /H$\alpha$,
[Nii]/H$\alpha$ and [O i]/H$\alpha$ ratios from all eight filaments and
different apertures satisfy the criteria and lie within the regime of expected
shock-heated regions.

However, a small discrepancy is observed when the [O iii]/H$\beta$ criterion
is used for slits 1 and 5. The measured intensities of the emission line(s) of
these slits are highly uncertain, with signal-to-noise values of 2-4 (see
Table 3). This is a possible interpretation why regions corresponding to Slit
1 and 5 fall outside the expected locus of shock-heated regions in the [O
iii]/H$\beta$ diagnostic diagrams. This can also be seen from the [O iii]
5007Å emission which is very weak or even absent at the position of the slits
1 and 5 (see RGB Figure 3 of HB3). A possible contribution in the H$\beta$
emission from the nearby Hii region W3 in slit 5, that would result in lower
[O iii]/H$\beta$ ratio, cannot be ruled out. We point out however that the new
diagnostics do not account for shocks with velocities lower than 100 km s-1
(only models with shock velocities between 100 and 1000 km s-1 are considered,
Kopsacheili et al. 2020). Furthermore, we note that all the diagnostics are
characterized by incompleteness, which is higher in the case of [N
ii]/H$\alpha$–[O iii]/H$\beta$ (for more details see, Kopsacheili et al.,
2020).

###  4.2 Origin and evolution of HB3

Adopted from the optical image extent of HB3 ($\sim 90^{\prime}\times
120^{\prime}$) and a distance of $d=$ 2 kpc, the average radius of the remnant
is $R_{s}=$ 31 pc. Assuming further that the remnant is still in the Sedov
phase and the magnetic field pressure is negligible, the values of the
preshocked ambient density ($n_{o}$) and the current shock’s velocity
($V_{s}$) can be constrained by using the relation (Dopita, 1979):

${\rm n_{[S{\sc II}]}\simeq\ 45\ \left(\frac{n_{o}}{\rm
cm^{-3}}\right)\left(\frac{V_{\rm s}}{100~{}km~{}s^{-1}}\right)^{2}}~{}{\rm
cm^{-3}},$ (1)

where $n_{\rm[SII]}=160\pm 80~{}\rm cm^{-3}$, extracted from the [S ii]
diagnostic lines. From Eq. (1), we get: $\left(\frac{n_{o}}{\rm
cm^{-3}}\right)\left(\frac{V_{\rm s}}{100~{}km~{}s^{-1}}\right)^{2}=3.6\pm
1.8$ . Such a result is in agreement with the independent estimations of HB3
shock velocity and pre-shock density by Lazendic & Slane (2006) who found
$\left(\frac{n_{o}}{\rm cm^{-3}}\right)\left(\frac{V_{\rm
s}}{100~{}km~{}s^{-1}}\right)^{2}=\left(\frac{0.32}{\rm
cm^{-3}}\right)\left(\frac{340}{100~{}km~{}s^{-1}}\right)^{2}=3.7\pm 1.4$.
Additional constraints on the ambient medium density -and thus on the current
shock velocity- we get by estimating the HB3 column density. Using the
relation of Ryter et al. (1975), we obtain an NH of $3.3(\pm 1.5)\times
10^{21}~{}{\rm cm}^{-2}$ and $6.0(\pm 1.7)\times 10^{21}~{}{\rm cm}^{-2}$ for
the minimum and maximum c(H$\beta$) values as calculated from our spectra,
respectively. The extracted column density range is consistent to the
estimated values of N${}_{\rm H}\sim 6.2\times 10^{21}$ cm-2 (Dickey &
Lockman, 1990) and N${}_{\rm H}\sim 7.7\times 10^{21}$ cm-2 (Kalberla et al.,
2005) in the direction of HB3, and the one extracted by Fesen et al. (1995)
using their optical spectra in HB3 western rim, who found N${}_{\rm H}\sim
5.7\times 10^{21}~{}{\rm cm}^{-2}$. The corresponding pre-shocked ambient
medium number density range is $0.5-1.0~{}d_{2}^{-1}\rm cm^{-3}$ having a mean
value of $n_{o}=\left(0.82\pm 0.15\right)~{}d_{2}^{-1}$ $\rm cm^{-3}$, where
$d_{2}$ the distance of HB3 in units of 2 kpc. Using Eq. (1) we extract a mean
shock wave velocity of $V_{s}=\left(208\pm 54\right)~{}d_{2}^{0.5}\rm
km~{}s^{-1}$. Employing the aforementioned results extracted by our optical
observations on the Sedov solution, we find an explosion energy for HB3 equal
to
$E~{}\sim~{}2.1~{}\times~{}10^{46}~{}\left(\frac{R_{s}}{pc}\right)^{3}~{}\left(\frac{n_{o}}{cm^{-3}}\right)~{}\left(\frac{V_{\rm
s}}{100~{}km~{}s^{-1}}\right)^{2}=\left(2.3\pm
1.1\right)~{}\times~{}10^{51}d_{2}^{3}~{}\rm erg$ (Ostriker & McKee, 1988)
while the SNR age is estimated to be $t\sim\left(5.8\pm 1.5\right)\times
10^{4}~{}d_{2}^{0.5}$ yrs.

The large size of the remnant and its bright optical line emission indicate
that HB3 has evolved beyond the SNR adiabatic phases and currently is in the
Pressure Driven Snowplow (PDS) phase. In order to assess the correctness of
this assumption, we estimate the transition radii from the Sedov to the PDS
phase extracted by Cioffi et al. (1988):

${\rm
R_{PDS}=14.0\left(\frac{E}{10^{51}~{}erg}\right)^{\frac{2}{7}}\left(\frac{n_{o}}{cm^{-3}}\right)^{-\frac{3}{7}}\zeta_{m}^{-\frac{1}{7}}},$
(2)

where $E$ is the SN energy and $\zeta_{m}$ a constant equal to unity for solar
metalicity. In the PDS phase the SN energy is found to be $E=6.8\times
10^{43}\left(\frac{n_{o}}{cm^{-3}}\right)^{1.16}\left(\frac{V_{s}}{km~{}s^{-1}}\right)^{1.35}\left(\frac{R_{s}}{pc}\right)^{3.16}\zeta_{m}^{0.161}=\left(3.7\pm
1.5\right)\times 10^{51}d_{2}^{2.7}$ erg (Cioffi et al., 1988) for which the
corresponding transition radius is $R_{\rm PDS}=\left(22.2\pm
3.1\right)d_{2}^{1.2}$ pc. The transition radius is slightly smaller that the
current radius of HB3 ($Rs\sim 31d_{2}$ pc) something that indicates that the
remnant has evolved into the PDS phase or is currently in the transition phase
between the Sedov and PDS stages. Intriguingly, this finding is commonly met
in the class of MMSNRs (Shelton et al., 1999). The corresponding time in which
HB3 started to enter into the PDS phase is $t_{\rm PDS}=1.33\times
10^{4}\left(\frac{E}{10^{51}~{}erg}\right)^{\frac{3}{14}}\left(\frac{n_{0}}{cm^{-3}}\right)^{-\frac{4}{7}}\zeta_{m}^{-\frac{5}{14}}=\left(2.0\pm
0.3\right)\times 10^{4}~{}d_{2}^{1.15}$ yrs, while the current age of the
remnant was found to be $t=~{}\frac{3}{4}t_{\rm
PDS}\left[\left(\frac{R_{s}}{R_{\rm
PDS}}\right)^{\frac{10}{3}}+\frac{1}{3}\right]=\left(5.1~{}\pm
2.1\right)~{}\times~{}10^{4}$ yrs. These timescales indicate that HB3 has
spent at least the 30% of its lifetime in the PDS phase.

###  4.3 On the mixed morphology of HB3

The previous approach provides a first order approximation on the current
evolutionary properties of the remnant as well as on its parent SN explosion.
Nevertheless, it neglects the following two major facts: a) the ambient medium
properties around HB3 deviate substantially from the homogeneity revealing
large spatial gradients. In particular, our optical study confirmed that HB3
interacts in its eastern region with a dense Hii/MC complex something that
affects its morphology, kinematics and emission properties. Thus, the
evolution of the remnant cannot be accurately approached by one-dimensional,
spherically symmetric models; b) HB3 has been classified as a mixed morphology
SNR being characterized by the coexistence of an optical/radio bright shell
with a centrally peaked thermal X-ray emission. These SNR properties cannot be
explained by the canonical SNR evolution models and therefore, additional
‘ingredients’ on its evolutionary history are required.

Lazendic & Slane (2006) provided an explanation on the mixed morphology
properties of HB3 by invoking the so called ‘evaporating clouds model’ (e.g.
White & Long, 1991; Slavin et al., 2017; Zhang et al., 2019) and adopting a
$C\over\tau$ ratio of 3-5, where $C$ is the mass ratio between the clouds and
the intercloud medium and $\tau$ is the ratio of the cloud evaporation
timescale to the current SNR age.

Here we discuss an alternative evolutionary scenario of HB3 that potentially
could led on its mixed morphology properties based on the ‘reflected shock
model’. According to this model MMSNRs are formed by the action of a reflected
shock, triggered by the collision of the SNR with the density walls of a pre-
existing circumstellar cavity sculptured by the mass outflows of the
progenitor star (Dwarkadas, 2005; Chen et al., 2008; Zhang et al., 2015).

If this was the case of HB3, the size of the formed wind-blown cavity should
be comparable to the current size of the remnant. This fact provides us an
estimation of the progenitor mass by employing the linear relation of Chen et
al. (2013):

${\rm R_{b}\simeq\ \left(1.22\pm
0.05\right)\times\left(\frac{M_{ms}}{M_{\odot}}\right)-\left(9.16\pm
1.77\right)~{}pc}.$ (3)

where $R_{b}$ the radius of the wind blown cavity and $M_{ms}$ the stellar
main sequence mass. Considering that the current radius of HB3 is $\sim
31~{}d_{2}$ pc we estimate that the remnant resulted by a progenitor star with
a main sequence mass of about $M_{ms}=33.9\pm 2.0~{}\rm M_{\odot}$. According
to the stellar evolution theory this star most likely passed through the Wolf-
Rayet phase before its final explosion. If true, then HB3 resulted by a Type
Ib/c SN (Smith, 2014, and references therein).

According to this scenario, the remnant initially evolved into the wind blown
cavity of a size $R_{b}\sim 31$ pc, shaped by the progenitor star. Describing
the SN ejecta density profile with a power law of $\rho_{ej}\propto r^{-n}$,
with n=7 (appropriate for Type I SNe; Chevalier, 1981), and assuming energy
conservation, the evolution of the SNR within the wind cavity till its
collision with the density walls of the circumstellar structure can be
approached by employing the self-similar solution of Chevalier (1982). The
radius ($R_{s}$) and velocity ($V_{s}$) of a SNR expanding in an wind blown
cavity with a density profile $\rho_{AM}=q~{}r^{-s}$, where s=2 and
$q=\frac{\bm{\dot{}}{M_{w}}}{4\pi u_{w}}$ and are given by:

${\rm
R_{s}=1.3\times\left(\frac{A~{}g^{n}}{q}\right)^{\frac{1}{n-s}}t^{\frac{n-3}{n-s}}}.$
(4)

${\rm
V_{s}=1.3\times\left(\frac{n-3}{n-s}\right)\times\left(\frac{A~{}g^{n}}{q}\right)^{\frac{1}{n-s}}t^{\frac{s-3}{n-s}}}.$
(5)

where for n= 7 we have A= 0.27 and
$g=\left[\left(\frac{25}{21\pi}\right)\left(\frac{E_{ej}^{2}}{M_{ej}}\right)\right]^{\frac{1}{7}}$
and t is the age of the remnant. By assuming the canonical energy for the SN
explosion of $E_{ej}=10^{51}$ erg, an ejecta mass of
$M_{ej}=9~{}\rm~{}M_{\odot}$ (as expected to be the final mass of a
$34~{}\rm~{}M_{\odot}$ progenitor e.g. Meyer et al., 2021) and adopting
typical WR wind properties $\left({\bm{\dot{}}{M}}\over
10^{-5}~{}\rm~{}M_{\odot}~{}yr^{-1}\right)\times\left(\dot{u_{w}}\over
10^{3}~{}\rm~{}km~{}s^{-1}\right)^{-1}=1$ we find that the remnant reached
with the density walls of the wind blown cavity after $t\approx 4500$ yrs of
evolution having an incident shock velocity of $Vs\approx
5400~{}\rm~{}km~{}s^{-1}$.

At the moment of the collision between the SNR and the density wall a pair of
shocks is formed: the transmitted shock that starts to penetrate the density
wall and the reflected shock that moves inwards and shocks the SN ejecta. Due
to the high wall-to-cavity density contrast the transmitted shock is expected
to decelerate substantially (Dwarkadas, 2005, 2007; Chiotellis et al., 2012,
2013). In particular, the transmitted shock velocity ($v_{t}$) right after the
collision of the SNR with the cavity walls as a function of the incident shock
velocity ($Vs$) is given by (Sgro, 1975):

${\rm V_{t}=\left(\frac{\beta}{A}\right)^{\frac{1}{2}}\times V_{s}}.$ (6)

where $A$ is the wall to cavity density contrast and $\beta$ the pressure
ratio of the post-transited-shock gas to the the post-incident-shock-gas. The
ratio $\beta$ is related to the density contrast between the post-reflected-
shock gas and post-incident shock gas ($A_{r}$) as:

${\rm\beta=\frac{4A_{r}-1}{4-A_{r}}},$ (7)

and $A_{r}$ in turn is related to $A$ as:

${\rm
A=\frac{3A_{r}\left(4A_{r}-1\right)}{\left[\left[3A_{r}\left(4-A_{r}\right)\right]^{\frac{1}{2}}-\sqrt{5}\left(A_{r}-1\right)\right]^{2}}}.$
(8)

For the adopted wind parameters the number density at the outer region of the
wind cavity will be in the order of $n_{cav.}\sim~{}10^{-3}~{}\rm cm^{-3}$.
The cavity wall density cannot be described by a single value as it is
strongly related to the ambient medium density, that as discussed above
deviates substantially from homogeneity. Adopting an ambient medium of
$n_{o}=0.5-1~{}\rm cm^{-3}$, as extracted by the column density estimations,
for the western portion of the remnant, and a $n_{cloud}\sim
10^{1}-10^{2}~{}\rm cm^{-3}$ for the eastern cloud and considering that the
ambient medium has been compressed by at least a factor of 4, we get $A\sim
10^{3}-10^{4}$ and $A\sim 10^{4}-10^{5}$ for the western and eastern regions
of the remnant, respectively. Importing these values to Eq. (5) - (8) we get a
transmitted shock velocity at west of $V_{t}\approx 130-400~{}\rm km~{}s^{-1}$
(consistent to the estimations of Lazendic & Slane, 2006) while to the eastern
region a much slower one in the order of $V_{t}\approx 40-130~{}\rm
km~{}s^{-1}$. Regarding the reflected shock, its initial velocity is given:

${\rm
V_{r}=\frac{1}{4}\left[3-\left(\frac{15A_{r}}{4-A_{r}}\right)^{\frac{1}{2}}\right]V_{s}},$
(9)

from which we get $V_{r}\approx 2300-2500~{}\rm km~{}s^{-1}$ at the east and
$V_{r}\approx 2500-2600~{}\rm km~{}s^{-1}$ at the west of HB3. A shock wave
with such a high velocity will compress and heat the gas to temperatures that
becomes X-ray bright.

Such a scenario seems consistent with the observed properties of HB3, since in
principle can explain the low expansion velocities of HB3 forward shock and at
the shame time its central X-ray emission. As the remnant evolved further the
transmitted shock penetrated more the density wall and it got further
decelerated up to its current expansion velocities. Simultaneously, the
reflected shock reached the center of the remnant producing the overall
centrally X-ray peak morphology of HB3. In order to verify whether such a
scenario could indeed reproduce the current properties of HB3 detailed
hydrodynamic modeling is required.

##  5 Conclusions

  1. 1.

HB3 (G 132.7$+$1.3), which is among the largest supernova remnants, was
studied in optical wavelengths in its full extent for the first time.
Filamentary and diffuse structures are both present in its optical images. The
[S ii] emission of HB3, though fainter, displays a similar morphology to that
of the H$\alpha+$[N ii] emission-line image, tracing all the filamentary
structures. On the other hand, the [O iii] emission line, displays a different
filamentary morphology compared to the former line images, indicating
different shock velocities and physical conditions.

  2. 2.

Spectroscopic data from eight filaments distributed across the whole remnant
exhibit strong H$\alpha$, [N ii] $\lambda\lambda$6548,6584, [S ii]
$\lambda\lambda$6716,6731 and [O i] $\lambda\lambda$6300,6363 emission lines
and low-to-moderate emission from the [O iii] $\lambda\lambda$4959,5007 lines.
These findings indicate shock waves with velocities $<$100 km s-1 for most of
the filaments with an exception of slit 2, which has higher [O iii]/H$\beta$
ratio and shock velocity up to 120 km s-1. The [S ii] lines ratio implies an
electron density $<$240 cm-3.

  3. 3.

The observed emission line ratios are also found to be in agreement with the
diagnostic diagrams that distinguish shock-heated gas (SNRs) and UV-
photoionized gas (Hii regions) verifying the origin of the optical emission
from shock-heated gas.

  4. 4.

We compared the H$\alpha$+[N ii] image extracted by our observations with the
emission that HB3 displays in the radio, X-rays, millimeter and $\gamma$-rays
bands. Overall, the radio emission of HB3 is co-spatial with the bright
optical filaments of the remnant something that indicates the same physical
origin of the two radiations and marks the border of the SNR in respect to the
adjacent cloud. The X-ray emission peaks roughly close to the optical center
of the remnant but its overall distribution tends to be shifted towards the
east, advocating that there is a link between the interaction of the remnant
with the Hii/MC complex and the mixed morphology properties of HB3. Finally,
regarding the CO emission line and $\gamma$-ray emission there are co-spatial
with the bright southeast optical filament of the remnant something that
offers an additional evidence on the interaction of HB3 with the W3 complex.

  5. 5.

Based on the [S ii] diagnostic lines and column densities obtained by our
observations, we estimated the current shock velocity and preshock ambient
medium density of HB3 and we found to be $V_{s}=208\pm 54~{}\rm km~{}s^{-1}$
and $n_{o}=0.82\pm 0.15~{}\rm cm^{-3}$, respectively. Regarding the
evolutionary state of HB3, we show that the remnant has most likely passed
into the PDS phase or it is in transition between the Sedov and PDS stages.
Under this assumption, the SN energy of HB3 was found to be $\left(3.7\pm
1.5\right)\times 10^{51}$ erg and its current age $\left(5.2\pm
2.1\right)\times 10^{4}$ yrs.

  6. 6.

We discussed the mixed morphology properties of HB3, under the reflected shock
model. We show that the overall properties of the remnant, namely its
optical/radio bright shocks and its central X-ray emission, can in principle
be explained considering that the remnant is currently interacting with the
density walls of a wind blown cavity shaped by a progenitor star with initial
mass of $\sim 32-36~{}\rm M_{\odot}$.

## Acknowledgements

This research is co-financed by Greece and the European Union (European Social
Fund-ESF) through the Operational Programme “Human Resources Development,
Education and Lifelong Learning 2014-2020” in the context of the project “On
the interaction of Type Ia Supernovae with Planetary Nebulae” (MIS 5049922).
A.C. acknowledge the support of this work by the project “PROTEAS II” (MIS
5002515),which is implemented under the Action “Reinforcement of the Research
and Innovation Infrastructure”, funded by the Operational Programme
”Competitiveness, Entrepreneur- ship and Innovation” (NSRF 2014–2020) and co-
financed by Greece and the European Union (European Regional Development
Fund). We thank (a) Pat Slane for providing us with the ROSAT PSPC X-ray image
in fits format, (b) Xin Zhou for providing us with the CO images in fits
format and (c) Hideaki Katagiri for providing us with the $\gamma$-ray image
in fits format. This work is based on observations made with the “Aristarchos”
telescope operated on the Helmos Observatory by the Institute of Astronomy,
Astrophysics, Space Applications and Remote Sensing of the National
Observatory of Athens. Skinakas Observatory is a collaborative project of the
University of Crete, the Foundation for Research and Technology-Hellas and the
Max-Planck-Institut für Extraterrestrische Physik. Observatory is a
collaborative project of the University of Crete, the Foundation for Research
and Technology-Hellas and the Max-Planck-Institut für Extraterrestrische
Physik. This research made use of data from SuperCOSMOS H$\alpha$ Survey
(AAO/UKST), from the ATNF Pulsar Catalogue and from the NASA/IPAC Infrared
Science Archive.

## Data Availability

The data underlying this article will be shared on reasonable request to the
corresponding author.

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