IFU spectroscopy of Southern Planetary Nebulae: III
In this paper we describe integral field spectroscopic observations of four southern Galactic Planetary Nebulae (PNe), M3-4, M3-6, Hen2-29 and Hen2-37 covering the spectral range; 3400-7000Å. We derive the ionisation structure, the physical conditions, the chemical compositions and the kinematical characteristics of these PNe and find good agreement with previous studies that relied upon the long-slit technique in their co-spatial area. From their chemical compositions as well as their spatial and kinematic characteristics, we determined that Hen2-29 is of the Peimbert Type I (He and N rich), while the other three are of Type II. The strength of the nebular He II line reveals that M3-3, Hen2-29 and Hen2-37 are of mid to high excitation classes while M3-6 is a low excitation planetary nebula (PN). A series of emission-line maps extracted from the data cubes were constructed for each PN to describe its overall structure. These show remarkable morphological diversity. Spatially resolved spectroscopy of M3-6, shows that the recombination lines of C II, C III, C IV and N III are of nebular origin, rather than arising from the central star as had been previously proposed. This result increases doubts regarding the weak emission-line star (WELS) classification raised by Basurah et al. (2016). In addition, they reinforce the probability that most genuine cases of WELS are arise from irradiation effects in close binary central stars (Miszalski, 2009).
keywords:ISM: abundances - Planetary Nebulae: Individual M3-4, M3-6, He2-29, Hen2-37
The vast majority of the spectroscopic studies of PNe up to the present day have relied upon long slit spectroscopic techniques. However these measurements sample only a portion of the complete nebula, and are of necessity weighted towards the high-ionisation regions around the central star (CS). An accurate determination of the physical and chemical nebular parameters, as well as the determination of their global parameters requires a knowledge of their integrated spectra and spatial structure. The advent of Integral Field Units (IFUs) now provides an opportunity to obtain this data, and to build fully self-consistent photoionization models for PNe. The IFU technique as applied to PNe was pioneered by Monreal-Ibero et al. (2005) and Tsamis et al. (2007). Recently, detailed physical and morpho-kinematical studies using optical IFU data have been taken by Danehkar & Parker (2015) and Danehkar (2015) to study the planetary nebulae Hen 3-1333, Hen 2-113 and Th 2-A.
The Wide Field Spectrograph (WiFeS) instrument mounted on the 2.3m ANU telescope at Siding Spring Observatory (Dopita et al., 2007; Dopita et al., 2010) offers the ability to perform such IFU spectroscopy, since it is capable of reaching a spatial resolution of 1.0″, a spectral coverage of 3200–8950Å and a spectral resolution of . Combined with a field of view of 25″x 38″, it is very well-suited to integral field spectroscopy of compact PNe. The first paper in this series by Ali et al. (2015b) used WiFeS to study the large, evolved and interacting planetary nebula PNG 342.0-01.7, generating an IFU mosaic to cover the full spatial extent of the object. In the second, Basurah et al. (2016) provided a detailed analysis of four highly excited non-type I PNe which casts doubt on the general applicability of the WELS classification. Here, and in upcoming studies, we aim to further exploit the capabilities of the WiFeS instrument to provide a new database on previously studied, bright and compact PNe. Specifically our objectives are to:
Create emission-line maps for PN in any diagnostic emission line within its spectral range.
Provide integrated spectra of the whole PN, and if possible, of its exciting star.
Analyse these spectra both in the forbidden and recombination lines to derive chemical abundances and to understand any differences between results obtained here and through long-slit observations of the same objects.
Determine expansion velocities and the kinematic nature of the PN.
Build self-consistent photoionisation models to derive abundances, physical conditions within the nebula, to determine distances, and to place the PNe on the Hertzsprung-Russell Diagram in order to derive central star masses, the evolutionary status and the nebular age.
Emission-line mapping of PNe has previously been obtained through the technique of narrow-band interference filter imaging, e.g. Manchado et al. (1996); Górny et al. (1999); Hajian et al. (2007); Miranda et al. (2010), Aller et al. (2015) and García-Rojas et al. (2016). The difficulty here is that each image is very expensive in observing time, and only one emission line can be be imaged at once. On the other hand, IFU spectroscopy makes available the narrow band images of all lines which are observed with sufficient signal to noise ratio. Emission line maps and line ratio maps allow us to structural details and physical processes. For example,[O III]/H ratio maps are useful to study the variation of ionization and chemical abundances and also to look for signatures of collimated outflows and shocks (Guerrero et al., 2013), [SII]/H ratio maps are particularly sensitive to shocked regions (Akras & Gonçalves, 2016; Akras et al., 2016), whereas HeII/H ratio maps identify the very high ionization regions (Vázquez, 2012).
The main objective of the current paper is to study four Galactic planetary nebulae – M3-4, M3-6, He2-29 and Hen2-37 – which have so far received relatively little attention. In this paper we will use WiFeS data to examine the correlations between nebular morphology and excitation. Furthermore, the WiFeS instrument is ideally suited to study spatial variations in nebular parameters, such as extinction, electron temperature, density and ionic abundances.
Górny et al. (1999) imaged Hen2-29 and Hen2-37 in a narrow band H filter and determined angular sizes of 20″16″and 27″ 24″, respectively. Corradi et al. (1998) classified Hen2-29 & Hen2-37 nebulae as elliptical PNe with modest ellipticity (major to minor axis length ratio 1.3) and outer irregular contours. Searching the literature and the ESO archive, it would appear that no narrow band images are available for M3-4 and M3-6. However, spectroscopic studies of all four PNe are to be found distributed across several papers, e.g., Milingo et al. (2002), Martins & Viegas (2000), Maciel & Quireza (1999), Kingsburgh & Barlow (1994), Chiappini & Maciel (1994) and Perinotto et al. (2004).
In this paper, we present excitation maps, integral field spectroscopy and an abundance analysis of these four PNe. The observations and data reduction are described in Section 2, while the emission-line maps are presented and discussed in Section 3. In Section 4, we use the spectrophotometry to derive the physical conditions, ionic and elemental abundances, and excitation class determinations. Section 5, gives kinematical signatures such as the expansion and radial velocities and the distances. In Section 6, we provide a discussion of the classification of the CS in M3-6, and our conclusions are given in Section 7.
|Nebula name||PNG number||No. of||PA ()||Exposure||Date||Airmass||Standard Star|
|M3-4||PN G241.0+02.3||3||90||600||31/3/2013||1.01||HD 111980 & HD 031128|
|M3-6||PN G253.9+05.7||6||0||300||31/3/2013||1.01||HD 111980 & HD 031128|
|3||0||50||31/3/2013||1.00||HD 111980 & HD 031128|
|Hen2-29||PN G275.8-02.9||3||45||600||01/4/2013||1.23||HD 111980 & HD 031128|
|1||45||300||01/4/2013||1.23||HD 111980 & HD 031128|
|Hen2-37||PN G274.6+03.5||3||90||600||31/3/2013||1.06||HD 111980 & HD 031128|
2 Observations & data reduction
The integral field spectra of the PNe were obtained over two nights of March 31 and April 01, 2013 with the WiFeS instrument. This instrument delivers a field of view of 25″ 38″at a spatial resolution of either 1.0″ 0.5″or 1.0″ 1.0″, depending on the binning on the CCD. The design of the image slicer means that there is essentially no clear space between pixels – the spatial filling is %. In these observations, we operated in the binned 1.0″x 1.0″mode. The data cover the blue spectral range of 3400-5700 Å at a spectral resolution of that corresponds to a full width at half maximum (FWHM) of km/s (Å), while in the red spectral range of 5500-7000 Å we used the higher spectral resolution grating corresponding to a FWHM of km/s (Å).
The wavelength scale was calibrated using the Cu-Ar arc Lamp with
40s exposures throughout the night, while flux calibration was
performed using the STIS spectrophotometric standard stars HD 111980
& HD 031128 111Available at :
www.mso.anu.edu.au/bessell/FTP/Bohlin2013/GO12813.html. In addition, a B-type telluric standard HIP 38858 was observed. Telluric absorption features from atmospheric oxygen (O) and water (HO) are corrected with PyWiFeS 222http://www.mso.anu.edu.au/pywifes/doku.php. data reduction pipeline (Childress et al. (2014)) as follows. First, the absorption for each telluric standard is measured by fitting a smooth low-order polynomial (typically a cubic function) to the stellar continuum redward of 6000Å, and dividing the observed spectrum by this smooth continuum fit. Then, the effective mean absorption at zenith for each component (O and HO) is computed for each observed telluric standard by scaling the observed absorption to airmass 1 using the appropriate scaling of optical depth with airmass for each component. We note that the O bands are typically saturated, while the HO bands are not, so they have different airmass dependences. In addition, the relative humidity in the various atmospheric layers ensures that the HO absorption components vary relative to the O obsorption. The mean zenith absorption profiles for the two components (O and HO) can then be similarly scaled the airmass of any observed science field, and the WiFeS data cube is divided by the expected telluric absorption profile. A similar technique can be employed with single telluric standards whose absorption profiles are used to correct an individual science data cube.
All data cubes were reduced using the PyWiFeS. A summary of the spectroscopic observations is given in Table 1. In some objects, due to the saturation of strong nebular emission lines such as [O III] 5007 and H, the fluxes in these lines were derived from additional frames with short exposure times.
The global spectra of each of the objects were extracted from their respective data cubes using a circular aperture matching the observed extent of the bright region of the PNe using QFitsView v3.1 rev.741333QFitsView v3.1 is a FITS file viewer using the QT widget library and was developed at the Max Planck Institute for Extraterrestrial Physics by Thomas Ott.. The line fluxes measured in blue and red spectra were slightly re-scaled (by a factor of ) using the emission lines in the overlapping spectral range (5500-5700 Å). This procedure allows for sub-arcsec. differences in the extraction apertures caused by differential atmospheric dispersion.
Emission-line fluxes and their uncertainties were measured, from the final combined, flux-calibrated blue and red spectra, using the IRAF splot task444IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy (AURA) under a cooperative agreement with the National Science Foundation.. Each line was fit with multi-gaussians, as necessary.
3 Nebular Morphology
Emission-Line maps of the brightest lines were constructed using the reduced WiFeS data cubes for each nebula in our set. These maps are used to describe their internal flux distributions and their excitation structure, by combining three lines into an RGB image. We have selected three such images for inclusion here, and these are presented for each nebula in Figures 1 to 4.
First, a combination of H, He I and He II (left panels in Figures 1 to 4) provides the emissivity distribution in the principal recombination lines. Since H is distributed throughout the nebula, in regions where helium is not ionised, these maps will be dominated by the H flux, and will appear red. In regions where Helium is singly ionised, the map will appear yellow (R+G), while in regions in which the He II line is seen, the map will appear mauve (R+B).
The second combination of the lines [O I] , [O II] and [O III] – the central panels in Figures 1 to 4 – brings out the ionisation stratification very clearly. Finally, as a different way of bringing out the ionisation structure, we give the [O I] , [N II] and He II image (right panels in Figures 1 to 4). This combination covers a more extreme range in ionisation potential than the central panels.
These four PNe have very different morphologies. M3-4 (Figure 1) has a very complex bipolar double-shell structure with two lobes in which He II is strong. The appearance of the inner bright rings suggest that they represent an incomplete bipolar shell. The ionisation stratification across these rings is clearly evident. This double-ring structure is embedded in an outer shell dominated by the [O I] emission. The outer shell has similar morphology to the inner, but it extends beyond the field of view of this WiFeS frame. The green tones in the upper left and right quadrants of the central panel shows that the outer shell is of intermediate excitation with strong [O II] emission.
M3-6 (Figure 2) appears as elliptical PN with brightness increasing towards the centre. Overall, the nebula is of lower excitation than M3-4, and the He II emission has its origin in the central star. It shows pronounced “ansae” along PA. At the ends of these ansae are two bright knots very prominent in [O I] and [S II] . These are almost certainly Fast Low-Excitation Emission Line Regions (FLIERs; Balick et al. (1987); Balick et al. (1993, 1994)). These regions are almost certainly shock-excited by a convergent fast stellar wind flow shaped by the density and/or magnetic structure of the earlier AGB wind (Dopita, 1997; Steffen et al., 2002). However, the radial velocities of these two knots are very little different – we measure the northern knot to have a heliocentric radial velocity of km s and the southern one at km s. If our hypothesis on their origin is correct, the knots must lie almost in the plane of the sky. In Table 7, we determine an overall heliocentric velocity of for the nebula which supports our interpretation.
At very low flux levels, we also detected a faint extended halo surrounding the bright nebula and extending over a region of 14″x 20″. This may be the ionised trace of the earlier AGB wind.
Hen2-29, Figure 3 reveals ”a reverse S-shape” or two-arm spiral structure formed by two point-symmetric arcs present along the minor axis within an overall elliptical morphology. The densest regions in the arms are also optically thick, while the overall elliptical outline is filled with high-excitation gas.
Finally, Hen2-37 (Figure 4) displays a like morphology overall. The bar of the contains optically thick material, and there are isolated lower recitation knots embedded in the body of the nebula in a region of otherwise low density. The high excitation regions are confined to a region within the outer circle of the , and is bounded at its outer boundary by lower excitation gas. Here, we might interpret the bar of the as a remnant of a dense equatorial ring of martial ejected during the AGB phase, the knots as being photevaporating inclusions embedded in the fast stellar wind region, and the outer ring as being the swept-up shell of denser AGB wind gas.
4 Determination of physical conditions and abundances
4.1 Line intensities and reddening corrections
The global emission line spectra are summarised in Table 2. These have been analysed using the Nebular Empirical Abundance Tool (NEAT 555The code, documentation, and atomic data are freely available at http://www.nebulousresearch.org/codes/neat/.; Wesson et al. (2012)). The code was applied to derive the interstellar extinction coefficient and subsequent nebular parameters such as temperature, density and ionic and elemental abundances. The line intensities have been corrected for extinction by adopting the extinction law of Howarth (1983), where the amount of interstellar extinction was determined from the ratios of hydrogen Balmer lines, in an iterative method. Firstly, an initial value for the extinction coefficient, , was calculated by the code. This value was derived from the intrinsic , , line ratios, assuming an electron temperature () of 10000K and electron density () of 1000 cm. Secondly, the code calculates the electron temperature and density as explained in the coming section. Thirdly, the intrinsic Balmer line ratios were re-calculated at the appropriate temperature and density, and again was recalculated. No spatial reddening variation was detected throughout any of the PN set, except in the case of M3-6. The two low-ionization condensations which lie along its major axis show a lower extinction compared to the main nebula.
Table 3, lists the estimated reddening coefficients and the observed and fluxes on a log scale. The complete list of the observed and de-reddened line strengths are given in Table 2. Here, columns (1) and (2) give the laboratory wavelengths () and the line identifications of observed emission lines respectively. Columns 3-14 give the observed wavelength (), fluxes F() and fluxes corrected for reddening I() relative to . The Monte Carlo technique was used by NEAT to propagate the statistical uncertainties from the line flux measurements through to the derived abundances. Both line ratios [O III] 5007/4959 and [N II] 6584/6548 were measured to be in the range of (2.9-3.1) which are comparable to the theoretical predictions of 2.98 (Storey & Zeippen, 2000) and 2.94 (Mendoza, 1983), respectively.
|4861.33||HI 4-2||4861.79||100.03.00||100.00.00||4861.55||100.05.00||100.00.00||4860.70||100.0 5.00||100.00.00||4860.60||100.05.00||100.00.00|
|4958.91||[O III]||4959.38||403.312.10||398.216.63||4959.11||232.211.61||227.315.95||4958.30||566.6 28.3||540.837.42||4958.20||722.836.14||695.548.36|
4.2 Optical thickness estimates
The low ionisation [N I] doublet and is present in the spectra of M3-4, Hen2-29 and Hen2-37. It is known that this line and other low ionization species such as [O I] arise in the outermost part of the nebula. Due to the charge-transfer reaction rate, it is expected that both [O I] and [N I] are abundant in the warm transition region between the neutral and ionized nebular envelopes (Liu, X-W et al. 1995). If shocks are present, the intensity of these lines is increased thanks to the increase in collision strength with electron temperature. In the case of M3-4 nebula, both [N I] and [O I] are present in the external parts of the object. By contrast, in Hen2-37 we found traces of [N I] and [O I] in the central region of the object (see Figure 4, central panel), showing that the bar of the is composed of material which is optically-thick to the radiation field. In the case of Hen2-29, the [N I] gas appears distributed in several regions of high emissivity (see Figure 3). Using [N I] line ratio /, we estimate ionic densities of 378cm and 412cm for M3-4 and Hen2-29, respectively. These estimates are in good agreement with the other nebular density measurements listed in Table 4.
In the case of Hen2-37, we can crudely estimate an electron temperature, K, from the [O I] line ratio for Hen2-37 nebula. This value is higher than the temperatures derived from [O III] and [N II] lines (Table 4), and may indicate the presence of a shock contribution.
The strength of [O II] and [N II] emission lines were used by Kaler & Jacoby (1989) to help define the optical thickness of PNe. Following their criteria, we verified that Hen2-37 nebula is probably optically thick to the ionising radiation field (radiation bounded). It displays strong low excitation lines, with F([N II] 6583)F(H) and F([O II] )/F(H) 1.5. By contrast, the nebula M 3-6 is optically thin (density bounded), except within the ansae. This is supported by several lines of evidence: (1) The weakness of the low excitation lines [O I], [O II] and [N II]; (2) A line ratio F([N II])/F(H; (3) The absence of detectable [N I] at 5200. The other two nebulae in the sample appear partially optically thick PNe. M3-4 and Hen2-29 have F[N II]()/F(H, F[O II](3727)F(H, F[N II](6583)F(H and F[O II](3727)F(H, respectively.
|Object||Log F||Log F|
|This work||Other works||This work||Other works||This work||Other works|
|M3-4||0.25||0.2, 0.38, 0.70, 0.23||-12.43||-12.4, -12.60||-11.49||-11.46|
|M3-6||0.43||0.64, 0.35, 0.58, 0.49, 0.53||-11.01||-11.1, -11.22, -12.0||-10.41||-10.36|
|Hen2-29||0.87||1.01, 0.63||-12.17||-12.1, -12.17||-11.44||-11.40|
|Hen2-37||0.74||0.80, 0.69, 1.03||-12.28||-12.4, -12.73, -12.40||-11.60||-11.53|
|Object||Temperature (K)||Density (cm)|
|[O III]||[N II]||[S II]||[Ar V]||[S II]||[O II]||[Ar IV]||[Cl III]|
4.3 Temperatures and densities
The observed emission lines in our sample (Table 2), permitted us to estimate the electron temperature and density from both the low and the medium ionisation zones. The low ionisation species provide the nebular temperature from the [N II] ()/5754 line ratio and the density from the [S II] 6716/6731 and the [O II] 3727/3729 line ratios. The moderately excited species provide the nebular temperature from the [O III] ()/4363 line ratio and the density from the [Cl III] 5517/5537 and [Ar IV] 4711/4740 line ratios. The temperatures derived from [O III] lines are consistently higher than those derived from the [N II] lines, except in the case of the optically-thin object M3-6.
In Table 4, we list the calculated nebular temperatures, densities and their uncertainties for each object. In addition, it also provides comparisons of our results with those which have previously appeared in the literature. In almost cases, we find good agreement with other works. As expected, temperatures derived from [O III] lines are consistently higher than those derived from the [N II] lines, except in the case of the optically-thin object M3-6.
4.4 Ionic and elemental abundances
Applying the NEAT, ionic abundances of Nitrogen, Oxygen, Neon, Argon, Chlorine and Sulfur were calculated from the collisional excitated lines (CEL), while Helium and Carbon were calculated from the optical recombination lines (ORL) using the temperature and density appropriate to their ionization potential. When several lines from a given ion are present, the adopted ionic abundance was taken averaging the abundances from each line weighted according to their line intensities. The total elemental abundances were calculated from ionic abundances using the ionization correction factors (ICF) given by Kingsburgh & Barlow (1994) to correct for unseen ions. The helium elemental abundances for all objects were determined from the He/H and He/H ions, assuming ICF (He) = 1.0. The carbon elemental abundances were determined from C/H for all objects, except in the case of Hen2-29 where we used both C/H and C/H taking ICF (C) = 1.0. In general, the chemical abundances are found to be in good agreement with other works which presented in Tables 5 and 6.
|M 3-4||M 3-6|
|Element||Neat(1)||Neat (2)||Ref 1||Ref 2||Ref 3||Ref 4||Neat(1)||Neat (2)||Ref 1||Ref 2||Ref 3||Ref 4|
|Element||Neat(1)||Neat (2)||Ref 1||Ref 2||Ref 3||Neat(1)||Neat (2)||Ref 1||Ref 2||Ref 3||Ref 5||Ref 6|
4.5 Excitation classes and Peimbert classification
Our objects (Table 2), display a mixture of both low- and high-excitation emission lines. Furthermore, the highly ionised species [Ne V] appears in M3-4, Hen2-29 and Hen2-37. The line strength of He II 4686Å relative to provides a best quantitative measure of the nebular excitation class (EC). The absence of the He II line from the nebular spectrum restricts the PNe to be of low excitation class (EC 5), and where it appears (EC 5), its strength relative to H can be used to define the excitation class. To derive the EC of our PN set, we were followed the methodology of Meatheringham & Dopita (1991) and Reid & Parker (2010). Both methods use the same scheme to derive the EC of low excitation PNe, and therefore give the same result. For PNe of EC, Reid & Parker (2010) have upgrade the scheme of Meatheringham & Dopita (1991) by considering also the increase of [O III] line ratio with nebular excitation. This line ratio was fixed in the scheme of Meatheringham & Dopita (1991), since in their optically-thick models this line ratio ‘saturates’. However, many PNe are optically thin, and in this case, the [O III] line ratio initially increases with decreasing optical thickness, before finally decreasing for high-excitation objects when the nebula becomes too optically-thin. Our results (Table 6) reveal that M3-6 is of low EC, while other three objects are of medium to high EC.
Hen2-29 is the only PN in our set which shows excess in He and N abundances. Thus, it can be classified as a Peimbert type I by applying the two criteria, He/H and N/O , as proposed by Maciel & Quireza (1999). These criteria are much rigid than the original criteria, He/H or N/O , suggested by Peimbert & Torres-Peimbert (1983). Based on the analysis of 68 PNe, Kingsburgh & Barlow (1994) found that the average helium abundance of Type I PNe is increased by a small amount (a factor of 1.2) over non-Type I objects. Therefore, they distinguish between Type I and non-Type according to the N/O ratio only. They defined Type I PNe as those PNe with N/O . Following this criterion, we classify all four PNe in our sample as non-Type I objects.
The Type I PNe were believed to have evolved from the most massive progenitor stars with initial main sequence masses range 2.4-8.0 M (Quireza et al., 2007) and consequently they should be associated with the thin Galactic disk and have low velocity dispersion. They also show a wide range of ionisation structures. Using the adopted distance and measured radial velocity, we derived a vertical Galactic height pc and a peculiar velocity666The difference between the observed radial velocity corrected for local standard of rest (LSR) and the velocity determined from the Galactic rotation curve. kms for Hen2-29. Previous measurements of the spatial ( pc, Gilmore & Reid (1983)) and kinematic characteristics ( kms, Maciel & Dutra (1992)) of this PNe show that it belongs to the thin Galactic disk population.
Faundez-Abans & Maciel (1987) developed the Peimbert classification scheme by dividing the Peimbert Type II into Type IIa and Type IIb. This division was proposed based essentially on the nitrogen abundances. They show that Type IIa are of nitrogen abundances (log (N/H)+12 ) larger than that of Type IIb (log (N/H)+12 ). Following the further re-analysis of Peimbert types suggested by Quireza et al. (2007), we find both M3-4 and Hen2-37 are consistent with being of Type IIa ( and N/O , but He/H ). Therefore, they should arise from an older and less massive main sequence population (1.2-2.4 M) than the Type I objects. Furthermore, the spatial and kinematic properties of M3-4 (336 pc and 29.9 km s) and Hen2-37(185 pc and 7.3 km s) show that they are thin Galactic disk members. M 3-6 nebula is the only PN in the set that shows (). The object lies at pc and has km s, hence it can be classified as a Type IIb PN.
5 Kinematical parameters and distances
|Object||(km S)||V (km S)||Distance (kpc)||Excitation class|
|Ours||literature||[S II]||[N II]||(I)||(II)||(I)||(II)|
|M3-4||25.03.24||7440 (1), 35.44.6(3)||18.7||17.5||3.82||4.27||6.2||7.8|
In order to detemine the evolutionary status of a PN, we need its expansion velocity () and its distance. Here, we determined from the two emission lines [S II] and [N II] which lie in the red channel of WiFeS instrument with the higher spectral resolution (). The FWHM of each line was measured using the IRAF splot task. The expansion velocity was corrected for instrumental and thermal broadening following (Gieseking et al., 1986). The results were listed in Table 7.
Hen2-37 has the highest expansion velocity, and in this object double peak nebular emission lines are clearly seen. For Hen2-29 we derive – greater than the standard value of 20 kms (Weinberger, 1989). Meatheringham et al. (1988) report expansion velocities of 23.6 kms and 30.7 kms for Hen2-29 and Hen2-37, respectively. These are smaller than we determine, but this is to be expected since the Meatheringham et al. (1988) results rely on the [O III] emission line which has higher ionisation potential than the lines we used, and is therefore produced closer to the centre of the nebula - see the central panels Figures 3 and 4.
The systemic velocities of the sample were measured using the emsao task of the IRAF package. The were derived from the Doppler-shift of [N II] 6548Å, 6562Å, [N II] 6583Å, He I 6678Å, [S II] 6716Å and [S II] 6730Å emission lines. We select these lines because they are locate in the high spectral resolution part of nebular spectra (). To derive the heliocentric radial velocity , we used the IRAF task RVCorr to correct for the effect of the Earth’s motion around the Sun. The results were presented in Table 7 and compared with the work of (Schneider & Terzian (1983), hereafter STPP83), Meatheringham et al. (1988) and Kniazev (2012). In general the determined of M3-4 nebula is relatively close to the recent estimation of Kniazev (2012) but it is smaller than of STPP83. The value of STPP83 was originally taken from the low dispersion spectra of “Mayall-1964” given as a private communication with Perek and Kohoutek (1967). The for M3-6 differs significantly in accuracy with STPP83. The value of STPP83 was taken as a weighted average of two unpublished measurements 5711 (Minkowski 1957) and 1225 (Mayall 1964) as a private communication with Perek & Kohoutek (1967). This average value was weighted according to the error of each measurement. For the object Hen2-29, our derived is consistent with that of MWF88 within the error range. For Hen2-37, we determined of -12.53.12 which is differs significantly than that of Meatheringham et al. (1988). Our value was checked by measuring = -16.05.0, from the blue part of the nebular spectrum (which has lower resolution, , than the red). Three possible explanations can be provided for this discrepancy: (1) The different spectrum resolution particularly most of nebular lines appear of double peaks due to the high expansion velocity of the nebula; (2) The radial velocity determined here was measured from integrated spectrum over the whole object; (3) The negative sign of the radial velocity measure given by Meatheringham et al. (1988) was simply missed.
None of the sample has a distance determined from either the trigonometric, spectroscopic, cluster membership, or expansion methods. Therefore, we must rely on the statistical approaches to estimate distances, recognising the large errors that this entails. We adopt here the average distance derived from the recent two distance scales of Ali et al. (2015a) and Frew et al. (2016), for each PN. The Ali et al. (2015a) scale depends on the mass-radius and radio surface brightness temperature-radius empirical relationships, and specifically on the nebular angular size and 5 GHz radio flux. The Frew et al. (2016) scale depends on the empirical relationship between H surface brightness and the radius of PN, using the nebular angular size and the H flux. The results of the two approaches are given in Table 6. Both methods give roughly the same distances. The angular radii of the PNe are taken from Frew et al. (2016), while radio fluxes at 5GHz were taken from Cahn et al. (1992) except for M3-4. We find this object has radio fluxes which differ between the different references. We adopted here the average value from Milne & Aller (1975), Milne (1979), Zijlstra et al. (1989) and Cahn et al. (1992).
6 The central star of M3-6
Tylenda et al. (1993), Marcolino & de Araújo (2003) and Weidmann & Gamen (2011) classified the CS of the M3-6 nebula as being of the WELS type. Miszalski (2009) reported that many of WELS are probably misclassified close binaries. Further Miszalski et al. (2011) and Corradi et al. (2011) observed many of WELS emission lines in the spectra of stars known to be close binary systems, and explained that these lines were originated from the irradiated zone on the side of the companion facing the primary.
Basurah et al. (2016) claimed that, for many objects, the WELS class of the PNe central star may be spurious. From WiFeS data of NGC 5979, M4-2 and My60, they showed that the characteristic CS recombination lines of WELS type are of nebular origin. Specifically, the lines used to provide the WELS classification are CII at 4267Å, NIII at 4634Å, 4641Å, CIII at 4650Å, CIV at 5801Å and 5812Å. The CS classification and the WELs class in general were discussed in detail by Basurah et al. (2016). Here, we have found a further example that increases the doubt regarding the reliability of the WELS classification. In Figure 5 we show emission line maps of M3-6 in four CS emission lines, CII, CIII, CIV, and NIII used to define the WELS class. It is apparent that the emission of these lines are spatial distributed over a large area of the nebula, and hence are clearly of nebular rather than of CS origin. Only in CIV at 5801Å and 5812Å is the emission limited to the central nebular region and consequently probably originates on the CS.
In fact, we can provide a revised classification for this central star based upon our own data. In Figure 6 we show our complete extracted spectrum of the central star of M3-6. To obtain this, we carefully removed the nebular emission determined from a zone around the CS from the spaxels defining the continuum image of the CS in the image cube. Many Balmer and He II lines clearly visible in absorption in the blue spectrum (upper panel), except He II at 4686 appears in emission. Further, many other emission lines are present such as N IV 4058, Si IV 4089, 4116, 4654, C IV 4658, and N III 4634, 4640. The N V doublet and the O V line at 5114 both appear in absorption. In the red spectrum (lower panel), the C IV doublet and the N IV doublet are clearly seen in emission, while He II 5412 and the interstellar lines of Na I D are visible in absorption.
Following the CS classification scheme of Mendez (1991) and improvements given to this scheme by Weidmann et al. (2015), we classify the CS of M3-6 nebula as a H-rich star of spectral type O3 I(f*). The relative weakness of the N V doublet compared to the He II line is an indication of a spectral type O3, and the presence of N III the doublet along with the He II line are features of Of(H) type. Further, the stronger emission of N IV relative to N III 4640 provides the * qualifier which is a unique property stars of the O3 type (Walborn & Howarth, 2000). In general, the Si IV emission is also present in Of* spectra. From Figure 6, it appears that the C IV recombination lines at are of CS origin – as is also revealed in the emission line map of the nebula constructed in these two lines and shown in Figure 5.
In this paper we presented the first integral field spectroscopy of the southern planetary nebulae M3-4, M3-6, Hen2-29 and Hen2-37 in the optical range 3400-7000 Å. We demonstrated the utility of these observations in both providing narrow-band data to probe the morphological and excitation structure of the nebulae as well as in deriving their dynamical nature, their optical thickness and their chemical enrichment characteristics.
The four PNe have different optical thickness, where M3-6 is optically thin, Hen2-37 is optically thick while M3-4 and Hen2-29 are partially optically thick. From the strength of He II line, we derived excitation class 6.2-7.8, 4.6, 8.0-10.5, and 8.4-11.4 for M3-4, M3-6, Hen2-29 and Hen2-37, respectively. From the chemical analysis of the sample based upon integral spectroscopy, we provided Peimbert types I, IIa, IIa, IIb for Hen2-29, M3-4, Hen2-37 and M3-6, respectively, and noted that the long-slit spectroscopic data can provide surprisingly good results even though only a sub-region of the PN is analysed by this technique.
In the case of M3-6, we find that the majority of the recombination lines used in literature to classify the CS as a weak emission-line star are in fact of nebular origin. Instead we classify the central star as H-rich and of spectral type O3 I(f*). This result extends to five (M3-6 and NGC 3211, NGC 5979, My 60, M 4-2 that mentioned in Basurah et al. (2016)) the number of cases of mis-classification of WELs stars discovered using integral field spectroscopy and served to increase doubts regarding the reliability of the WELS classification in general. In this we support the conclusions of Weidmann et al. (2015) who observed 19 objects with the WELS classification amongst the total of 72 so classified. They concluded that ”the denomination WELS should not be taken as a spectral type, because, as a WELS is based on low-resolution spectra, it cannot provide enough information about the photospheric H abundance.
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