# First optical identification of the SRG/eROSITA-detected supernova remnant G
116.6$-$26.1 I. Preliminary results.

E.V. Palaiologou,1,2 I. Leonidaki,${}^{1,2\href
https://orcid.org/0000-0001-5547-3401}$ M. Kopsacheili${}^{1,2\href
https://orcid.org/0000-0002-3563-819X}$
1University of Crete, Department of Physics, Voutes University Campus,
GR-70013 Heraklion, Greece
2Foundation for Research and Technology – Hellas, P.O. Box 1385, GR-71110
Heraklion, Greece e-mail<EMAIL_ADDRESS>
(Accepted 2022 June 02. Received 2022 June 01; in original form 2022 April 20)

###### Abstract

The supernova remnant (SNR) candidate G 116.6$-$26.1 is one of the few high
Galactic latitude (|b| > 15∘) remnants detected so far in several wavebands.
It was discovered recently in the SRG/eROSITA all-sky X-ray survey and
displays also a low-frequency weak radio signature. In this study, we report
the first optical detection of G 116.6$-$26.1 through deep, wide-field and
higher resolution narrowband imaging in H$\upalpha$ , $[\text{S\,{ii}}]$ and
$[\text{O\,{iii}}]$ light. The object exhibits two major and distinct
filamentary emission structures in a partial shell-like formation. The optical
filaments are found in excellent positional match with available X-ray, radio
and UV maps, can be traced over a relatively long angular distance
($38\arcmin$ and $70\arcmin$) and appear unaffected by any strong interactions
with the ambient interstellar medium. We also present a flux-calibrated,
optical emission spectrum from a single location, with Balmer and several
forbidden lines detected, indicative of emission from shock excitation in a
typical evolved SNR. Confirmation of the most likely SNR nature of G
116.6$-$26.1 is provided from the observed value of the line ratio [S ii] /
H$\upalpha$ = $0.56\,\pm\,0.06$, which exceeds the widely accepted threshold
0.4, and is further strengthened by the positive outcome of several diagnostic
tests for shock emission. Our results indicate an approximate shock velocity
range 70–100 km s-1 at the spectroscopically examined filament, which, when
combined with the low emissivity in H$\upalpha$ and other emission lines,
suggest that G 116.6$-$26.1 is a SNR at a mature evolutionary stage.

###### keywords:

ISM: supernova remnants – ISM: individual objects: G 116.6$-$26.1 .

††pubyear: 2022††pagerange: First optical identification of the SRG/eROSITA-
detected supernova remnant G 116.6$-$26.1 I. Preliminary results.–B

##  1 Introduction

G 116.6$-$26.1 is a newly discovered Galactic supernova remnant (SNR) through
soft X-ray imaging and spectra obtained with the Russian-German observatory
SRG/eROSITA (Churazov et al., 2021). According to the X-ray observations, it
is a faint and nearly circular object, with large angular size (about 4∘ in
diameter), located in high Galactic latitude. The morphological structure of
the SNR candidate indicates a marginal brightening along the periphery (mostly
in the southern and western portions) and signs of mild surface brightness
enhancement at the innermost 20′of the object. Its soft X-ray imaging spectrum
is dominated by O vii and O viii lines, similar to the surrounding background
spectrum. Churazov et al. (2021) suggest that the supernova (SN) explosion
took place in the halo of Milky Way, characterized by hot [$\sim$(1–2)
$\times\,10^{6}\,K$] and low density (about $10^{-3}~{}cm^{-3}$) plasma.

They propose a relatively old remnant, originating from a thermonuclear
explosion (SN Ia) that happened about 40 000 years ago, at an approximate
distance of 3 kpc from us and about 1.3 kpc out of the Galactic disk. However,
they do not reject the possibility of G 116.6$-$26.1 being the result of a
core collapse SN explosion. In favor of this scenario is the fact that
probably the dust distribution in its neighborhood is affected by G
116.6$-$26.1 , which would imply that this object is placed within 300 pc from
us, resulting from a core collapse supernova explosion (SN II). In a very
recent work, Churazov et al. (2022) report radio emission from G 116.6$-$26.1
detected in the LOFAR Two-metre Sky Survey (LoTTS-DR2, Shimwell et al., 2022).
It presents a faint, shell-like structure with good positional coincidence
between the radio boundary and the X-ray limb, as presented in Churazov et al.
(2021), while no optical counterparts have been reported until now.

In our Galaxy there are roughly 300 known SNRs (Green, 2019), with more than
90 per cent of them located within 5 degrees off the Galactic plane (Kothes et
al., 2017), while $\sim$ 10 of the currently known remnants reside between
latitudes 5 and 10 degrees. However, only a handful of these objects have been
found in higher Galactic latitudes. Boumis et al. (2002) reported the
discovery of optical filaments in Pegasus Constellation suggesting that are
part of one (or more) SNRs. Its nature as such was also confirmed by Fesen et
al. (2015) based on its morphology and its spectral characteristics. This is
the G 70.0$-$21.5 remnant which was studied more recently also by Raymond et
al. (2020), and it is believed to be the result of a type Ia SN. McCullough et
al. (2002) discovered the Antlia SNR (or G 275.5$+$18.4), which was detected
in H$\upalpha$ and X-rays. Later, Shinn et al. (2007) and Fesen et al. (2021)
confirmed this identification, using GALEX FUV imagery and optical
observations, respectively. Its progenitor is probably a B-type star (Shinn et
al., 2007) and hence, it comes likely from a core collapse SN explosion.

Another high latitude SNR is the Hoinga or G 249.7$+$24.7 remnant discovered
also in the X-ray SRG/eROSITA All-Sky Survey eRASS1 (Becker et al., 2021),
which presents also radio, optical, and UV emission (Becker et al., 2021;
Fesen et al., 2021). According to Becker et al. (2021), the fact that no
pulsar has been associated with this remnant so far, in combination with its
high latitude, indicate that it is probably the remnant of a type Ia SN
explosion. The highest Galactic latitude SNR found yet, is G 354.0$-$33.5,
observed in FUV, H$\upalpha$ , and radio wavelengths (Testori et al., 2008;
Fesen et al., 2021). No evidence about its progenitor and/or the explosion
type of this SNR is found.

In this work we report for the first time optical emission from G 116.6$-$26.1
, based on both wide-field and higher resolution imaging as well as
spectroscopic observations. The results seem to satisfy most of the criteria
for the optical identification of SNRs. This, in combination with its spatial
correlation with other wavelengths, enhances the belief of G 116.6$-$26.1
being a new Galactic SNR. We also discuss the possible nature of its
progenitor but we did not reach a secure conclusion. The structure of the
paper is as follows: In Section 2 we describe the observations and the
reduction of the imaging and spectroscopic data. In Section 3 we present the
results of the aforementioned analysis. We quote available observations of G
116.6$-$26.1 in other wavelengths and we explore their spatial correlation
with the optical emission in Section 4. In Section 5 we discuss our results,
while in Section 6 we summarise our conclusions.

##  2 Observations and Data Reduction

###  2.1 Imaging

####  2.1.1 Wide-field imagery

The first step in our search for possible optical line emission from the
newly-discovered SNR G 116.6$-$26.1 , was to perform wide-field imaging, since
the reported angular diameter of the source in X-rays was estimated to be
${\sim}4\degr$ (Churazov et al., 2021). Such a capability is provided by the
fast-optics (f/3.2), Schmidt-Cassegrain 0.3 m telescope (a Lichtenknecker Flat
Field Camera, hereafter LFFC) at Skinakas Observatory in Crete, Greece.
Coupled with a back-illuminated 2k $\times$ 2k CCD camera (by Andor Tech.),
the telescope–instrument system offers a field of view (FoV) $100\arcmin\times
100\arcmin$ on the sky. The thermoelectric-cooled CCD sensor (operating
temperature $-70\degr$ Celsius) shows negligible dark current, has pixel size
13.5 $\micron$ which corresponds to $2\aas@@fstack{\prime\prime}94$ on the
celestial sphere.

This preliminary observing run aimed at the construction of a mosaic of
partially overlapping images, wide enough to cover the reported X-ray extent
of the source, in the light of H$\upalpha$ emission, in a $3\times 3$ frame
pattern. The images were obtained with the LFFC telescope on 2021 September 1,
3–5 and October 7–8, through a narrow-band H$\upalpha$+$[\text{N\,{ii}}]$
filter. A red broadband continuum filter (SDSS$-$r’) was also used to provide
images for subsequent subtraction of the background continuum and especially
for the removal of the stellar component in the images in order to improve the
visibility of very faint structures. Filter characteristics are given in Table
1. Each of the nine fields in the mosaic was observed for a total exposure
time of 9000 s (15$\times$600 s) in H$\upalpha$+$[\text{N\,{ii}}]$ and 300 s
(15$\times$20 s) through the continuum filter, respectively.

All nights were photometric and target fields ranged in airmass between 1.0
and 1.4 during observations. Multiple bias exposures and twilight flat frames
of very high signal-to-noise (S/N) ratio were taken on a daily basis. We did
not observe any spectrophotometric standard stars, because we were not
planning to flux-calibrate our images taken with the LFFC 0.3-m telescope,
serving the single purpose to visually examine the area for any H$\upalpha$
emission signature.

Table 1: The characteristics of the optical filters used in this research.

[b]  Filter $\uplambda_{c}$a FWHM Transmb (Å) (Å) (%)
H$\upalpha$+$[\text{N\,{ii}}]$ $\uplambda\uplambda$ 6548, 6584 Å 6582 82 99
$[\text{S\,{ii}}]$ $\uplambda\uplambda$ 6716, 6731 Å 6730 32 99
$[\text{O\,{iii}}]$ $\uplambda$ 5007 Å 5008 25 52 SDSS$-$r’ 6214 1290 96
Johnson/Bessell V 5380 980 88

  * a

The central wavelength

  * b

Peak filter transmittance

Figure 1: The continuum-subtracted 9-frame mosaic towards the SNR G
116.6$-$26.1 in the light of H$\upalpha$+$[\text{N\,{ii}}]$ . Superposed
contours (solid green lines) from the X-ray observations (Churazov et al.,
2021) show the correlation between the optical filaments with the X-ray
emission. Several artefacts in the form of dark round halo appear around
bright stars which are not emission features but ghosts due to scattered light
from internal reflections in the narrow-band filter.

The raw images were reduced using the IRAF software (Tody, 1986, 1993)
procedures for bias subtraction and flat-field correction. Individual
exposures of each of the 9 frames observed were registered to a reference
image in the set and combined through a 3-sigma clipping average algorithm in
order to remove cosmic rays and artificial satellite trails – a not so rare
situation in modern era, especially in wide-field imaging. Astrometric
calibration was performed with the aid of astrometry.net (Lang et al., 2010)
web service, interfaced through AstroImageJ software (Collins et al., 2017).
The next step was to use the narrow-band (H$\upalpha$+$[\text{N\,{ii}}]$ ) and
continuum (SDSS-r’) images in order to subtract the sky background and
eliminate as many field stars as possible to reduce crowding confusion,
following a procedure described in Appendix A. Finally, the continuum-
subtracted frames were assembled together into the final mosaic utilizing the
Montage111http://montage.ipac.caltech.edu software with adjustment for smooth
background levels between overlapping images. The final mosaic in
H$\upalpha$+$[\text{N\,{ii}}]$ emission is shown in Fig. 1, where the presence
of two major optical filaments is evident.

The detected sharp filaments shown in Fig. 1 were further investigated, using
the LFFC telescope and H$\upalpha$+$[\text{N\,{ii}}]$ , $[\text{S\,{ii}}]$ and
$[\text{O\,{iii}}]$ filters, on 2021 November 6 and 8, for a total exposure
time of 3 h (18 images of 600 s each), 4 h (24$\times$600 s) and 3 h
(9$\times$1200 s), through the respective filters. Multiple short exposures
were also obtained in the respective continuum filters for background and
starlight subtraction. The final background-subtracted images are shown in
Fig. 2.

Figure 2: Low resolution negative images of the major optical filaments’ area,
in H$\upalpha$+$[\text{N\,{ii}}]$ (top), $[\text{S\,{ii}}]$ (middle) and
$[\text{O\,{iii}}]$ (bottom panel). Note that the dark spots and disc-like
features scattered around the images, are residuals from imperfect subtraction
of bright stars. These leftover artefacts in the bottom panel
($[\text{O\,{iii}}]$ filter) have no counterpart in the top and middle (red
filters) images. Figure 3: The area of the major optical filaments (Fig. 2) in
H$\upalpha$+$[\text{N\,{ii}}]$ (red), $[\text{S\,{ii}}]$ (green) and
$[\text{O\,{iii}}]$ (blue), to indicate the relative position of the emission
in these three lines near the shock fronts.

####  2.1.2 Higher angular resolution imaging

Guided by the wide-field mosaic of the object’s H$\upalpha$+$[\text{N\,{ii}}]$
emission, additional images were taken of the upper and lower optical
filaments at the south-eastern edge of the X-ray structure, using the 1.29 m
(f/7.6) Ritchey–Chrétien telescope at Skinakas Observatory. Images were
acquired with a 2048 $\times$ 2048 back-illuminated, deep-depletion CCD camera
(Andor iKon-L) with pixel size 13.5 $\micron$. The system provided a
9$\aas@@fstack{\prime}$6 $\times 9\aas@@fstack{\prime}6$ FoV and image scale
$0\aas@@fstack{\prime\prime}28$ per pixel. The narrow-band filters used are
these in Table 1. Observations were held on 2021 July 15, September 10 and 12,
October 4–5 and November 5, 7 and 11, under good seeing conditions
($0\aas@@fstack{\prime\prime}8$–$1\aas@@fstack{\prime\prime}2$) and high
target elevation (airmass range 1.00–1.35). Images for a 3-frame mosaic were
obtained in the H$\upalpha$+$[\text{N\,{ii}}]$ filter covering the brighter
part of the upper filament (region UF in Fig. 1) with 9600 s (16$\times$600 s)
total integration time for each frame. Another series of 16 $\times$ 600 s
exposures in $[\text{S\,{ii}}]$ combined into a single-frame image of the
brighter part of UF. The lower filament (designated region LF in Fig. 1) was
also imaged – only with a partial filter coverage due to weather conditions –
in H$\upalpha$+$[\text{N\,{ii}}]$ with 2 overlapping frames of 2 h total
exposure time each ($12\times 600$ s).

Table 2: Spectroscopic observations log.

[b]  RegionID Date Aperture center Aperture Exposuresa $\mathrm{X_{eff}}$b
R.A. (J2000) Dec (J2000) height (UT) (″) (s) UF 2021 Nov 3
$\mathrm{00^{h}29^{m}03\aas@@fstack{s}7}$ $\mathrm{+35^{d}51^{m}50^{s}}$ 22.5
$3\times 2400$ 1.035, 1.008, 1.002

  * a

Number of individual exposures and common exposure time

  * b

Effective airmass at each exposure.
[$\mathrm{X_{eff}~{}=~{}\frac{1}{6}(X_{start}\,+\,4\,X_{mid}\,+\,X_{end})}$]

Table 3: Locations of fainter filaments detected in wide-field images, but not
examined yet in detail.

[b]  Filament center Approximate Appears Image R.A. (J2000) Dec (J2000) length
(′) in figure filter $\mathrm{00^{h}25^{m}15^{s}}$
$\mathrm{+37^{d}15^{m}28^{s}}$ 6.7 Fig. 1 H$\upalpha$+$[\text{N\,{ii}}]$
$\mathrm{00^{h}25^{m}22^{s}}$ $\mathrm{+36^{d}52^{m}40^{s}}$ 18.5 Fig. 1
H$\upalpha$+$[\text{N\,{ii}}]$ $\mathrm{00^{h}24^{m}22^{s}}$
$\mathrm{+34^{d}57^{m}33^{s}}$ 7.7 Fig. 1 H$\upalpha$+$[\text{N\,{ii}}]$
$\mathrm{00^{h}24^{m}46^{s}}$ $\mathrm{+34^{d}55^{m}09^{s}}$ 6.9 Fig. 1
H$\upalpha$+$[\text{N\,{ii}}]$ $\mathrm{00^{h}24^{m}08^{s}}$
$\mathrm{+35^{d}01^{m}07^{s}}$ 2.6 Fig. 1 H$\upalpha$+$[\text{N\,{ii}}]$
$\mathrm{00^{h}30^{m}10^{s}}$ $\mathrm{+35^{d}51^{m}28^{s}}$ 8.6 Fig. 2
$[\text{O\,{iii}}]$ $\mathrm{00^{h}30^{m}40^{s}}$
$\mathrm{+35^{d}58^{m}48^{s}}$ 13.0 Fig. 2 $[\text{O\,{iii}}]$
$\mathrm{00^{h}25^{m}00^{s}}$ $\mathrm{+35^{d}18^{m}13^{s}}$ 12.6 Fig. 2
$[\text{O\,{iii}}]$ $\mathrm{00^{h}31^{m}38^{s}}$
$\mathrm{+37^{d}08^{m}18^{s}}$ 21.3 Fig. 12 Far UV

###  2.2 Spectrophotometry

Long-slit, low-resolution spectroscopic observations of the brightest part of
the upper filament were made on 2021 November 3, using the 1.29 m telescope at
Skinakas Observatory and a CCD camera of the same type and characteristics as
in the higher resolution direct imaging, described in section 2.1.2. A 1302
grooves mm-1 grating was used, blazed at 5500 Å, giving an average dispersion
0.94 Å pixel-1 and spectral coverage 4820–6750 Å. The object spectra were
taken with a slit of width 3$\aas@@fstack{\prime\prime}$9
(7$\aas@@fstack{\prime\prime}$3 slit used for the spectrophotometric standard
stars), and usable length 14$\aas@@fstack{\prime}$5 – long enough to allow for
background subtraction – always oriented along the North-South direction.
Seeing varied between 1$\aas@@fstack{\prime\prime}$1 and
1$\aas@@fstack{\prime\prime}$5 under dark and photometric sky conditions.

Several zero-exposure (bias) images were collected each night along with two
sets of exposures for flat-fielding. One set consisted of images taken with a
quartz-lamp as light source inside the spectrograph and a diffuser in front of
the slit, while the second set was a series of twilight sky exposures, in
order to perform a slit illumination correction on the first set. Spectra were
wavelength calibrated using a Ferrum-Helium-Neon-Argon comparison lamp and arc
images were obtained after the end of each target exposure. Flux calibration
of the spectral lines was achieved through exposures of the spectrophotometric
standard stars HR718, HR1544, HR3454, HR7596, HR7950, HR9087 (Hamuy et al.,
1992, 1994) and BD$+25{\degr}4655$, BD$+33\degr 2642$, G 191-B2B (Oke, 1990;
Massey et al., 1988). Standards were observed every night in at least three
different time slots, in groups of 4–6 stars for better airmass coverage.

One more set of images was acquired, necessary for the correction of optical
distortion and curvature of the spectral features along the dispersion and
spatial directions, which is quite noticeable in our long-slit spectra. A
preselected nearby (angular distance ≲ 8° from the observed target), very
bright field star (visual magnitude 3.5–5.5) was positioned near one of the
slit’s long-dimension ends and a series of spectral images was obtained,
offseting the telescope in steps of $\approx 30\arcsec$ along the slit
direction in-between exposures, until the other end of the slit was reached.
This task is accomplished in less than 5 minutes, since typical exposure times
for high signal-to-noise ratio individual spectra is 3–5 s. The images are
added together during data reduction and give the equivalent of a multi-star
spectrum, which, along with a comparison lamp exposure, can determine the
geometric transformation needed to correct distortions in the two-dimensional
science spectra.

Standard IRAF procedures (Valdes, 1992) were utilised for spectra reduction
and extraction of calibrated line emission fluxes. The reduction steps include
bias subtraction, flat-field correction, cosmic-ray removal using the software
L. A. Cosmic (van Dokkum, 2001) and geometric rectification
222https://iraf.net/irafdocs/spect.pdf. Background estimation was performed in
areas close to the spectral apertures and free of diffuse emission, selected
with the guidance of the higher resolution images. Center coordinates and
height of the aperture used in the spectra extraction, exposure date,
integration times and effective airmasses are presented in Table 2.

Figure 4: Higher angular resolution mosaic in H$\upalpha$+$[\text{N\,{ii}}]$
of the upper filament’s (UF) brighter segment. An orderly network of similarly
aligned, hair-thin filaments of varying degree of curvature is prominent in
this image.

Figure 5: The middle frame of the H$\upalpha$+$[\text{N\,{ii}}]$ mosaic in
Fig. 5, as seen in $[\text{S\,{ii}}]$ emission light. Although images in both
filters were obtained under the same seeing conditions, the filaments’
sharpness in $[\text{S\,{ii}}]$ is not comparable to that in H$\upalpha$
light.

Figure 6: Uncalibrated bi-colour composite image of the area in Fig. 5 with
H$\upalpha$+$[\text{N\,{ii}}]$ (red) and $[\text{S\,{ii}}]$ (green). The
respective emissions at the brightest filament area seem to start at the same
location (dark orange coloured post-shock shell), but the H$\upalpha$
intensity drops faster in the downstream direction compared to
$[\text{S\,{ii}}]$ . The yellow, cyan and blue rectangles mark the locations
of the slit, aperture for the extracted spectrum and sky area for background
subtraction, respectively.

Figure 7: Two-frame mosaic in H$\upalpha$+$[\text{N\,{ii}}]$ of the brighter
segments in region LF. Several thin filaments arranged along NE-SW direction,
suggesting the presence of expanding emission shells, are mixed with fainter
diffuse light, mostly in the downstream direction of the associated shock
fronts. It should be noticed that the perfectly straight feature, running
diagonally across the upper-right part of the image and marked with two yellow
arrows, is neither an artificial satellite trail nor real thin filament. It is
an optical artefact, a diffraction spike produced by copious light from the
bright star HD 2093 (V = 7.36), just outside the lower right edge of the
image.

##  3 Results

###  3.1 Direct Imaging

####  3.1.1 Wide-field coverage

The deep (total exposure time 22.5 h) H$\upalpha$+$[\text{N\,{ii}}]$
continuum-subtracted 9-frame mosaic of the area where candidate SNR G
116.6$-$26.1 was discovered in X-rays, is shown in Fig. 1. The image center is
at R.A. = $00^{\rm h}23^{\rm m}14^{\rm s}$, Dec. = $36^{\rm d}25^{\rm
m}45^{\rm s}$ and has an angular length of $4\aas@@fstack{\circ}25$ on each
side. Overlaid are X-ray contours from the SRG imaging observations (Churazov
et al., 2021). The prominent set of two optical filaments (marked as UF –
upper filament – and LF – lower filament) at the southeastern part of the
X-ray image, running almost in parallel directions, are the major features
seen in this wide-field image and there is very good positional match with the
brightest contours of the X-ray map.

A few much fainter filaments are seen scattered, mostly at the eastern part of
the SNR (near the left edge of the image, at Declination
$36\aas@@fstack{\circ}5$). Two more, very faint, straight and parallel to each
other, filaments are located just north-east of the presumed center of the
remnant [at (RA, Dec) $\simeq$ ($00^{\text{h}}25^{\text{m}}$, $37\degr
10\arcmin$)]. Another wider diffuse stripe is extending radially outwards the
north-western corner of the image, which is also barely visible in the DSS red
images and in the All-Sky H$\upalpha$
map333https://lambda.gsfc.nasa.gov/product/foreground/fg_diffuse.cfm with
$6\arcmin$ resolution (Finkbeiner, 2003). This latter feature seems to be
unrelated to the SNR emission. We did not investigate these fainter structures
further, but will be included in our follow-up observations.

Fig. 2 presents a close-up (but still in low angular resolution) of the
filaments, in H$\upalpha$+$[\text{N\,{ii}}]$ , $[\text{S\,{ii}}]$ and
$[\text{O\,{iii}}]$ . UF is brighter in H$\upalpha$+$[\text{N\,{ii}}]$ and
$[\text{S\,{ii}}]$ and marginally visible in $[\text{O\,{iii}}]$ . It appears
slightly curved at its brighter portion, with the convex part facing towards
the center of the SNR, and extends (in the red emission lines) for
$\simeq\,38\arcmin$. On the other hand, in the $[\text{O\,{iii}}]$ light, UF
appears fainter but seems it can be traced to a longer extend, getting a bit
brighter towards its southwestern projected end, to the west of star HD 2154
(marked in bottom panel of Fig. 2).

The lower filament (LF) is more elongated compared to UF, with an estimated
angular extend at least $70\arcmin$ (measured in the $[\text{S\,{ii}}]$
image), and despite its high angular length, it appears to be very straight,
indicating it is part of a large-radius expanding shell and minimal
interaction with any large-scale denser gaseous bodies (diffuse or molecular
clouds). If we trace visually the lower filament, starting from its brighter
projected end next to HD 2093 (see top panel in Fig. 2), its brightness
remains more-or-less unchanged in all three emission lines, for about one
third of its angular length, where a bifurcation occurs and
$[\text{O\,{iii}}]$ emission is dimmed heavily. The northern branch, much
fainter and sharper than the lower one, follows a remarkably straight path and
is visible in both H$\upalpha$+$[\text{N\,{ii}}]$ and $[\text{S\,{ii}}]$
light. The lower branch appears brighter, a little wider and its continuity is
interrupted for a short distance before its brightness weakens completely.
South of the brighter part of LF and close to it (angular separation $\lesssim
6\arcmin$) a few short and fainter filaments are seen, running along in
parallel to the major feature. Another network of curved filaments is visible
– mostly in H$\upalpha$+$[\text{N\,{ii}}]$ – to the west of HD 2093 (top panel
in Fig. 2).

We summarise in Table 3 the locations of fainter filamentary structures
appearing in our wide-field imagery, which were not investigated any further
in this preliminary report. The celestial coordinates of each filament’s
estimated center along with their approximate angular length (in arcminutes)
are listed there. In the last two columns of Table 3, the corresponding Figure
number and the filter used in obtaining the associated image are given as
well.

Figure 3 shows in colour representation the spatial distribution of the line
emissions in H$\upalpha$+$[\text{N\,{ii}}]$ (red), $[\text{S\,{ii}}]$ (green)
and $[\text{O\,{iii}}]$ (in blue hues). The shock fronts associated with the
two filaments in this picture, are moving towards the lower left direction. In
the UF, both H$\upalpha$+$[\text{N\,{ii}}]$ and $[\text{S\,{ii}}]$ contribute
in brightness, with an apparent spatial precedence of the H$\upalpha$
emission. No appreciable light is seen in $[\text{O\,{iii}}]$ there. However,
evidence for $[\text{O\,{iii}}]$ emission is provided both in the bottom frame
of Fig. 2 and, more convincing, in the spectrum taken in region UF (see Sec.
3.2). Also, in the bright section of LF, the $[\text{O\,{iii}}]$ emission lags
behind the H$\upalpha$ emitting front.

####  3.1.2 Higher angular resolution imaging

The two major filaments detected in our deep wide-field images can now be seen
in greater detail in higher resolution H$\upalpha$+$[\text{N\,{ii}}]$ mosaics,
in Figs. 5 and 7. A multitude of thin emission filaments exist, arranged
almost in parallel directions with no apparent major distortion through
interactions with the surrounding medium or signs of recent strong large-scale
instabilities. Fig. 5 shows the $[\text{S\,{ii}}]$ emission at the region
covered by the middle frame in the H$\upalpha$+$[\text{N\,{ii}}]$ mosaic in
Fig. 5. Only the brighter filaments present in the H$\upalpha$ image are seen
in $[\text{S\,{ii}}]$ , where the intensity scale-length in the emitting shell
behind the shock front seems to be greater than that in H$\upalpha$ . This is
clearly depicted in Fig. 7, where the H$\upalpha$ (red) emission brightness
falls-off more quickly behind the shock than $[\text{S\,{ii}}]$ does.

At the brighter part of emission in the LF region a similar picture of several
thin H$\upalpha$+$[\text{N\,{ii}}]$ filaments in close spacing arrangements is
shown in Fig. 7. Near the southern frame’s lower edge, a short segment of the
smaller companion filament is also visible (see Figs. 2 and 3).

Figure 8: Emission line spectrum across filaments in region UF. Figure 9: The
colour excess $E(B-V)$ as a function of distance towards the region UF in G
116.6$-$26.1 , according to the 3D dust maps of Lallement et al. (2019) and
Green (2019).

[b]

Table 4: Normalised emission line fluxes from the extracted spectrum in the UF
filament.

${\rm\uplambda}$ (Å), ID | Extinction function, $\text{f(}\uplambda)$a | F($\uplambda$)b | I($\uplambda$)c | S/Nd
---|---|---|---|---
4861 ${\rm H\upbeta}$ | 0.0000 | 100.0 | 100.0 | 8.1
4959 $[\text{O\,{iii}}]$ | -0.0297 | 51.3 | 51.1 | 3.5
5007 $[\text{O\,{iii}}]$ | -0.0446 | 123.6 | 122.6 | 7.7
6300 $[\text{O\,{i}}]$ | -0.2835 | 151.3 | 143.6 | 8.1
6364 $[\text{O\,{i}}]$ | -0.2924 | 55.2 | 52.3 | 6.1
6563 ${\rm H\upalpha}$ | -0.3224 | 382.1 | 360.0 | 39.2
6583 $[\text{N\,{ii}}]$ | -0.3256 | 75.9 | 71.4 | 12.1
6716 $[\text{S\,{ii}}]$ | -0.3458 | 138.5 | 129.9 | 8.7
6731 $[\text{S\,{ii}}]$ | -0.3479 | 75.5 | 70.8 | 5.2
Observed ${\rm F(H\upbeta})$e |  | $7.75\,\pm\,0.9$ |  |
${\rm c(H\upbeta)}$f |  | $0.080\,\pm\,0.17$ |  |
${E(B-V)}$g |  | $0.054\,\pm\,0.12$ |  |

  * a

Derredening function according to Fitzpatrick et al. (2019) extinction law and
extinction ratio R(V) = 3.1. See Appendix B for details.

  * b

Observed line flux normalised to ${\rm H\upbeta=100}$.

  * c

Intrinsic flux ratio, i.e. flux relative to ${\rm H\upbeta=100}$ corrected for
interstellar extinction.

  * d

Signal to noise ratio (S/N) for the observed flux ratio. Error contribution
from the ${\rm H\upbeta}$ flux is not included in the noise term.

  * e

In units of $10^{-18}$ erg Å-1 s-1 cm-2 arcsec-2.

  * f

Logarithmic ${\rm H\upbeta}$ extinction coefficient: ${\rm
c(H\upbeta)=log[(H\upbeta)_{intrinsic}/(H\upbeta)_{observed}]}$.

  * g

$B-V$ colour excess.

Table 5: Major dereddened flux ratios of the UF spectrum. Line ratio | Value
---|---
[S ii] $\uplambda\uplambda$ 6716, 6731 / H$\upalpha$ | $0.56\,\pm\,0.06$
$[\text{N\,{ii}}]$ $\uplambda$ 6583 / ${\rm H\upalpha}$ | $0.20\,\pm\,0.02$
$[\text{O\,{iii}}]$ $\uplambda$ 5007 / ${\rm H\upbeta}$ | $1.23\,\pm\,0.22$
$[\text{O\,{i}}]$ $\uplambda$ 6300 / ${\rm H\upalpha}$ | $0.40\,\pm\,0.05$
${\rm I(6716)/I(6731)}$ | $1.83\,\pm\,0.41$
$[\text{N\,{ii}}]$ $\uplambda$ 6583 / [S ii] $\uplambda\uplambda$ 6716, 6731 | $0.36\,\pm\,0.05$

###  3.2 Optical spectroscopy

A low resolution, long-slit spectrum was acquired at a bright filament
position in region UF (see Table 2) with the Skinakas spectrograph, mount on
the 1.29 m RC telescope. Individual exposures on the target were aligned
spatially (via spectra of stars in the slit) and along the dispersion
direction (using sky emission lines), corrected for atmospheric extinction and
the 2D images were averaged. Extraction aperture was selected to sample
relatively brighter region across the filaments and background was estimated
from emission free regions, at the closest possible distance from the
extraction aperture. The exact location of the slit aperture used for object
and background extraction is shown overlaid in Fig. 7. Line fluxes were
measured fitting Gaussians over a linear local background estimate.

The resulting flux-calibrated spectrum is presented in Fig. 8. Correction of
the line emission fluxes for attenuation effects due to interstellar
extinction was performed with the Balmer decrement method. We used the
Fitzpatrick et al. (2019) extinction law corresponding to the mean galactic
ratio of total-to-selective extinction R(V) = 3.1, and details on the
derivation of the extinction function are given in Appendix B.

The usual practice in performing flux dereddening is to select, as a first
step, a value for the intrinsic Balmer decrement
$\text{I(H}\upalpha)/\text{I(H}\upbeta)$, appropriate for the type of gaseous
nebula studied. In the SNR case, the adopted values are either
$\text{I(H}\upalpha)/\text{I(H}\upbeta)$ = 2.86 – the theoretical Case B
recombination (nebula optically thick to the Lyman series photons) – seen
mostly in H ii regions, or $\text{I(H}\upalpha)/\text{I(H}\upbeta)$ = 3.0 when
conditions favour the likelihood of significant collisional excitation of
hydrogen in the post-shock region (e.g. Raymond et al., 2020; Fesen et al.,
2021). Indeed, models with self-consistent treatment of the precursor
ionization for a wide range of shock velocities (v${}_{s}=$ 10–1500 km
$s^{-1}$) indicate that the intensity ratio
$\text{I(H}\upalpha)/\text{I(H}\upbeta)$ exceeds always the Case B value 2.86
and varies between $\sim$ 3.0–5.0. In lower-velocity shocks the H$\upalpha$
/H$\upbeta$ flux ratio is strongly enhanced by collisional excitation in the
postshock gas (Raymond, 1979; Chevalier et al., 1980; Cox & Raymond, 1985;
Hartigan et al., 1987; Sutherland & Dopita, 2017; Dopita & Sutherland, 2017).
The derived value of the extinction $A(V)$ and colour excess $E(B-V)$ are then
compared to estimates from independent sources.

Figure 10: Part of the long slit spectral image in region UF, showing the
spatial correlation of major emission lines identified with labels at the
bottom. The [O i] $\uplambda\uplambda 6300,6364$ lines are not included for
reasons explained in the text. The vertical axis scale corresponds to the
angular offset from the extraction aperture center (value 0) across the slit
aligned along the N-S direction. Positive numbers correlate with northern
displacements (post-shock region). The white vertical stripe separates the
bluer (left) from the redder wavelengths and the two horizontal black bands
running across the image width are spectra of stars aligned in the slit.
Letters and short horizontal line segments in red indicate the location of
filaments discussed in the text.

In this study, we decided to take a different approach and adopt a value for
the Balmer decrement $\text{I(H}\upalpha)/\text{I(H}\upbeta)$ such that the
resulting colour excess $E(B-V)$ agrees with the extinctions predicted by the
3D dust maps of Lallement et al. (2019)444https://astro.acri-st.fr/gaia_dev/
and Green (2019) 555http://argonaut.skymaps.info/query in the direction of G
116.6$-$26.1 . Both dust maps show a gradual increase of $E(B-V)$ to 0.055 at
a distance $\sim 400$ pc, and remains flat to larger distances (see Fig. 9).
This requirement leads to intrinsic Balmer decrement
$\text{I(H}\upalpha)/\text{I(H}\upbeta)\,=\,3.6$, which corresponds to shock
velocity range 75 $\lesssim$ (vs / km s-1) $\lesssim$ 120, as predicted by
relevant models (see Sutherland & Dopita, 2017, fig. 18). It should be
emphasized that the emission line ratios, useful in diagnostic criteria, are
insensitive to the reddening parameters used, because by construction, they
correspond to wavelengths close to each other, and thus line fluxes are almost
equally affected by interstellar extinction.

Dereddened fluxes were calculated with the aid of Equations 9 and B, derived
in Appendix B, and the colour excess applying Eq. B. Listed in Table 5 are the
measured $\text{F}(\uplambda)$ and extinction corrected $\text{I}(\uplambda)$
fluxes, normalised to $\text{F(H}\upbeta)=100$ and $\text{I(H}\upbeta)=100$,
respectively, as well as the logarithmic H$\upbeta$ extinction coefficient
c(H$\upbeta$) and colour excess $E(B-V)$. The S/N values given in the last
column of Table 5 do not include calibration errors, estimated to be less than
15 per cent, neither the contribution of the normalising H$\upbeta$ flux
error. It should be also noticed the relatively high uncertainty in the
interstellar extinction related values of c(H$\upbeta$) and $E(B-V)$,
reflecting the high relative error in the H$\upbeta$ flux (8 per cent).

The spectrum of a filament in region UF presents clear evidence that the
associated emission results from shock-heated gas, commonly seen in SNRs. The
widely accepted criterion supporting this is the observed value of [S ii] /
H$\upalpha$ = 0.56 $\pm$ 0.06, exceeding the threshold 0.4, which separates
SNRs from H ii regions (Dodorico, 1978; Blair et al., 1981; Dopita et al.,
1984; Fesen et al., 1985). Another sensitive shock emission test is provided
through the [O i] / H$\upbeta$ ratio, found to be 0.40 at a significance level
of $8\,\upsigma$. This line ratio does not exceed 0.1 in H ii regions while
higher values are associated with emission from shock heated interstellar
nebulosities, like the shocks in SNRs (e.g. Fesen et al., 1985; Kopsacheili et
al., 2020).

The electron density-sensitive [S ii] ratio $\uplambda\uplambda$ 6716 / 6731
was found to be above the low-density limit, [0.436 $\leq$ 6716 / 6731 $\leq$
1.434 (Proxauf et al., 2014; Osterbrock & Ferland, 2006)], but with a high
uncertainty,  I(6716) / I(6731) = $1.83\,\pm\,0.41$. Therefore, we can not
estimate the electron density (in the zone where S+ recombines) with
confidence but only state that n${}_{e}\,\lesssim 100$ cm-3.

Another remarkable property revealed from the spectrum sampled in the UF
filament is its low brightness. The measured flux F(H$\upalpha$) = (2.96 $\pm$
0.08) $\times 10^{-17}$ erg Å-1 s-1 cm-2 arcsec-2 is 1–3 orders of magnitude
fainter than most of the Galactic SNRs observed in the optical (e.g. Alikakos
et al., 2012; Boumis et al., 2005, 2008; Boumis et al., 2009; Mavromatakis et
al., 2002; Mavromatakis et al., 2005; Mavromatakis et al., 2009; How et al.,
2018).

###  3.3 Post-shock emission morphology

Figure 10 presents two sections of the negative, background-subtracted 2D
spectral image in the region UF filaments. The spatial variation of emitted
fluxes in several emission lines – marked at the bottom of the picture – is
shown with respect to angular distance from the aperture center (at value 0)
used for the spectrum extraction, along the slit, aligned in the N-S
direction. Positive numbers on the left axis scale correspond to northern
displacements (down-stream direction relative to the shock fronts). Marked
also with letters and short horizontal red line segments are the locations of
filaments, as identified from the higher resolution
H$\upalpha$+$[\text{N\,{ii}}]$ image in Fig. 5. It should be mentioned that
only light emitted from filaments A and B is included in the aperture of the
extracted spectrum in Fig. 8.

Figure 11: GALEX-FUV image covering the area of G 116.6$-$26.1 . The dashed
blue circle defines the X-ray extend of the supernova remnant and the $\times$
symbol marks its center. The filamentary structures near the center of the
south-eastern quadrant are located at the same position as in our
H$\upalpha$+$[\text{N\,{ii}}]$ wide-field image in Fig.1.

Figure 12: Contours of radio emission in 120–168 MHz from LOFAR LoTSS DR2
images superposed on our continuum-subtracted wide field
H$\upalpha$+$[\text{N\,{ii}}]$ image (Fig. 1). The coincidence of the major
optical filaments with the enhanced radio brightness at the south-eastern edge
is evident. Contours shown correspond to 1, 5 and 20 $\upmu$Jy beam-1 levels
of flux density in the radio data.

Despite the fact that [O i] $\uplambda\uplambda$ 6300,6364 emission lines are
equally important in SNRs, we do not include the corresponding part in the 2D
spectral image in Fig. 10. The reason is that these lines are also quite
strong in the night sky spectrum and the background subtraction procedure does
not provide a smooth result at the positions of the [O i] lines. Thus, in our
low resolution spectroscopy, the absence of considerable velocity component in
the SNR [O i] lines to produce sufficient Doppler shift, combined with the
intrinsic faintness of the nebular emission (even in H$\upalpha$) yield
relatively ‘noisy’ background-subtracted [O i] lines and result in a rather
confusing picture.

Starting from the most southern filament C, this appears as a diffuse and very
faint feature which is seen in the spectrum to emit basically H$\upalpha$
light and just a weak hint of H$\upbeta$ emission at nearly the noise level.
Filaments A and B are projected close to each other, at an angular distance
$\simeq 19\arcsec$. H$\upalpha$ emission starts ahead of the brightest parts
of the filaments, along with some H$\upbeta$ and weak $[\text{O\,{iii}}]$
light, but no detection of emission from other low-excitation lines, such as
[$\text{N\,{ii}}]$ $\uplambda\uplambda$ 6548,6583 and [S ii]
$\uplambda\uplambda$ 6716,6731. The H$\upalpha$ intensity increases for a
short distance and then slowly decreases over a relatively long interval
($\simeq\,100\arcsec$) before it fades completely. The second Balmer line
(H$\upbeta$ ) follows a similar variation but becomes undetectable much
earlier, shortly after its peak. Heavier ions, like N+ and S+, start emitting
further downstream, behind the location of maximum H$\upalpha$ emission and
for a short angular distance, roughly 30$\arcsec$.

In contrast, in the line $[\text{O\,{iii}}]$ $\uplambda$ 5007, emission can be
traced for much longer distance, seen almost as far as the H$\upalpha$ light,
but in low intensity. The detected $[\text{O\,{iii}}]$ emission in the post-
shock region suggests a lower limit on the shock velocity vs of about 70 km
s-1, while its low brightness implies that this value is not expected to be
much higher than 90 km s-1. These limits were found from the diagram in fig. 6
of Dopita & Sutherland (2017), using the measured line ratio log([O iii]
$\uplambda$5007 / H$\upbeta$) = 0.09 $\pm$ 0.08 (see Table 5), and come to
agreement with the implied shock velocity range through our earlier choice for
the intrinsic Balmer decrement $\text{I(H}\upalpha)/\text{I(H}\upbeta)$, in
the extinction correction procedure of the spectral line fluxes.

One additional advantage of using the line ratio [O iii] $\uplambda$5007 /
H$\upbeta$ to constrain the shock velocity range is the very weak dependency
of its value on density (Dopita & Sutherland, 2017). However, the lower limit
on the shock velocity implied by the detection of [O iii] $\uplambda$5007
emission, is based on the assumption that O++ ions are not present in the
preshock gas, or else [O iii] emission lines can be produced in much slower
shocks (Raymond, 1979). This latter situation cannot be excluded from our
spectrum, since the preshock area is projected on the postshock region of
another, much weaker shock-front, leading ahead of the brighter filament in
this region.

A similar structure behind relatively slow shocks (v${}_{s}\lesssim$ 150 km
s-1) was found in optical spectroscopic investigations of other high galactic
latitude SNRs, such as G 70.0$-$21.5 (Raymond et al., 2020), G 107.0$+$9.0
(Fesen et al., 2020) and in the Hoinga or G 249.5$+$24.5 SNR (Becker et al.,
2021; Fesen et al., 2021). The coexistence of radiative and Balmer dominated
shock fronts in these supernova remnants indicate variable velocities of
expansion in filamentary structures at different locations of the remnants.

##  4  Observations at other wavelengths

###  4.1 X-rays

G 116.6$-$26.1 has been discovered in X-rays, through the SRG/eROSITA all-sky
survey (Churazov et al., 2021), as a faint, very extended circular feature
with a total 0.3–2 keV flux of about 3 $\times$ 10-11 erg s-1cm-2 distributed
over a projected area $\sim$ 11.8 deg2. The soft X-ray image depicts the
circular source with a slightly brightened rim, especially across the south-
eastern edge and two compact knots (seen in the smoothed version of the X-ray
image), one near the western ridge and the other close to the center of the
object. Comparing the location of the two major optical filaments in Fig. 1
with the overlaid X-ray brightness contours, there seems to be a clear
correlation of the optical features with one of the stronger X-ray emission
parts of the remnant at south-west. A second area of enhanced X-ray brightness
appears at the northern part of the remnant, where no optical counterpart was
found thus far.

###  4.2 Ultraviolet data

Our search for ultraviolet counterpart to the optical filaments has been
performed with the publicly available imagery of the Galaxy Evolution Explorer
(GALEX) space telescope All-Sky Imaging survey (AIS), described in detail in
Morrissey et al. (2007). The GALEX archival database contains far-UV (FUV,
$\mathrm{\uplambda_{eff}\sim}$1528 Å, 1344–1786 Å) and near-UV (NUV,
$\mathrm{\uplambda_{eff}\sim}$2310 Å, 1771–2831 Å) direct imaging and grism
spectra.

Figure 12 presents the GALEX FUV intensity mapping of the SNR G 116.6$-$26.1
area, enclosed in a dashed circle of radius 2∘ determined from the X-ray
discovery image. Evidently, the two parallel filaments present in the middle
of the south-east quadrant in the FUV image, are registered at exactly the
same location as in our optical image and have similar lengths. These
filaments are also visible in the GALEX NUV image, although they appear quite
fainter. The NUV image is not shown here because the numerous artefacts in it
produce a rather confusing picture of the field. One more, much shorter but
quite sharp filament can be seen near the eastern edge of the FUV image
(Declination $37\degr 00\arcmin$–$37\degr 20\arcmin$), which is not clearly
seen in our wide field H$\upalpha$+$[\text{N\,{ii}}]$ image, a point that will
be re-examined during our follow-up observations.

###  4.3 Observations in Radio frequencies

We have also used archival data from the large scale survey GALFA-H i DR2 of
Galactic neutral hydrogen in the 21 cm radio line (Peek et al., 2018), to
search for possible radio emission in the supernova remnant area, partially
included in declination within the field covered by the survey. The available
data cover a velocity range from $-$188.4 to $+$188.4 km s-1 with respect to
the Local Standard of Rest, in slices with spacing 0.184 km s-1. Several
features are present in the area but no association with the visible filaments
or the overall X-ray contour morphology could be established.

However, Churazov et al. (2022) reported recently detection of extremely faint
radio emission towards G 116.6$-$26.1 from the LOFAR Two-metre Sky Survey
(LoTSS) DR2 images (Shimwell et al., 2022). These data include 120–168 MHz
images, mapping 27 per cent of the northern sky in low frequency at 6$\arcsec$
resolution, offering full bandwidth Stokes I total intensity maps, linear
polarization image cubes and circular polarization continuum images.

In order to examine the spatial correlation between the optical filaments with
the radio emission, we constructed a mosaic from archival LOFAR images of the
G 116.6$-$26.1 projected area of size $5^{\circ}\times 5^{\circ}$, partially
removed point-like radio sources and created a network of contours at flux
density levels 1, 5 and 20 $\upmu$Jy beam-1. Fig. 12 shows these radio
contours overlaid on our continuum subtracted H$\upalpha$+$[\text{N\,{ii}}]$
wide field image (in Fig. 1). The major parts of the radio emission structure
consist of two elongated linear features, almost parallel to each other and
running along the NE to SW direction, located at the diagonals of the SE and
NW quadrants of the image. These linear structures’ northern ends are
connected by a faint circular arc, completing a figure reminiscent of a horse-
shoe. The eastern branch of this feature is the brightest and coincides with
the location of the two major optical filaments, as seen in Fig. 12.

The north-western segment may indicate the location for another possible
optical feature. Close inspection of our H$\upalpha$+$[\text{N\,{ii}}]$ image
reveals the presence of some diffuse and extended optical emission, displaced
towards the NW relative to the radio contours, making the association between
radio and optical emission questionable. It should be also noticed that in
this area material of elevated opacity is present – indicated by the white
shade in the negative representation of the H$\upalpha$+$[\text{N\,{ii}}]$
image – which, if located in the foreground of the remnant, makes any possible
optically enhanced emission in this region hardly visible.

###  4.4 Infrared and sub-millimetre wavelengths

Churazov et al. (2021) presented the IRAS $100\micron$ image of the G
116.6$-$26.1 area and compared the infrared emission map to the X-ray
structures, noticing the existence of a partial anticorrelation between the
two. They raised the possibility that the dust distribution may have been
influenced by G 116.6$-$26.1 , in which case the dust and the remnant are
located roughly at the same distance from us, about 300 pc. Based on possible
future confirmation of this neighbouring, the authors propose an alternative
scenario for G 116.6$-$26.1 being a nearby, core-collapse SNR. Our preliminary
data do not suffice to confirm or disprove this possible scenario.

We conducted an extensive search for counterparts at other infrared and sub-
millimeter wavebands, but did not find any signs of morphological similarity
with either the optical or X-ray structure of G 116.6$-$26.1 . The data used
for the search stem from surveys such as the WISE All-Sky (all four bands, at
3.4, 4.6, 12 and 22 $\micron$), the AKARI Far-infrared All-Sky Survey (the
WideS band with centre wavelength 90 $\micron$), Plank Legacy Archive (at
frequencies 143, 217, 353, 545 and 857 GHz and the CO[$1\rightarrow 0$],
CO[$2\rightarrow 1$], CO[$3\rightarrow 2$] emission maps derived from the 100,
217 and 353 GHz channels of the Planck High Frequency Instrument).

##  5 Discussion

As already noted in Section 1, G 116.6$-$26.1 belongs to the few SNRs at high
Galactic latitudes (|b| > 15∘) that have been observed and detected in the
optical, along with other wavelengths. These are G 70.0$-$21.5 (Boumis et al.,
2002; Fesen et al., 2015; Raymond et al., 2020), G 249.7$+$24.7 (the so-called
Hoinga SNR, Becker et al., 2021; Fesen et al., 2021), G 275.5$+$18.4 (Antlia
SNR, McCullough et al., 2002; Shinn et al., 2007; Fesen et al., 2021) and G
354.0$-$33.5 (Testori et al., 2008; Fesen et al., 2021). Given their large
size, faint emissivity and estimates for low shock velocities ($\sim$ 70–100
km s-1), these SNRs seem to go through their mature phases of their lives.

The current work comes to present evidence, for the first time, for the
existence of optical emission from G 116.6$-$26.1 . Through our imaging and
spectroscopic observations, two filamentary structures are revealed (UF and
LF) in a partial shell formation. Several fine filaments in the UF region
denote optical emission originating from shock-heated gas, usually associated
with the shock fronts of the remnants. At the same time, the compelling
sequence of emission lines can be easily traced, depicting the post-shock
(emission) morphology of the source (Figs 7 and 10). From the available
spectrum, we see that the one out of the two strong optical filaments examined
presents a low shock velocity ($\sim$70–100 km s-1). The latter, in
conjunction with the referred X-ray/radio angular size, seems to also place G
116.6$-$26.1 in the class of older SNRs.

G 116.6$-$26.1 was primarily identified as a potential SNR on the basis of its
soft X-ray properties (Churazov et al., 2021). Sources that present low-
temperature ($<$ 2 keV), emission line spectrum from hot, ionized gas are
regarded as thermal X-ray emitting SNRs (e.g. Garofali et al., 2017; Leonidaki
et al., 2010). Optical evidence for shock-excited gas arises from the
$[\text{S\,{ii}}]$ / H$\upalpha$ emission line ratio ($>$0.4) in one of its
two distinct filaments. Further confirmation regarding the nature of this
filament comes from the fulfillment of 9 out of the 10 applicable diagnostic
tests described in Kopsacheili et al. (2020) that contain the emission line
ratios [O iii] / H$\upbeta$ , [O i] / H$\upalpha$ , $[\text{S\,{ii}}]$ /
H$\upalpha$ , and [N ii] / H$\upalpha$ . All the aforementioned, in
combination with the spatial correlation of the optical structures with some
of the brightest X-ray, UV and radio emission features (see Section 4 and
Figs. 1 and 12) provide robust evidence for establishing G 116.6$-$26.1 as a
new Galactic SNR.

Owing to their position and low surface brightness, the group of faint and
large remnants at high Galactic latitudes have escaped detection and an in-
depth study, since most of the SNR surveys focus on the vicinity of the
Galactic Plane, where wealth of dust and gas help stars to flourish and
eventually die. Especially in the case of core-collapse SNRs, this area is the
ideal nursery for their birth while their existence above the usual regions of
research (|b| $\leq 10^{\circ}$) is less expected, since white dwarf binary
systems, responsible for Type Ia supernovae explosions, are more likely to be
found at higher galactic latitudes. However, three out of five high galactic
latitude SNRs mentioned previously, have been classified in the literature as
core-collapse SNRs.

In an attempt to cautiously advocate in favor or against a specific SNR
progenitor, various criteria should be taken into account. The basic
characteristic of Type Ia SNRs comes from their spectrum: i.e. the existence
of prominent narrow and broad emission line components of H$\upalpha$ emission
and the absence (or sometimes weakness) of forbidden lines. This occurs
because they are thought to be surrounded by a medium consisting mainly of
neutral hydrogen, producing non-radiative shocks (e.g. Chevalier et al.,
1980): i.e. shocks for which their radiative cooling time is much longer than
their characteristic dynamical time scale (e.g. age). Given the
aforementioned, they are expected to be found within regions isolated from
massive stars / OB associations that ionize the ambient medium (Maggi et al.,
2016; Franchetti et al., 2012). Furthermore, the well-established criterion
for identifying core-collapse SNRs ($[\text{S\,{ii}}]$ /H$\upalpha$ $>$0.4)
drops down to less than 0.05 in the case of Type Ia SNRs (Lin et al., 2020).
Another criterion that has been used to separate core-collapse from Type Ia
SNRs is their morphology. Assuming an unperturbed surrounding, it has been
demonstrated (Lopez et al., 2011) that Type Ia SNRs present a symmetric
explosion and morphology. However, the notion that Type Ia SNRs do not shape
their environment has started to be challenged the last few years and various
exceptions cannot be ruled out (Zhou & Vink, 2018). For example, based on
models reproducing the observables, it has been argued that famous cases such
as the historical Kepler and Tycho SNRs, present extensive massive outflows
from their progenitor binary, strongly modifying this way their circumstellar
medium (e.g. Chiotellis et al., 2012; Zhou et al., 2016).

On the basis of the derived spectroscopic properties of the UF optical
filament, G 116.6$-$26.1 presents a typical spectrum of a core-collapse SNR,
with enhanced $[\text{S\,{ii}}]$ / H$\upalpha$ emission line ratio and
various, definite forbidden lines denoting the shock excitation of the
surrounding medium. Although it is only one spectrum we have in hand,
significant information can be drawn since it is part of the two prominent
optical filaments of the entire SNR. However, the extracted H$\upalpha$ flux
of the optical upper filament, with absolute flux 1-3 orders of magnitude
lower compared to other Galactic SNRs, supports the notion of a Type Ia
progenitor (Franchetti et al., 2012). What could further endorse this scenario
is the almost circular/symmetric morphology seen in the X-ray/radio images. As
can be realised, the existing data do not allow us to explicitly decide over
one progenitor scenario or the other. The determination of the distance of G
116.6$-$26.1 (whether it is part of the disc or the Galactic halo, for
example) is still ambiguous and which is, by the way, a well-known caveat in
the research of Galactic SNRs. A twofold scenario may also be the case for G
116.6$-$26.1 as reported in other Galactic SNRs (e.g. Fesen et al., 2020):
shock-heated line emission may coexist with Balmer-dominated filaments.

##  6 Conclusions

In this preliminary investigation for the existence of an optical counterpart
to the candidate SNR G 116.6$-$26.1 , we successfully detected emission from
related filamentary stuctures, through deep wide-field and higher angular
resolution CCD imaging of the area, followed by supplementary
spectrophotometric observations.

First detection was achieved after the construction of a deep, wide-field
image mosaic ($4\aas@@fstack{\circ}25\times 4\aas@@fstack{\circ}25$, scale
${\sim}3\arcsec$ pixel-1) through a narrowband H$\upalpha$+$[\text{N\,{ii}}]$
filter, covering the footprint of the X-ray discovering image. Two distinct,
major filaments were revealed, located at the southeastern ridge of the X-ray
structure, in a partial shell-like formation. These features were detected in
$[\text{S\,{ii}}]$ and $[\text{O\,{iii}}]$ emission as well. Images acquired
in higher angular resolution, mostly in H$\upalpha$+$[\text{N\,{ii}}]$ ,
present networks of sharp filaments in both areas, running almost parallel to
each other with no significant signs of disturbance from dynamical
interactions with the ambient medium or internal large-scale instabilities
(indicating a rather old remnant).

Our deep flux-calibrated and dereddened low-resolution spectrum, obtained at a
single location of the northern filament, allowed us to confirm the nature of
G 116.6$-$26.1 as a most likely Galactic SNR. The spectroscopic data show
faint but significant emission from lines typically observed in SNR sources,
such as H$\upalpha$ , H$\upbeta$ , $[\text{O\,{iii}}]$ ($\uplambda\uplambda$
4959, 5007), [O i] ($\uplambda\uplambda$ 6300, 6364), $[\text{N\,{ii}}]$
($\uplambda$ 6583) and $[\text{S\,{ii}}]$ ($\uplambda\uplambda$ 6716, 6731).
The $[\text{S\,{ii}}]$ /H$\upalpha$ line ratio was found greater than 0.5,
indicating emission from shock-heated gas, and 9 out of the 10 applicable
emission-line ratio diagnostic tests presented in Kopsacheili et al. (2020),
turned out to be positive. The measured $[\text{O\,{iii}}]$ / H$\upbeta$ line
ratio suggests relatively low shock-front velocities (roughly 70–100 km s-1).

Comparison with multiwavelength data confirmed collocation between the optical
filamentary signature of the SNR G 116.6$-$26.1 with emission in X-rays and
low-frequency radio, detected and reported by Churazov et al. (2021) and
Churazov et al. (2022), respectively. A search in the GALEX online archive
helped us to identify fine FUV-emitting filaments, which overlap almost
perfectly with the brightest portions of their optical counterparts.

We also discussed possible scenarios about the nature of the progenitor star
but could not reach a definite conclusion due to luck of supporting data.

Although we have presented strong evidence that G 116.6$-$26.1 belongs to the
Galactic SNR group, several of its properties remain undetermined. Our ongoing
follow-up imaging and spectroscopic investigation will help to uncover the
physical properties of the entire remnant (e.g. plasma conditions, kinematics,
environment, distance, evolutionary stage) and shed more light on its nature.

## Acknowledgements

We would like to thank the anonymous referee for his/her valuable feedback and
constructive comments that improved the content of our research paper.

This research made use of Montage. It is funded by the National Science
Foundation under Grant Number ACI-1440620, and was previously funded by the
National Aeronautics and Space Administration’s Earth Science Technology
Office, Computation Technologies Project, under Cooperative Agreement Number
NCC5-626 between NASA and the California Institute of Technology.

This publication utilizes data from Galactic ALFA HI (GALFA HI) survey data
set obtained with the Arecibo L-band Feed Array (ALFA) on the Arecibo 305m
telescope. The Arecibo Observatory is operated by SRI International under a
cooperative agreement with the National Science Foundation (AST-1100968), and
in alliance with Ana G. Méndez-Universidad Metropolitana, and the Universities
Space Research Association. The GALFA HI surveys have been funded by the NSF
through grants to Columbia University, the University of Wisconsin, and the
University of California.

LOFAR data products were provided by the LOFAR Surveys Key Science project
(LSKSP; https://lofar-surveys.org/) and were derived from observations with
the International LOFAR Telescope (ILT). LOFAR (van Haarlem et al., 2013) is
the Low Frequency Array designed and constructed by ASTRON. It has observing,
data processing, and data storage facilities in several countries, which are
owned by various parties (each with their own funding sources), and which are
collectively operated by the ILT foundation under a joint scientific policy.
The efforts of the LSKSP have benefited from funding from the European
Research Council, NOVA, NWO, CNRS-INSU, the SURF Co-operative, the UK Science
and Technology Funding Council and the Jülich Supercomputing Centre.

This research has made use of the SIMBAD database, operated at CDS,
Strasbourg, France (Wenger et al., 2000).

M.K. acknowledges funding from the European Research Council under the
European Union’s Seventh Framework Programme (FP/2007-2013)/ERC Grant
Agreement No 617001. She also acknowledges support from the European Research
Council under the European Union’s Horizon 2020 research and innovation
program, under grant agreement No 771282.

## Data Availability

The optical images and spectrum taken at Skinakas Observatory and used in this
research will be available upon a reasonable request to the corresponding
author.

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##  Appendix A Image subtraction procedure

The image subtraction procedure described here was used _only_ when the
background-subtracted image was needed for visual examination and detection of
possible faint optical emission. The continuum (C) image was registered to the
narrow-band (NB) image using the astropy (Astropy Collaboration et al., 2013,
2018) package astroalign (Beroiz et al., 2020), which is achieved by finding
similar 3-point asterisms (triangles) in both images and estimating the affine
transformation between them.

Next we performed a background noise estimation to set the threshold pixel
value for source detection that follows. This was based on procedures of the
package SEP (Barbary, 2016), a stand-alone PYTHON adaptation of the widely
used program SEXTRACTOR (Bertin & Arnouts, 1996) for source detection and
photometry. The method used consists of the following steps, applied to each
of the two images separately. (1) Saturated pixels are masked and a spatially
variable background is fit. (2) A background-subtracted image is computed
along with an estimate of the background noise for each image pixel. (3)
Source detection follows with a threshold set to 3$\sigma$, where $\sigma$ is
the global root-mean-square (RMS) of the background noise. (4) Sources are
deblended and a catalogue is prepared with source centroid positions and
fluxes (equivalent to the FLUX_AUTO in SEXTRACTOR), measured over elliptical
apertures centered at each detected source and corrected for local background
contribution.

The sources detected in the NB and C images (already aligned at an earlier
stage) were matched on source centroid proximity criteria with STILTS, the
command-line sister package to TOPCAT (Taylor, 2005). The number of matched
stars per single frame (FoV $100\arcmin\times 100\arcmin$) was in the range of
5500–11000, depending on the narrow-band filter type, and their flux ratio
flux(NB) / flux(C) was computed. This set of flux ratios was iteratively
sigma-clipped, with 3$\sigma$ as the limit; the removed outliers were mostly
the brighter of the stars. The average value of the sigma-clipped set was used
to normalise the continuum image before subtracting it from the narrow-band.
The continuum-subtracted result was always inspected visually. In a few cases,
instead of the mean value, we used either the median or the 75 percentile of
the cleaned ratios set.

In order to reduce the effect of over- or under- subtraction of the star light
and enhance visibility of faint emission, we applied a 3$\times$3 median
filter to the whole continuum-subtracted image. When this smoothing did not
produce a satisfactory result (too many ’dark holes’ left in the image), we
replaced the pixel values within the elliptical apertures – used previously in
the flux estimation – with the median value of pixels within a narrow
elliptical ring around the aperture and an additive Gaussian noise component
with local characteristics.

##  Appendix B Interstellar extinction correction

Emission line ratios were corrected for interstellar reddening using the
normalized extinction curve of Fitzpatrick et al. (2019) (F19 hereafter),
produced by combining Hubble Space Telescope/STIS and International
Ultraviolet Explorer spectrophotometry and Two Micron All-Sky Survey broadband
NIR photometry. Their extinction curves are defined through the quantities:

$\displaystyle k(\uplambda-55)$
$\displaystyle\equiv\frac{E(\uplambda-55)}{E(44-55)}=\frac{A(\uplambda)-A(55)}{A(44)-A(55)}$
(1) $\displaystyle=k^{\prime}(\uplambda)\quad\text{(for shorter notation)}$
$\text{and}\quad R(55)\equiv\frac{A(55)}{E(44-55)}$ (2)

where $A(\uplambda)$ is the total extinction (in magnitudes) at wavelength
$\uplambda$, $E(\uplambda_{1}-\uplambda_{2})$ denotes the colour excess (that
is, the difference between the observed colour and the expected colour in the
absence of interstellar absorption) and the numbers 44 and 55 correspond to
monochromatic wavelengths of 4400 Å and 5500 Å, respectively, instead of the
most commonly used Johnson B and V filters.

If $F(\uplambda)$ and $I(\uplambda)$ denote the observed and intrinsic flux,
respectively, of a nebular emission line at wavelength $\uplambda$, then the
colour excess of the Balmer decrement can be written as:

$\displaystyle E({\rm H\upbeta-H\upalpha})$ $\displaystyle=A({\rm
H\upbeta})-A({\rm H\upalpha})$ $\displaystyle=[A({\rm
H\upbeta})-A(55)]-[A({\rm H\upalpha})-A(55)]$ $\displaystyle=[k^{\prime}({\rm
H\upbeta})-k^{\prime}({\rm H\upalpha})]E(44-55)$ (3)

or

$E(44-55)=\frac{E({\rm H\upbeta-H\upalpha})}{k^{\prime}({\rm
H\upbeta})-k^{\prime}({\rm H\upalpha})}$ (4) $\text{Similarly,}\quad
E(\uplambda-{\rm H\upbeta})=[k^{\prime}(\uplambda)-k^{\prime}({\rm
H\upbeta})]E(44-55)$ (5) $\displaystyle\text{ On the other hand, since\
}I(\uplambda)=F(\uplambda)\ 10^{0.4\,A(\uplambda)},\text{it follows that}$
$\displaystyle E({\rm H\upbeta-H\upalpha})$ $\displaystyle=-2.5~{}\left\\{{\rm
log}\left[\frac{F({\rm H\upbeta})}{F({\rm H\upalpha})}\right]-{\rm
log}\left[\frac{I({\rm H\upbeta})}{I({\rm H\upalpha})}\right]\right\\}$
$\displaystyle=2.5~{}{\rm log}\left[\frac{F({\rm H\upalpha})/F({\rm
H\upbeta})}{I({\rm H\upalpha})/I({\rm H\upbeta})}\right]$ (6)

Combining equations B, 4 and B, we get:

$\displaystyle{\rm log}\left[\frac{I(\uplambda)}{I({\rm H\upbeta})}\right]$
$\displaystyle={\rm log}\left[\frac{F(\uplambda)}{F({\rm H\upbeta})}\right]$
$\displaystyle+\frac{k^{\prime}(\uplambda)-k^{\prime}({\rm
H\upbeta})}{k^{\prime}({\rm H\upbeta})-k^{\prime}({\rm H\upalpha})}{\rm
log}\left[\frac{F({\rm H\upalpha})/F({\rm H\upbeta})}{I({\rm
H\upalpha})/I({\rm H\upbeta})}\right]$ (7)

Two more useful quantities need to be evaluated; the logarithmic H$\upbeta$
extinction coefficient, ${\rm c(H\upbeta)=log\left[\frac{I({\rm
H\upbeta})}{F({\rm H\upbeta})}\right]}$, and the colour excess $E(B-V)$. ${\rm
c(H\upbeta)}$ can be obtained from equations 1, 2, 4 and B:

$k^{\prime}({\rm H\upbeta})=\frac{A({\rm
H\upbeta})-A(44)}{E(44-55)}=\frac{2.5\ {\rm c(H\upbeta)}}{E(44-55)}-R(55)$

which gives

$\displaystyle{\rm c(H\upbeta)}$ $\displaystyle=0.4\ E(44-55)\
[k^{\prime}({\rm H\upbeta})+R(55)]$ $\displaystyle=\frac{k^{\prime}({\rm
H\upbeta})+R(55)}{k^{\prime}({\rm H\upbeta})-k^{\prime}({\rm H\upalpha})}\
{\rm log}\left[\frac{F({\rm H\upalpha})/F({\rm H\upbeta})}{I({\rm
H\upalpha})/I({\rm H\upbeta})}\right]$ (8)

Then, the dereddening relation B can be written as:

$\frac{I(\uplambda)}{I({\rm H\upbeta})}=\frac{F(\uplambda)}{F({\rm
H\upbeta})}10^{{\rm
c(H\upbeta)}f(\uplambda)},\text{where}\>f(\uplambda)=\frac{k^{\prime}(\uplambda)-k^{\prime}({\rm
H\upbeta})}{k^{\prime}({\rm H\upbeta})+R(55)}$ (9)

We used a series of cubic splines to fit the F19 average extinction curve in
$k^{\prime}(\uplambda)$ for the Milky Way, which corresponds to $R(V)=3.1$ and
$R(55)=3.02$, in the wavelength range 3300–10000 Å. The anchor points were
placed at values of $x$ (inverse wavelength) from Table 3 in F19 and
interpolation on the fitted curve gives $k^{\prime}({\rm H\upbeta})=0.59183$
and $k^{\prime}({\rm H\upalpha})=-0.57246$. Then, equation B becomes:

${\rm c(H\upbeta)}=3.1022\ {\rm log}\left[\frac{F({\rm H\upalpha})/F({\rm
H\upbeta})}{I({\rm H\upalpha})/I({\rm H\upbeta})}\right]$ (10)

The authors in F19 also provide values for the colour excess ratio
$r=E(B-V)/E(44-55)$ which, however, depend on the effective temperature of the
observed stars and the overall extinction [as measured by $E(44-55)$]. For
each value of $E(44-55)$, we have chosen as representative the fixed value
that would give the same area under the curve ($r,~{}T_{\rm eff}$), that is:

$\langle r\left[E(44-55)\right]\rangle=\frac{\int_{T_{\rm eff,min}}^{T_{\rm
eff,max}}r\,dT_{\rm eff}}{\int_{T_{\rm eff,min}}^{T_{\rm eff,max}}\,dT_{\rm
eff}}\text{, and taking the average over the}$

two $E(44-55)$ values, we obtain

$\displaystyle E(B-V)$ $\displaystyle=0.98\ E(44-55)=\frac{0.98\times
2.5}{k^{\prime}({\rm H\upbeta})+R(55)}$ $\displaystyle=0.6873\ {\rm
c(H\upbeta)}$ (11)

We give in Table 6 values of the extinction function $f(\uplambda$) (Eq. 9)
for several optical emission lines of astrophysical interest.

Table 6: Extinction function (defined in Eq. 9) for several optical emission lines. $\lambda$(Å) ID | f($\uplambda$) |  $\uplambda$(Å) ID | f($\uplambda$) |  $\uplambda$(Å) ID | f($\uplambda$) |  $\uplambda$(Å) ID | f($\uplambda$) |  $\uplambda$(Å) ID | f($\uplambda$)
---|---|---|---|---|---|---|---|---|---
3703.86 H i | 0.2879 | 4387.93 He i | 0.1160 | 6101.83 [K iv] | -0.2580 | 7331.40 [Ar iv] | -0.4266 | 8776.97 He i | -0.5657
3705.02 He i | 0.2875 | 4437.55 He i | 0.1016 | 6118.20 He ii | -0.2601 | 7451.43 [Ar iii] | -0.4380 | 8799.00 He ii | -0.5681
3711.97 H i | 0.2852 | 4471.50 He i | 0.0908 | 6233.80 He ii | -0.2747 | 7499.84 He i | -0.4421 | 8816.64 He i | -0.5700
3721.63 [S iii] | 0.2822 | 4541.59 He ii | 0.0693 | 6300.34 [O i] | -0.2835 | 7530.83 [Cl iv] | -0.4448 | 8845.37 He i | -0.5731
3721.94 H i | 0.2821 | 4625.53 [Ar v] | 0.0478 | 6310.80 He ii | -0.2850 | 7592.74 He ii | -0.4503 | 8914.77 He i | -0.5806
3726.03 [O ii] | 0.2809 | 4685.68 He ii | 0.0349 | 6312.10 [S iii] | -0.2851 | 7751.06 [Ar iii] | -0.4685 | 8929.11 He ii | -0.5822
3728.82 [O ii] | 0.2801 | 4711.37 [Ar iv] | 0.0298 | 6363.78 [O i] | -0.2924 | 7751.12 [Ar iii] | -0.4685 | 8997.02 He i | -0.5894
3734.37 H i | 0.2785 | 4713.17 He i | 0.0295 | 6406.30 He ii | -0.2986 | 7751.43 [Ar iii] | -0.4686 | 9068.60 [S iii] | -0.5968
3750.15 H i | 0.2741 | 4714.17 [Ne iv] | 0.0293 | 6527.11 He ii | -0.3169 | 7816.16 He i | -0.4775 | 9108.54 He ii | -0.6009
3770.63 H i | 0.2684 | 4715.66 [Ne iv] | 0.0290 | 6548.10 [N ii] | -0.3201 | 8045.63 [Cl iv] | -0.5050 | 9123.60 [Cl ii] | -0.6024
3797.90 H i | 0.2600 | 4724.15 [Ne iv] | 0.0273 | 6562.77 ${\rm H\upalpha}$ | -0.3224 | 8116.60 He i | -0.5103 | 9174.49 He i | -0.6074
3805.74 He i | 0.2575 | 4725.62 [Ne iv] | 0.0271 | 6583.50 [N ii] | -0.3256 | 8168.91 He i | -0.5137 | 9210.33 He i | -0.6108
3819.62 He i | 0.2531 | 4740.17 [Ar iv] | 0.0243 | 6678.16 He i | -0.3400 | 8236.77 He ii | -0.5177 | 9303.19 He i | -0.6191
3835.39 ${\rm H\upeta}$ | 0.2481 | 4861.33 ${\rm H\upbeta}$ | 0.0000 | 6683.20 He ii | -0.3408 | 8237.15 He i | -0.5177 | 9344.93 He ii | -0.6226
3868.75 [Ne iii] | 0.2383 | 4921.93 He i | -0.0177 | 6716.44 [S ii] | -0.3458 | 8265.71 He i | -0.5194 | 9367.03 He ii | -0.6244
3888.65 He i | 0.2330 | 4931.80 [O iii] | -0.0209 | 6730.82 [S ii] | -0.3479 | 8421.99 He i | -0.5306 | 9463.58 He i | -0.6317
3889.05 ${\rm H\upzeta}$ | 0.2329 | 4958.91 [O iii] | -0.0297 | 6795.00 [K iv] | -0.3573 | 8444.69 He i | -0.5325 | 9516.63 He i | -0.6353
3926.54 He i | 0.2241 | 5006.84 [O iii] | -0.0446 | 6890.88 He i | -0.3710 | 8451.20 He i | -0.5331 | 9530.60 [S iii] | -0.6361
3964.73 He i | 0.2162 | 5015.68 He i | -0.0471 | 7005.67 [Ar v] | -0.3868 | 8480.85 [Cl iii] | -0.5357 | 9603.44 He i | -0.6404
3967.46 [Ne iii] | 0.2157 | 5047.74 He i | -0.0555 | 7065.25 He i | -0.3948 | 8486.31 He i | -0.5362 | 9625.70 He i | -0.6416
3970.07 ${\rm H\upepsilon}$ | 0.2151 | 5191.82 [Ar iii] | -0.0926 | 7135.80 [Ar iii] | -0.4038 | 8519.35 He ii | -0.5392 | 9702.71 He i | -0.6451
4009.26 He i | 0.2067 | 5197.90 [N i] | -0.0943 | 7160.56 He i | -0.4069 | 8528.99 He i | -0.5401 | 9762.15 He ii | -0.6474
4026.21 He i | 0.2026 | 5200.26 [N i] | -0.0950 | 7177.50 He ii | -0.4090 | 8566.92 He ii | -0.5438 | 9824.13 [C i] | -0.6492
4068.60 [S ii] | 0.1912 | 5411.52 He ii | -0.1451 | 7237.26 [Ar iv] | -0.4162 | 8578.69 [Cl ii] | -0.5450 | 9850.26 [C i] | -0.6498
4076.35 [S ii] | 0.1890 | 5517.66 [Cl iii] | -0.1681 | 7262.76 [Ar iv] | -0.4191 | 8581.87 He i | -0.5453 | 10023.10 He i | -0.6514
4101.74 ${\rm H\updelta}$ | 0.1815 | 5537.60 [Cl iii] | -0.1732 | 7281.35 He i | -0.4212 | 8608.31 He i | -0.5479 | 10027.72 He i | -0.6514
4120.84 He i | 0.1758 | 5577.34 [O i] | -0.1819 | 7298.04 He i | -0.4231 | 8626.19 He ii | -0.5498 | 10123.61 He ii | -0.6503
4143.76 He i | 0.1692 | 5754.60 [N ii] | -0.2059 | 7318.92 [O ii] | -0.4253 | 8632.97 He i | -0.5505 | 10311.00 He i | -0.6442
4168.97 He i | 0.1625 | 5875.66 He i | -0.2240 | 7319.99 [O ii] | -0.4254 | 8648.25 He i | -0.5520 |  |
4340.47 ${\rm H\upgamma}$ | 0.1268 | 6036.70 He ii | -0.2490 | 7329.67 [O ii] | -0.4264 | 8733.43 He i | -0.5610 |  |
4363.21 [O iii] | 0.1218 | 6074.10 He ii | -0.2543 | 7330.73 [O ii] | -0.4266 | 8739.97 He i | -0.5617 |  |