Twenty years of observations of PM 1-188: Its chemical abundances and extraordinary kinematics
Miriam Peña, Liliana Hernández-Martínez, Francisco Ruiz-Escobedo
MMNRAS , 1–15 (2015) Preprint 22 February 2021 Compiled using MNRAS L A TEX style file v3.0
Twenty years of observations of PM 1-188: Its chemicalabundances and extraordinary kinematics
Miriam Peña, (cid:63) Liliana Hernández-Martínez, , Francisco Ruiz-Escobedo Instituto de Astronomía, Universidad Nacional Autónoma de México, Ap. 70-264, Ciudad Universitaria, 04510, Ciudad de México Instituto de Ciencias Nucleares, Universidad Nacional Autonoma de México, Ap. 70-543, 04510, Ciudad de México Facultad de Ciencias, Universidad Nacional Autonoma de México, 04510, Ciudad de México
Accepted 2021 February 17. Received 2021 February 15; in original form 2020 September 14
ABSTRACT
The analysis of 20 years of spectrophotometric data of the double shell planetarynebula PM 1-188 is presented, aiming to determine the time evolution of the emissionlines and the physical conditions of the nebula, as a consequence of the systematicfading of its [WC 10] central star whose brightness has declined by about 10 mag inthe past 40 years. Our main results include that the [O iii ], [O ii ], [N ii ] line intensitiesare increasing with time in the inner nebula as a consequence of an increase in electrontemperature from 11,000 K in 2005 to more than 14,000 K in 2018, due to shocks.The intensity of the same lines are decreasing in the outer nebula, due to a decrease intemperature, from 13,000 K to 7,000 K, in the same period. The chemical compositionof the inner and outer shells was derived and they are similar. Both nebulae presentsubsolar O, S and Ar abundances, while they are He, N and Ne rich. For the outernebula the values are 12+log He/H= 11.13 ± ± ± ± ± −
150 to 100 km s − range, and both shellshave expansion velocities of about 40 km s − . Key words: planetary nebulae: individual: PM 1-188 – ISM: abundances – ISM:kinematics and dynamics
Planetary nebulae (PNe) are formed from stars of low-intermediate masses (1–8 M (cid:12) ) in an advanced stage of evolu-tion (post AGB). They have completely burnt hydrogen andhelium in their cores and possess dense carbon-oxygen coresof about 0.6 M (cid:12) . At these late stages, the star loses partof its atmosphere through intense winds. It evolves towardslarger effective temperatures and starts ionising the ejectedenvelope. Thus the PN was part of the stellar atmosphereand its chemical composition depends on the processes ofdredge-ups that carry elements of the interior to the sur-face.In general, central stars of PNe possess a thin H-rich en-velope, but about 15% of them are known to be H-deficient,showing atmospheric instabilities and developing strong stel-lar winds similar to the massive Wolf-Rayet stars. They havebeen named [WR], to differentiate them from the massive (cid:63)
E-mail: [email protected] (MP)
WR. Almost all of these stars are of the sequence of Car-bon, showing strong helium, carbon and oxygen emissionlines, therefore they are designated as [WC] stars. It is foundthat most of them are distributed in the early spectral types[WC 2–4] which show lines of C iv , He ii and O v and haveeffective temperatures larger than 100,000 K, and in the latespectral types [WC 8 – 11], showing lines of lower ions andmuch lower effective temperatures, from 25,000 K to 80,000K (Koesterke 2001; Acker & Neiner 2003). There are veryfew stars in the intermediate classes. In the following theassociated nebula will be called a [WC]PN.PN G012.2+04.9, with common names PM 1-188,HuBi 1 and IRAS 17514-1555, is a well known quite inter-esting planetary nebula first found by the IRAS satellite,which during several decades has shown a central star ofthe rare [WC 10] class. This was discovered by Hu & Bibo(1990) who also reported a strong infrared excess in this ob-ject. Since this discovery, PM1-188 has been actively studiedto reveal its nature and evolution.Pollacco & Hill (1994) found that PM 1-188 consists of © 2015 The Authors a r X i v : . [ a s t r o - ph . S R ] F e b M. Peña et al. a faint extended nebula surrounding a dense bright zone,with apparently bipolar structure. Peña (2005b) reportedthe extraordinary ionisation structure of the nebula, showingthat the [N ii ] and [O iii ] emission lines appear concentratedin the inner zone, while the H lines are more intense in theouter zone, presenting a cavity in the centre (see her Fig.1). Also in this paper, the author indicated the systematicfading of the central star.Peña (2005b) reported that the [WC] central star ofPM 1-188 shows a systematic fading. The WR phenomenonis known to be variable. Indeed, a few [WC] PNe show vari-ations that could be intrinsic to the star or due to externalcauses (Werner et al. 1992; Peña et al. 1997).Leuenhagen & Hamann (1998) derived physical param-eters of the central star of PM 1-188 from non-LTE theo-retical stellar atmosphere models finding a mass loss ratelog ˙M= − − , a surface temperature of 35,000 K, a luminosity logL/L (cid:12) =3.70 and a chemical composition, in mass fraction,of β (He) = 42%, β (C) = 50%, β (O) = 7%, and a very smallamount of H. The remarkable He- and C-rich and H-poorphotospheric chemical composition is typical of [WC] stars(Koesterke 2001) and it has been found in some born-againPNe such as Abell 30, although in this case, β (He) amounts63%, β (C), 20% and β (O), 15%, by mass (Todt et al. 2015),chemistry more typical of a [WC-early] star. PM 1-188 mighthave experienced an unusual evolution as Abell 30.Peña et al. (2001) and Gorny et al. (2001) analysed somenebular and dust properties in a number of [WC]PNe, andthey suggest a possible evolutionary sequence from [WC-late] to [WC-early] objects, except for PM 1-188 and K 2-16,both with a very late [WC] star. Peña et al. (2001) suggestedthat these exceptions could have experienced a late-helium-flash and have returned to the AGB for a born-again evo-lution, or they could be stars evolving very slowly from theAGB due to their low mass.Recently, Guerrero et al. (2018) deeply analysed thenebular structure and the evolution of the central star,claiming that the inner nebula is ionised by shocks withvelocity of about 70 km s − ; the outer nebula is recombin-ing and, according to models of stellar evolution, the centralstar is the descendent of a low-mass star that experienced aborn-again event whose ejecta shock-excite the inner shell.In summary of these prior works, PM1-188 is definitelya [WC 10] PN exhibiting time variations which might berelated to unusual stellar evolution. While the central starseems to be more deeply understood, the nature of the neb-ula remains unclear. To further advance understanding ofPM 1-188 and gain insights on evolution of born-again phe-nomenon, time temporal evolution of the nebular parame-ters and internal gas motion are necessary. For that purpose,we gathered 20 years of low and high-resolution spectro-scopic data and also newly secured high-dispersion spectrato perform spatially-resolved studies of the nebula.This paper is organized as follows: In §2 the observa-tions are described. In §3, line intensities and their time evo-lution in the inner and outer nebulae are analysed. In §4, wepresent the physical conditions (electron temperatures anddensities) and their time evolution, and the derived ionic andchemical abundances for both, inner and outer nebulae. In§5 we discuss the evidence for shocks in the inner shell. Thepresence of an outflow, derived from high-resolution spec- tra, is shown in §6, while in §7 Position-Velocity diagramsconstructed from data obtained in two position angles, arepresented. The discussion and conclusions are displayed in§8. In an appendix we present the Ionisation CorrectionsFactors (ICFs) used to derive the chemical abundances andthe list of atomic parameters employed to determine thephysical conditions and chemistry of the nebula. As said above, more than 20 years of spectroscopic ob-servations of PM 1-188 have been gathered. Data prior to2005 were already published in Peña et al. (2001) and Peña(2005a), and are reanalysed here. New observations reportedin this work include spectrophotometric data obtained in2005 with the Las Campanas Observatory (LCO) Clay tele-scope and the LDSS3-Two spectrograph; observations withthe OAN-SPM 2-m telescope and the Boller & Chivens(BCh) spectrograph in 2005 and 2017, and observations withthe Manchester Echelle Spectrograph (MES) in 2017. Alsoobservations with the Gran Telescopio Canarias (GTC) andOSIRIS spectrograph in long slit mode in 2018 (programGTC112-18A, P.I. M. Guerrero), retrieved from the publicarchives, are analysed here.The log and characteristics of the observations are pre-sented in Table 1 where we list the telescope and instrumentused, the covered wavelength range, the spectral resolution(in Å/pix) and resolving power R ( λ / δλ ), the spatial scale,the total exposure time, the slit size (length and width),the position angle of the slit, P.A., the seeing during theobservations and some comments. Most of the data along these years was obtained at Observa-torio Astronómico Nacional, San Pedro Mártir, B.C., Méx-ico (OAN-SPM), with the 2.1-m telescope and the Boller &Chivens (BCh) spectrograph, at intermediate resolution (4– 7 Å/pix). Most of the time, the observed wavelength rangewas from about 3700 to 7400 Å which allows to obtain alarge number of lines emitted in the visual zone. The spec-tral resolution of these observations is adequate to resolvethe important [N ii ] and [S ii ] lines employed in temperatureand density diagnostic.In two epochs (1997 and 2004) the high resolutionechelle REOSC spectrograph, that achieves a spectral reso-lution of about 0.3 Å/pix (R=18000 at 5000 Å), was used. In2017, the Manchester Echelle Spectrograph (MES) was em-ployed in order to derive the interesting kinematics of thisnebula; these observations are described in the next subsec-tion.In 2005 spectra with the LDSS3-Two long-slit spectro-graph, attached to the LCO 6.5-m telescope Clay were ob-tained. This spectrograph works with two arms operatingsimultaneously. The blue one covered from 3830 to 5150Å and the red one, from 5160 to 6595 Å. The grism VPHBlue was employed. Standard stars and a He-Ne-Ar lampwere used for flux and wavelength calibrations. The usualprocedures with IRAF were performed for data reduction. IRAF is distributed by the National Optical Astronomy Ob-MNRAS , 1–15 (2015) Table 1.
Log of observations and their characteristicsDate telescope instrument λ range spectral resolution plate scale total exp slit size P.A., seeing, commentsÅ Å/pix R (5000Å) arcsec/pix time (s)1997/08/04 a OAN-2m echelle REOSC 3600–6800 0.3 18000 1.20 1800 13.3” ×
2” 90 ◦ , 1.5”, Peña et al. (2001)2000/09/25-26 a OAN-2m BCh-600 l/mm b (cid:48) ×
4” 90 ◦ , 2”, Peña (2005a)2002/06/05-07 a OAN-2m BCh-300 l/mm 3450–7400 2.23 567 1.50 4200 5’ ×
4” 90 ◦ , 1.5”, Peña (2005a)2004/04/23-26 a OAN-2m echelle-REOSC 3800–6800 0.3 18000 1.20 2700 13.3” ×
2” 90 ◦ , 2”, Peña (2005a)2005/05/15-17 OAN-2m BCh-600 l/mm 3800–5950 1.54 685 1.20 7200 5’ ×
4” 90 ◦ , 1.5”,2005/05/18-19 OAN-2m BCh-600 l/mm 5200–7300 1.54 685 1.20 3600 5’ ×
4” 90 ◦ , 1.5”2005/08/13 LCO Clay LDSS3-Two 3830–5150 0.8 1900 0.38 900 4’ ×
2” 173 ◦ (parallactic), 1”-1.2”” ” ” 5160–6595 0.8 1900 0.38 900 4’ ×
2” 173 ◦ (parallactic), 1”-1.2”2017/06/24,25 OAN-2m BCh-300 l/mm 3700–7360 2.23 685 1.08 7200 5’ ×
2” 90 ◦ , 1.5”2017/06/29 OAN-2m MES order-87 6545–6595 0.06 25000 0.35 7200 5.2 (cid:48) ×
2” 90 ◦ , 1.5”2017/06/30 OAN-2m MES order-87 6545–6595 0.06 25000 0.35 7200 5.2 (cid:48) ×
2” 0 ◦ , 2.5”2018/05/14 GTC c OSIRIS LS 3630–7500 2.12 1018 0.254 6300 6.5 (cid:48) × ◦ , 1”a: published in Peña et al. (2001) or in Peña (2005a)b: For the BCh spectrograph, the used grid is indicatedc: GTC112-18A program, P.I. M. Guerrero, data retrieved from the public archives The nebular line fluxes were measured in the calibrated spec-tra, integrating all the emission over a local continuum esti-mated by eye.The central star is not visible in the spectra of 2017 or2018 due to its faintness. But it was well detected in 2005,when data with LCO telescope Clay were obtained. Fig. 1shows a zone around H α , where the stellar continuum andthe extended nebula are visible. Nebular [N ii ] lines are veryintense in the inner zone, [N ii ] λ ii λ ii ] linesare faint in the outer zone, but perfectly detected. H α ap-pears intense in the centre, overlapped to the stellar con-tinuum, then it shows a cavity immediately outside and in-creases its emission further away in the zone of the outernebula. The nebular lines are marked in the figure. The stel-lar emission shows several WR lines of C ii , also marked inthe figure.The most recent observations presented here were ob-tained in 2018 with the GTC OSIRIS spectrograph in longslit mode. Data were retrieved from the archives and datareduction was performed with the standard procedures us-ing IRAF. Flux and wavelength calibrations were carriedout with the standard star GD140 and the Hg-Ar-Ne lampprovided in the same observations. These spectra covereda wide wavelength range from 3700 to 7300 Å. The largeexposure time (105 min in total, distributed in 4 individualobservations of 1575 s each) helped to detect a large num-ber of nebular lines, which allowed us to determine confi-dent physical conditions and ionic abundances. Nebular linefluxes were measured in the calibrated spectra, integratingall the emission over a local continuum estimated by eye.All the spectra of a given epoch were combined in asingle spectrum to increase the signal-to-noise. The totalobserving time for each epoch is listed in Table 1. For theanalysis of lines, we made an effort to obtain separatelythe line intensities of the inner and outer nebulae in allthe spectra, although in the epochs previous to 2005, the servatories, which is operated the Association of Universities forResearch in Astronomy, Inc., under contract to the National Sci-ence Foundation. [WC] central star was very bright and some of its emissionlines are contaminating some nebular lines such as the He i λλ ii ] λ + ( λλ + ( λλ ++ ( λλ ++ ( λ ++ ( λ + ( λλ ++ ( λ ( λ ( λ (cid:48)(cid:48) , and all but [O ii ] λ α , H β , He i i lines are mainly concentrated in the outer nebula, and theirdistribution show a central cavity of about 2.5” in the zonewhere low-ionisation lines are intense.Calibrated extracted spectra for both nebulae, as ob-tained from GTC OSIRIS data, are presented in Fig. 2,where the different behaviour of [N ii ] λλ α , in the inner and outer nebula, is noticeable. High resolution spectra were obtained with the OAN 2.1mtelescope and the Manchester Echelle Spectrometer (MES,Meaburn et al. 1984; Meaburn et al. 2003) on 2017-06-29.Four spectra of 1800 s exposure time each were acquired,with the slit oriented E-W. These spectra were combined inone of 7200 s of total exposure time. The slit was 2 (cid:48)(cid:48) wideand 6.5 (cid:48) length. A δλ = 90 Å bandwidth filter was used toisolate the order 87, covering the interval from 6545 Å to6595 Å, which then includes H α and [N ii ] 6548 Å and 6583Å lines. A binning of 2 × MNRAS , 1–15 (2015)
M. Peña et al.
Figure 1.
The nebular lines H α λ ii ] λλ ◦ . The spatial scale is indicated in theleft side, with the origin in the stellar continuum. The extendednebula is clearly visible in these lines. The compact inner nebulapresents intense [N ii ] lines, diminishing towards the outer nebula,while H α is more intense in the outer nebula. The wavelengthsof nebular lines are marked in black. The stellar continuum andseveral [WR] lines of C ii are detected and marked in red. Inparticular, the C ii λ ii ] λ resolution of 0.06 Å/pix, equivalent to 11 km s − , and aspatial resolution of 0.35 (cid:48)(cid:48) /pix. The wavelength calibrationwas performed with a Th-Ar lamp, no flux calibration wasdone.Another spectrum was obtained with MES on 2017-06-30, with total exposure time of 3600 s, and the slit orientedN-S (P.A. 0 ◦ ). On this and the other occasions the slit wascentred in the central star position. Emission of the inner nebula was extracted from the twocentral arcsecs. Fluxes measured for this zone, for the dif-ferent epochs, relative to H β , not corrected by reddening,F( λ ), are presented in Table 2. The latter two lines in thistable present, for each epoch, the total flux measured for H β through the slit and the logarithmic reddening correction,c(H β ) with its uncertainties. For each observation, the log-arithmic reddening correction was derived by using Cardelliet al. (1989) reddening law by assuming a ratio of total toselective extinction R V = 3 . and by using the theoreticalH α /H β ratios given by Storey & Hummer (1995) for tem-peratures of 7,500 K, 10,000 K, 12,500 K and 15,000 K, ade-quate to the value derived for each observation. For all cases,case B recombination theory and a density of 100 cm − wereassumed. In the calculus of c(H β ) we did not take into ac-count the minimal contamination of He ii λ λ β and H α which in both cases amountless than 0.5%, for the inner nebula.The observed fluxes were dereddened by using the men-tioned reddening law, and the c(H β ) value determined ineach case. The results, relative to I(H β ), are presented inthe same table, under the header I( λ ). The uncertainties for I( λ ), calculated by considering the uncertainties in the fluxesand in c(H β ), are presented in %, enclosed in parenthesis.The emission of the outer nebula was extracted from thezones around the central cavity presented by the H lines.Both zones at each side of the centre were extracted andthen combined. The fluxes F( λ ), relative to H β , not cor-rected by reddening, are presented in Table 3. The same asin the case of the inner nebula, fluxes were dereddened withCardelli et al. (1989) reddening law and c(H β ) derived ineach case. These intensities, relative to I(H β ), and their re-spective uncertainties, in percentage, are listed in the sametable. The latter two lines of the table indicate the observedH β fluxes and the derived logarithmic reddening correctionsc(H β ) with their uncertainties.The analysis of the line intensities in Tables 2 and 3indicates that prior to 2005, the He ii λ λλ ii linebrights alone with an intensity of about 0.1 H β in 2017 -2018. It could be that the line intensity has been slightlyincreasing with time from 0.10 in 2005 to 0.13 in 2018, rel-ative to H β , but the uncertainties (of about 15 % for GTCOSIRIS value and larger for OAN-BCh and LCO-LSST2values) make this not conclusive. In the outer nebula He ii λ ii ], [N ii ], and [S ii areintense in the inner nebula. In particular, the [N ii ] λ α , indicating that thenebular [N ii ] λ λ α . This is indicativeof possible shocks as it will be discussed in §5.The [O iii ] λ β , as expected due to thisnebula is ionised by a 35,000 - 38,000 K effective tempera-ture star which does not produce enough photons to ionisea large fraction of O + .In the outer nebula all these low ionisation lines aredetected, but much fainter, for instance the [O ii ] λ β intensity ratio changes from about 7.9 in the inner zone to0.4 in the outer zone, and I([N ii ] λ β ) has a valueof about 2.7 in the inner zone and about 0.4 in the outernebula (both values from the dereddened line observationson 2018). [O iii ] λ iii ] λ iii ] λ iii ] λ ii ] λλ ii ] λ i ] λ ii ] λλ ii ] λλ MNRAS , 1–15 (2015) Wavelenght (Å) F l u x ( × e r g c m s ) GTC OSIRIS PM1-188 (Inner)
Wavelenght (Å) F l u x ( × e r g c m s ) GTC OSIRIS PM1-188 (Inner)
Wavelength (Å) F l u x ( × e r g c m s ) GTC OSIRIS PM1-188 (Outer)
Wavelength (Å) F l u x ( × e r g c m s ) GTC OSIRIS PM1-188 (Outer)
Figure 2.
The calibrated spectra of the inner (up) and outer nebulae (down), obtained from GTC OSIRIS observations, are shown.Details around H α in order to notice the behaviour of [N ii ] lines are shown on the right side. ature in this nebula. The increase in temperature is affectingall the collisionally excited lines, but it is affecting in partic-ular the [N ii ] lines that are very sensitive to electron tem-peratures and shocks. Unfortunately the [O ii ] λλ i lines are intense and theyare affected by blends with stellar lines in the spectra previ-ous to 2005. That is the reason of the large flux of He i λ i λ β ) in both, the inner andouter nebulae. The He i lines have not varied since 2005. Assaid before by (Peña 2005a) He i lines are intense in bothnebulae, indicating a large He abundance.The logarithmic reddening correction, c(H β ), has notchanged with time, with an average value of 1.04 ± ± From the nebular lines, in particular from those collisionallyexcited lines, physical conditions such as electron densitiesand temperatures can be derived from some diagnostic lineratios. In the inner and outer nebulae, densities can be de-termined from the [S ii ] λλ ii ] λλ λ ii λ ii ] λλ ii ] λλ ii ] and [S ii ] diagnostic lines. The results arepresented in Table 4. Densities derived from [S ii ] lines arealways low, but the inner nebula is slightly denser (a fewhundred particles per cm ) than the outer nebula that hasa density of about a hundred particles per cm .Regarding the temperature, the value derived from the[N ii ] λλ / line ratio for the outer nebula is near7,000 K in 2018, and more than 14,000 K in the inner nebula.Both temperatures are higher (but very uncertain) from the MNRAS000
The calibrated spectra of the inner (up) and outer nebulae (down), obtained from GTC OSIRIS observations, are shown.Details around H α in order to notice the behaviour of [N ii ] lines are shown on the right side. ature in this nebula. The increase in temperature is affectingall the collisionally excited lines, but it is affecting in partic-ular the [N ii ] lines that are very sensitive to electron tem-peratures and shocks. Unfortunately the [O ii ] λλ i lines are intense and theyare affected by blends with stellar lines in the spectra previ-ous to 2005. That is the reason of the large flux of He i λ i λ β ) in both, the inner andouter nebulae. The He i lines have not varied since 2005. Assaid before by (Peña 2005a) He i lines are intense in bothnebulae, indicating a large He abundance.The logarithmic reddening correction, c(H β ), has notchanged with time, with an average value of 1.04 ± ± From the nebular lines, in particular from those collisionallyexcited lines, physical conditions such as electron densitiesand temperatures can be derived from some diagnostic lineratios. In the inner and outer nebulae, densities can be de-termined from the [S ii ] λλ ii ] λλ λ ii λ ii ] λλ ii ] λλ ii ] and [S ii ] diagnostic lines. The results arepresented in Table 4. Densities derived from [S ii ] lines arealways low, but the inner nebula is slightly denser (a fewhundred particles per cm ) than the outer nebula that hasa density of about a hundred particles per cm .Regarding the temperature, the value derived from the[N ii ] λλ / line ratio for the outer nebula is near7,000 K in 2018, and more than 14,000 K in the inner nebula.Both temperatures are higher (but very uncertain) from the MNRAS000 , 1–15 (2015)
M. Peña et al.
Year l o g ( I / I ( H )) PM1-188 (Inner)
Year l o g ( I / I ( H )) PM1-188 (Outer) [OII] 3727 [OIII] 5007 [NII] 5755 HeI 5875 [NII] 6548 [SII] 6716 [OII] 7325
Figure 3.
Time evolution of the intensity of most important lines are presented, for the inner and outer nebulae. Line intensities arepresented in log-scale. [S ii ] λλ / (6716 + 31) and [O ii ] λλ ii ] diagnostic ratios, is plotted asa function of time in Fig. 4. Despite the errors that canbe as large as 1,000 K, it is evident an opposite systematicbehavior in both nebulae. In the inner one the electron tem-perature has increased from about 11,000 K in the epoch2000 - 2005 to more than 14,000 K in 2018. According tothe photo-ionisation model including shock excitation, pre-sented by Guerrero et al. (2018), this inner nebula is appar-ently heated by a shock with velocity ∼
70 km s − travelingoutwards the ionised shell. The origin of the temperaturerise in this zone should be the traveling shock and the out-flow expanding at about 150 km s − in this zone (see §6).The rise in electron temperature is producing the increasein line intensities of heavy elements in this zone, affecting inparticular the [N ii ] lines, as shown in Fig. 3.On the opposite, the electron temperature has de-creased in the outer nebula, from about 13,000 K in theepoch 2000-2005 to about 7,000 K in 2017-2018. This phe-nomenon is certainly showing the effect of the fading of thecentral star. Less and less ionising photons have been ar-riving to the external nebula since few decades ago, andas a consequence, the gas is cooling fast, and also recom-bining, but at a much slower rate. Supplementary Fig 4.by Guerrero et al. (2018) shows a simple photo-ionisationmodel evaluating the time evolution of electron temperature,the ionisation fraction of H, and the emissivities of variouslines in a spherical photo-ionised nebula, after the turn off ofthe ionising source. This model predicts that emissivities of[N ii ], [O ii ] and [S ii ] lines decrease slowly up to about 100yr, and at this point all these emissivities decrease abruptly,when the electron temperature diminishes to less than about7,000 K. On the other hand, the emissivity of [O iii ] λ Ionic abundances were calculated by using the physical con-ditions derived for each epoch and the dereddened line in-tensities presented in Tables 2 and 3. The electron densityderived from [S ii ] λλ ii ] λλ / line ratio wereused in each case. When no density value was available, weassumed a value of 100 cm − . The results are presented inTable 4.Total abundances were derived from the ionic abun-dances and the ionisation correction factors (ICF) fromKingsburgh & Barlow (1994) and Delgado-Inglada et al.(2014) to correct for the ions not visible in the optical spec-trum. The expressions used for calculating the ICFs and toderive the total abundances are listed in the Appendix. Thevalues obtained in each case are presented in Table 4, be-low the ionic abundances of each element. The derived totalabundances are also listed in Table 4.In the following, for the discussion, we adopt the abun-dances obtained from the GTC 2018 observations for bothnebulae, which present the lowest uncertainties and are morecomplete. PM 1-188 outer nebula shows sub-solar O, S andAr abundances, with 12+log O/H = 8.05 ± ± ± ± − ± ± ± − ± ± − ± − ± MNRAS , 1–15 (2015) Year T e ( × K ) T e [N II] (PM1-188) T e [N II] (Inner) T e [N II] (Outer) Figure 4.
Time evolution of electron temperature in the inner(blue) and outer (green) nebulae.
BPT diagnostic diagrams (Baldwin et al. 1981; Veilleux &Osterbrock 1987) have been used for many years to analysethe ionising mechanism producing an ionised region. In fact,Baldwin et al. (1981), suggested a two-dimensional quan-titative classification scheme to separate, verbatim, “nor-mal H ii regions, planetary nebulae, objects photoionisedby a power-law continuum, and objects excited by shock-wave heating”. The most used line intensity ratios are log[O iii ] λ β vs. log [N ii ] λ α , log [O iii ] λ β vs. log [S ii ]( λλ α , and log [O iii ] λ β vs.log [O i ] λ α .More recently Kewley et al. (2001) built a theoreticalclassification scheme to establish an upper limit for starburstmodels on the optical BTP diagnostic diagrams. While BPTdiagrams were constructed to separate nebulae with differ-ent ionisation mechanisms and the theoretical Kewley et al.(2001) models where constructed to separate starbursts fromAGNs, the physics of the line formation in PNe is the same,therefore we will study the nature of the line mechanism pro-duction in the inner and outer nebulae in PM1-188 throughthese diagrams. This is not the first time that diagnostic di-agrams are used to this purpose in PNe, (see e.g., Raga etal. 2008; Akras & Goncalves 2016, and references therein).In Fig. 5, we present diagrams of [O iii ] λ β vs.[N ii ] λ α , [S ii ] λ (6717+6731)/H α , and [O i ] λ α for the inner and outer nebulae. The dereddened intensityvalues are taken from Tables 2 and 3 and are plotted withfilled and open symbols for the inner and outer nebulae,respectively. The different epochs are marked with differentsymbols and colours as explained in the figure caption.It should be noticed that, as we said before, the [N ii ] λ ii λ ii ] λ ii ] λ iii ] λ β intensity ratio and lower values of the low-ionisation lines, relative to H α . This is particularly notice-able for the GTC-2018 data (magenta diamonds) where thefilled symbols (inner nebula) lie to the right of the Kewley’smodels in the cases of [N ii ] λ α and [O i ] λ α ,(the case of [S ii ]/H α line ratio is marginal), indicating thatthe ionising mechanisms in this zone is not photoionisationbut a different one, like heating by shocks. As this occurs forall the epochs, it implies that the inner nebula appears al-ways as heated by shocks. In the case of [S ii ] lines (panel b),almost all the points are placed in the photoionisation zone.This would be related to the radiation strength in UV wave-length adopted by Kewley´s models which were built withsolar metallicities. In a metal-deficient environment suchas in PM 1-188 which shows low S abundance, the respec-tive [O iii ] λ β and [S ii ]/H α intensities are expectedto become stronger and lower than the values indicated byKewley´s models. Thus, the observed low [O iii ] λ β and strong [S ii ]/H α values in the inner nebula could be ex-plained by shocks rather than photoionisation.Thus, by means of BPT diagrams, we have corroboratedthat the inner nebula is mainly heated by shocks, while theouter nebula is photoionised. We first noticed an outflow ejected from the central star inthe OAN-SPM echelle REOSC spectrum obtained in 2004,with the slit position in the E-W direction (P.A. 90 ◦ ). Theoutflow was clearly visible in both [N ii ] λ λ α line. The jet emissionwas more intense and extended to the blue side of the lines,being the red emission fainter.These results were confirmed in the high resolutionspectrum obtained with the MES on 2017-06-29 (see Fig.6). The outflow found in the 2004 echelle REOSC spectrumappears better defined in the MES spectrum due to the bet-ter spectral resolution of MES, and to a larger exposuretime. This outflow is evidently ejected from the central star.The spectrum obtained for H α , in this orientation, showsan ellipse with a cavity in the centre where the expansionvelocity is the highest.The spectrum obtained with MES at P.A. of 0 ◦ does notshow the outflow (see Fig. 6). Interestingly the two lobes ofthe H α emission, corresponding to the outer nebula, appeartilted with a displacement of about 15 km s − , being thenorthern zone to the blue. This is showing that the outernebula is not spherical as it appears in the images (and soit was classified by Sahai et al. 2011 as R,*,ib: round withan inner bubble, a visible star). Our H α images are equalto the PV diagrams presented by Guerrero et al. (2018)who claimed that this behaviour corresponds to a barrel-like structure whose symmetry axis is tilted to the line ofsight by 25 ◦ . MNRAS , 1–15 (2015)
M. Peña et al.
Figure 5.
BPT diagrams, showing the behaviour of [O iii ] λ β vs. [N ii ] λ ii ] λ (6716+6731), and [O i ] λ α .The solid and dashed lines in this figure correspond to Kewley et al. (2001) models and their uncertainties, respectively, and separate H ii region-like objects from AGNs. The different epochs are marked with different symbols, OAN-2000: vertical triagles, OAN-2002: squares,OAN-2005: stars, LCO-2005: crosses, OAN-2017: horizontal triangles, GTC-2018: diamonds. The filled symbols correspond to the innernebula and the empty symbols, to the outer zone. The points to the right of the solid line are interpreted as regions heated by shocks. In the N-S oriented image (P.A. 0 ◦ ), the [N ii ] emissionis strong and wide in the centre and faint and elongatedtowards the N-S direction, showing some faint filamentarystructure.The observed radial velocity of the object, derived fromthese high resolution spectra, is +56 . ± − and theheliocentric velocity is v sys = 53.9 ± − .The velocity fields obtained from the MES high resolu-tion spectra are analysed in the next section, by means ofPosition-Velocity diagrams. Position-velocity diagrams (PVD) were constructed for theMES 2017 data at slit positions P.A. 90 ◦ and P.A. 0 ◦ , forthe [N ii ] λ α lines in order to explore the nebu-lar kinematics. It is safe to use these lines because at thisepoch the star was too faint, it was not detected and doesnot contaminate the nebular lines. These PVD were pro-duced in Python, using the Astropy library. Our code readsour MES data and plots them using the proper scale plateand the radial velocity resolution. As the slit was placed onthe central star for the two positions, P.A. 90 ◦ and P.A. 0 ◦ ,the spatial origin is the position of the central star and thespatial distribution takes into account a pixel size of 0.36 (cid:48)(cid:48) .To reconstruct the radial velocities, we used the rest wave-lengths of H α and [N ii ] λ sys =53.9 km s − ), thus we can analyse the kine-matics of these lines.The PV diagrams are presented in Fig. 6. The line con-tours in each diagram represent the intensities of the lines,each line has its ad-hoc contour separation, as explainedlater. ◦ In Fig. 6 (up, left), we can observed the H α emission pre-senting two peaks and a cavity in the centre. The contours show a maximum expansion velocity HWHM of about 40km s − in the centre.The contours for the H α PVD represent the 10%, 20%,30%, 40%, and 50% of the total intensity from the outerto the inner contour respectively. The 10% contour shows afaint elongation to the blue.The [N ii ] λ α , but only one cen-trally concentrated bright emission (slightly blue shifted).This zone has an expansion velocity of about 40 km s − .The most spectacular fact in this diagram is the evident out-flow emerging from the central star, with velocities rangingfrom −
150 to 100 km s − . The outflow to the blue is moreextended and brighter than the red one, possibly due to alarger extinction in the back. The outflow is visible mainlyin the [N ii ] lines. It is only barely observed in H α becauseits position coincides with the cavity shown in H α in thecentral zone. ◦ The H α and [N ii ] λ α PVD represent the same in-tensity levels as in the E-W PV diagram. In this case theH α lobes, corresponding to the north and south sides of theouter nebula, are displaced in velocity by about 15 km s − one from the other. The northern lobe is shifted to the blue.This behaviour indicates that the nebula is not spherical,but an open-ended (barrel-like) elongated structure.The [N ii ] λ ≤
40 kms − . A faint filamentary structure extending from North toSouth, of about 20 arcsec long, corresponding with the outerzone, is noticed. No outflow is found in this orientation.From the PV diagrams the kinematic ages for the inner MNRAS , 1–15 (2015) Figure 6.
Position-velocity diagrams for the H α and [N ii ] λ ◦ ) and lower panels, with the slit oriented N-S (P.A. 0 ◦ ). and outer nebulae can be derived. A distance of 5.36 kpc toPM 1-188 as published by Frew et al. (2016) was adopted.Therefore, from our H α image at P.A.=90 ◦ which shows aprojected nebular size of 10 arcsec and an expansion velocityof 40 km s − , we compute a kinematic nebular age of 6,400yr for the outer nebula. This is a lower limit if the structureis elongated and inclined with respect to the line of sightas proposed by Guerrero et al. (2018). For the inner nebulawe used the [N ii ] λ ◦ image and performed thesame calculations taking into account a size of 3 arcsec andan expansion velocity of 40 km s − , from these data we com-puted a lower-limit kinematic age of 1,900 yr. A kinematicage can be determined also for the outflow by assuming thatwe are seeing only one pulse. This structure shows a size of2 arcsec and an expansion velocity of 150 km s − , then akinematic age equal to 335 yr is derived. PM 1-188 is a planetary nebula with a central star of therare spectral type [WC 10], whose brightness has been di-minishing with time, losing about 10 magnitudes in the last40 years. The central star seems to have had an event of theborn-again type as claimed by Guerrero et al. (2018). A fewborn-again central stars have been reported in the literature.The most studied cases are A 30, A 78, A 58 and recentlyWR 72 (Wesson et al. 2008; Gvaramatze et al. 2020, and ref-erences therein). They have been explained through a verylate thermal pulse (VLTP) of a post-AGB star. The maincharacteristics of these born-again central stars (mostly clas-sified as [WC early] or [WC early]-PG1159 stars) and theirassociated PNe are the existence of an outer evolved H-richnebula, and an inner H-deficient shell which in many casesshows knots with cometary tails extended radially. Interest-ingly, a binary central star was found in A 30 (Jacoby etal. 2020). Some of these characteristics have been found in
MNRAS000
MNRAS000 , 1–15 (2015) M. Peña et al.
PM 1-188, in particular the existence of an outer low densitynebula, and a compact denser inner shell showing an outflowsimilar to the cometary tails. However the effective temper-ature of its central star is much lower than that found inthe mentioned born-again stars, and such objects have notshown the decrease in brightness displayed by PM 1-188 cen-tral star.In this work the spectrophotometric behaviour of theionised gas in PM 1-188 is analysed by means of medium andhigh resolution spectra obtained from observations along 20years. This allows us to study the time evolution of emissionlines in the inner and outer nebulae. Also, the time evolutionof physical parameters is analysed.The outer extended nebula emits H and He lines, andfaint low-ionised heavy-element (N + , O + , and S + ) lines. Wefound that the heavy-elements emission lines have been de-creasing with time and also the electron temperature hasdiminished from about 13,000 K previous 2005 to less than7,000 K in 2018. Therefore this section of the nebula, ejectedduring the first time the star was in the AGB zone, has beenrecently cooling and recombining with time as a consequenceof the fading of the central star because less and less ionisingphotons arrive to the outer nebula with time. A kinematicage of 6,400 yr was calculated from the projected nebulardimension, the distance and the expansion velocity. This agecould be larger if the structure is elongated and inclined, aspropose Guerrero et al. (2018).The inner nebula shows very intense lines of low ionisedspecies such as N + , O + , and S + . Also weak lines of O ++ ,Ne ++ and He ++ are found in this zone. It is particularlyinteresting the faint emission of the He ii λ ++ should be produced in the shocked region, as predicted bythe shock model presented by Guerrero et al. (2018). BPTdiagrams constructed by us also indicate that the ionisingmechanisms in the inner nebula is not photoionisation but adifferent one as heating by shocks. Besides the unexpectedemergence of He ii λ iii ] λ ii ] λ ii ] λ ± ± − ± ± − ± − ± ± ± ± ± ± ± α elements O,S and Ar indicate that PM1-188 was formed in a low abun-dance medium, while the enrichment of He, N and Ne seemsa consequence of the stellar nucleosynthesis. The chemistryof PM 1-188 is reminiscent of the chemical abundances ofsome halo PNe (see the analysis for DdDm-1 and other haloPNe presented by Henry et al. 2008), in particular it is remi-niscent of BoBn 1 which is very O, S and Ar poor and showslarge C, N and Ne enrichment. Otsuka et al. (2010) derived12+log O/H = 7.74, log N/O = 0.29, log Ne/O = 0.22, logS/O = -2.42, and log Ar/O = -3.41, and to explain the ele-mental abundances in BoBn 1 Otsuka et al. (2010) proposedfor this nebula a progenitor of 1 – 1.5 M (cid:12) initial mass staror a 0.75+1.5 M (cid:12) binary system. To suggest a binary sys-tem as origin of PM 1-188 is tempting considering the rarebehaviour of the central star which ejected a nebula morethan 6,000 yr ago and actually, after an apparent born-againepisode, presents atmospheric instabilities, a second shockedyounger nebula and outflows. However, so far there is notevidence for a binary central star.An outflow emerging from the central zone was foundin our high resolution spectra obtained in P.A. 90 ◦ . It isbright in [N ii ] lines and does not appear in H α due to Hlines show a cavity in this zone. From our position-velocitydiagrams, it is found that this outflow shows a velocity from −
150 to 100 km s − . The blue side is brighter and the redside is fainter possibly due to large reddening in the back-side. We also estimated the age of the outflow to be 335 yr.Rechy-García et al. (2020) reported to have found a shell-like structure expanding at 300 km s − with an age of 200yr in PM 1-188. This coincides very well with the outflow indirection N-S we are finding. Due to we have not obtainedMES observations in other directions we can not assert thatthe ejection has a shell-like structure, but we do not detectsuch an outflow in the observations with P.A. 0 ◦ . Thereforethe structure detected by Rechy-García et al. (2020) shouldbe a broken shell. If this structure were a complete shell, theenergy required to keep it in motion should be larger than10 erg, which is too large for the central star to provide it.A hydrodynamical model taking into account the detectedoutflow and the photoionisation by the central object is inprogress to study in detail the kinematics in the inner shelland its evolution (Rodríguez-González et al. in prep.)Other result obtained in this work, is the systemic radialvelocity of 56 ± − and a heliocentric velocity of 53.9 ± − obtained from our high resolution spectra.Since its discovery, the planetary nebula PM 1-188 hasgiven us many surprises and we have slowly advanced in theunderstanding of this object, but still several facts remainunexplained. Deeper observations in the optical and otherwavelength ranges are necessary to get a deep insight of thispeculiar object. ACKNOWLEDGEMENTS
This work is based upon observations carried out at the Ob-servatorio Astronómico Nacional on the Sierra San PedroMártir (OAN-SPM), Baja California, México, and at LasCampanas Observatory with the Magellan telescope Clay.Work based on data from the GTC Archive at CAB (INTA-
MNRAS , 1–15 (2015) CSIC). We thank the daytime and night support staff at theOAN-SPM for facilitating and helping to obtain our obser-vations, along 20 years. J.S. Rechy-García and V. Gómez-Llanos are thanked for helpful discussion and support duringobservations. L.H.-M. is grateful to Dra. Maria del Pilar Car-reón for the hospitality at ICN, UNAM. F.R.-E. acknowl-edges scholarship from CONACyT, México. This work re-ceived partial support from DGAPA-PAPIIT IN103117 andIN105020, UNAM.Data Availability Statement: The data underlying thisarticle will be shared on reasonable request to the corre-sponding author.
REFERENCES
Acker, A., Neiner, C., 2003, A&A, 403, 659Asplund, M., Grevesse, N., Sauval, A. J., & Scott, P. 2009,ARA&A, 47, 481Akras, S. & Goncalves, D., 2016, MNRAS, 455, 930Cardelli, J. A. A., Clayton, G. C., Mathis, J. S. 1989, ApJ, 345,245Baldwin, J. A., Phillips, M. M., & Terlevich, R. 1981, PASP, 93,5Cohen, M., Barlow, M. J., Liu, X.-W., & Jones, A. F. 2002, MN-RAS, 332, 879Delgado-Inglada, G., Morisset, C., & Stasińska, G. 2014, MNRAS,440, 536Froese Fischer, C., & Tachiev, G. 2004, At. Data Nucl. Data Ta-bles, 87, 1Frew, D. J., Parker, Q. A., & Bojicic, I. S., 2016, MNRAS, 455,1459.Galavís, M. E., Mendoza, C., & Zeippen, C. J. 1995, A&AS, 111,347Galavís, M. E., Mendoza, C., & Zeippen, C. J. 1997, A&AS, 123,159Gvaramadze, V. V., Kniazev, A. Y., Gräfener, G.,& Langer, N.2020, MNRAS, 492, 3316Gorny, S. K., Stasińska, G, Szczerba, R., & Tylenda, R. 2001,A&A, 377, 1007Guerrero, M. A., Fang, X, Miller-Bertolami, M. M., et al. 2018,NatureAstronomy, 2, 784Jacoby, George H., Hillwig,Todd C., & Jones, David, 2020, MN-RAS, 498, L114Henry, R.B.C:, Kwitter, K., Dufour, R. J., & Skinner, J. N. 2008,ApJ, 680, 1162Hu, J. Y., Bibo, E. A., 1990, A&A, 234, 435Kaufman, V., & Sugar, J. 1986, J. Phys. Chem. Ref. Data, 15,343Kewley, L. J., Dopita, M. A., Sutherland, R. S., Heisler, C. A., &Trevena, J., ApJ, 556, 121Kingsburgh, R. L., & Barlow, M. J. 1994, MNRAS, 271, 267Kisielius, R., Storey, P. J., Ferland, G. J., & Keenan, F. P. 2009,MNRAS, 397, 903Koesterke, L. 2001, ApSS, 275, 41Leuenhagen, U., & Hamann, W.-R. 1998, A&A, 330, 265Luridiana, V. , Morisset, C., Shaw, R. A. 2015, A&A, 573, A42McLaughlin, B. M., & Bell, K. L. 2000, J. Phys. B, 33, 597Meaburn, J., Blundell, B., Carling, R., Gregory, D. F., Keir, D.,et al. 1984, MNRAS, 210, 463Meaburn, J., López, J. A., Gutiérrez, L., Quiroz, F., Murillo, J.M., et al. 2003, Rev. Mex. Astron. Astrofis., 39, 185Mendoza, C. 1983, IAU Symp., 103, 143Osterbrock, D. E., & Ferland, G. J., Astrophysics of Gaseous Neb-ulae and Active Galactic Nuclei (2nd ed.), University ScienceBooks, 2006 Otsuka, M., Tajitsu, A., Hyung, S., & Izumiura, H. 2010, ApJ,723, 6580Peña, M. 2005a, Rev. Mex. Astron. Astrofis., 41, 423Peña, M. 2005b, Rev. Mex. Astron. Astrofis.CS, 23, 120Peña, M., Torres-Peimbert, S., & Ruiz, M. T. 1991, PASP, 103,865Peña, M., Hamann, W.-R., Koesterke L., et al., 1997, ApJ, 491,233Peña, M., Stasińska, G., & Medina, S. 2001, A&A, 367, 983Podobedova, L. I., Kelleher, D. E., & Wiese, W. L. 2009, J. Phys.Chem. Ref. Data, 38, 171Pollacco, D.L, Kilkenny, D., Marang, F., van Wyk, F., & Roberts,G. 1992,MNRAS, 256, 671Pollacco, D. L., & Hill P. W. 1994, MNRAS, 267, 692Porter R. L., Ferland G. J., Storey P. J., & Detisch M. J., 2012,MNRAS, 425, L28Porter R. L., Ferland G. J., Storey P. J., & Detisch M. J., 2013,MNRAS, 433, L89Raga, A.C., Riera, A., Mellema, G., Esquivel, A., and Velázquez,P. F., 2008, A&A, 489, 1141Rechy-García, J. S., Guerrero, M. A., Santamaría, E., Gómez-González, V. M. A., Ramos-Larios, G., Toalá, J. A., Cazzoli,S., Sabin, L., Miranda, L. F., Fang, X., Liu, J., 2020, ApJL,903, id.L4Sahai, R., Morris, M. R., & Villar, G. G. 2011, AJ, 141Storey P. J., & Hummer D. G., 1995, MNRAS, 272, 41Storey P. J., & Zeippen C. J., 2000, MNRAS, 312, 813Storey, P. J., Sochi, T., & Badnell, N. R. 2014, MNRAS, 441,3028Tayal, S. S. 2011, ApJS, 195, 12Tayal, S. S., & Zatsarinny, O. 2010, ApJS, 188, 32Todt, H., Guerrero, M. A., Fang, X. et al. 2015, ASPC, 493, 141Veilleux, S., & Osterbrock, D. E., 1987, ApJS, 63, 295Werner K., Hamann W.-R., Heber U. et al., 1992, A&A, 259, L69Wesson, R., Barlow, M.J., Liu, X.-W., Storey, P.J., Ercolano, B.,De Marco, O. 2008, MNRAS, 383, 1639MNRAS , 1–15 (2015) M. Peña et al. T a b l e . F l u x e s o f t h e i nn e r n e bu l a i nd i ff e r e n t e p o c h s , F ( λ ) a . D a t a p r e v i o u s t o2005 i n c l ud e s t e ll a r e m i ss i o n li n e s . T h ec o l u m n I ( λ ) b c o rr e s p o nd s t o t h e d e r e dd e n e dflu x e s . T e l e s c . O AN - m O AN - m O AN - m O AN - m L C O - C l a y O AN - m G TC Sp ec tr . B C h B C h ec h e ll e B C h L SS T - B C h O S I R I S - L S d a t e I o n λ F ( λ ) I ( λ ) F ( λ ) I ( λ ) F ( λ ) I ( λ ) F ( λ ) I ( λ ) F ( λ ) I ( λ ) F ( λ ) I ( λ ) F ( λ ) I ( λ ) [ O ii ] O o R O o R . . ( ) O o R O o R O o R O o R O o R O o R . . ( ) . . ( ) [ N e iii ] O o R O o R n o i s y n o i s y O o R O o R O o R O o R O o R O o R n o i s y n o i s y . . ( ) H + H e i O o R O o R n o i s y n o i s y O o R O o R O o R O o R O o R O o R n o i s y n o i s y . . ( ) H n o i s y n o i s y n o i s y n o i s y O o R O o R O o R O o R O o R O o R n o i s y n o i s y . . ( ) H e i n o i s y n o i s y n o i s y n o i s y O o R O o R O o R O o R n o i s y n o i s y n o i s y n o i s y . : . : [ S ii ] n o i s y n o i s y n o i s y n o i s y O o R O o R O o R O o R n o i s y n o i s y n o i s y n o i s y . . ( ) H δ . . ( ) n o i s y n o i s y O o R O o R . . ( ) n o i s y n o i s y . . ( ) . . ( ) H γ b l e ndb l e nd . . ( ) O o R O o R . . ( ) n o i s y n o i s y . . ( ) . . ( ) H e i . . ( ) . . ( ) O o R O o R . . ( ) n o i s y n o i s y . . ( ) . . ( ) H e ii c W R - b W R - b W R - b W R - b < < . . . ( ) . . ( ) . . ( ) . . ( ) H β . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) [ O iii ] . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) [ O iii ] . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) [ N i ] . . ( ) . . ( ) n o i s y n o i s y . . ( ) n o i s y n o i s y . . ( ) . . ( ) [ N ii ] . . ( ) . . ( ) n o i s y n o i s y . . ( ) . . ( ) . . ( ) . . ( ) H e i . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) [ O i ] . . ( ) . . ( ) n o i s y n o i s y . . ( ) . . ( ) . . ( ) . . ( ) [ O i ] b l e ndb l e ndb l e ndb l e ndn o i s y n o i s y . : . : n o i s y n o i s y . . ( ) . . ( ) [ N ii ] . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) H α . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) [ N ii ] . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) . . ( ) H e i O o R O o R . . ( ) n o i s y n o i s y . . ( ) O o R O o R . . ( ) . . ( ) [ S ii ] O o R O o R . . ( ) . . ( ) . . ( ) O o R O o R . . ( ) . . ( ) [ S ii ] O o R O o R . . ( ) . . ( ) . . ( ) O o R O o R . . ( ) . . ( ) H e i O o R O o R . . ( ) —— . . ( ) O o R O o R . . ( ) . . ( ) [ A r iii ] O o R O o R n o i s y n o i s y —— . . ( ) O o R O o R . . ( ) . . ( ) C ii O o R O o R . . ( ) —— . . ( ) O o R O o R . . ( ) . . ( ) [ O ii ] + O o R O o R . . ( ) —— O o R O o R O o R O o R . . ( ) . . ( ) F ( H β ) d . . . . . . . c ( H β ) . ± . . ± . . ± . . ± . . ± . . ± . . ± . O o R i s O u t o f R a n g e ; — i nd i c a t e s li n e n o t d e t ec t e d o r n o t m e a s u r a b l e . a T h ee rr o r s i n fl u x a r e a b o u t % w h e n F ( λ ) / F ( H β ) ≥ , a b o u t % w h e n F ( λ ) / F ( H β ) ∼ . , a b o u t % w h e n F ( λ ) / F ( H β ) ≤ . a li n e s m a r k e d w i t h : a r e v e r y un ce rt a i n . b T h e un ce rt a i n t i e s i nd e r e dd e n e d i n t e n s i t i e s I ( λ ) , c a l c u l a t e db y c o n s i d e r i n g t h e un ce rt a i n t i e s i n c ( H β ) , a r e g i v e n i n % a nd e n c l o s e d i np a r e n t h e s i s . c W R - b i s W o l f - R a y e t bu m p . d T h e a b s o l u t e fl u x i n H β i s i nun i t e s o f E - e r g c m − s − a nd i t i s o n l y i nd i c a t i v e a s i t d e p e nd s o n t h ee x tr a c t i o n w i nd o w , t h e o b s e r v i n g c o nd i t i o n s ( ph o t o m e tr i c o r n o t) , e t c . MNRAS , 1–15 (2015) Table 3.
Fluxes of the outer nebula in different epochs, F( λ ) a . Data previous to 2005 can include stellar emission. The column I( λ ) b corresponds to the dereddened fluxes.Telesc. OAN-2m OAN-2m OAN-2m LCO-Clay OAN-2m GTCSpectr. BCh BCh BCh LSST-2 BCh OSIRIS-LSDate 2000/09 2002/06 2005/05 2005/08 2017/06 2018/06Ion λ F( λ ) I( λ ) F( λ ) I( λ ) F( λ ) I( λ ) F( λ ) I( λ ) F( λ ) I( λ ) F( λ ) I( λ )[O ii ] 3727 OoR OoR 98.0 303.2(30) noisy noisy OoR OoR noisy noisy 21.0 41.0(20)H9 3835 OoR OoR noisy noisy noisy noisy noisy noisy noisy noisy 4.3: 8.0:H8+He i i ii ] 4068 OoR OoR noisy noisy noisy noisy noisy noisy noisy noisy 3.3: 4.4:H δ γ i β i iii ] 4959 20.6 19.3(30) 14.2 13.0(30) 7.2 6.7(20) noisy noisy noisy noisy noisy noisy[O iii ] 5007 29.3 26.5(25) 41.3 36.2(25) 17.1 15.4(15) 15.2 13.9(15) 6.6 6.1(20) 3.7 3.4(25)He i i ] 5198 9.2 7.4(35) 8.5 6.4(35) 5.9 4.7(25) noisy noisy 4.5 3.8(25) 3.9 3.3(25)[N ii ] 5755 8.1 5.0(35) 8.0 4.2(35) 5.8 3.5(25) 4.0 2.6(25) 1.2 0.8(30) 0.8 0.6(25)He i i ] 6300 28.4 14.2(35) 14.1 5.6(40) 9.6 4.6(30) noisy noisy 6.1 3.6(30) 3.8 2.2(25)[O i ] 6364 28.5 13.9(35) 7.8: 3.0: 4.7: 2.2: noisy noisy noisy noisy 1.4: 0.8:[N ii ] 6548 132.0 60.5(20) 174.0 61.7(20) 125.7 55.4(10) 79.9 39.5(10) 52.4 28.7(10) 69.6 37.7(8)H α ii ] 6583 674.2 305.1(14) 708.2 246.9(14) 417.1 181.4(4) 272.0 132.9(4) 194.0 105.3(4) 214.5 114.9(3)He i ii ] 6716 73.9 31.9(40) 73.0 24.3(35) 60.4 25.0(25) OoR OoR 37.4 19.6(25) 38.9 20.1(10)[S ii ] 6731 71.4 30.7(35) 60.0 19.9(40) 49.6 20.4(30) OoR OoR 29.9 15.6(30) 29.1 15.0(15)He i iii ] 7136 OoR OoR noisy noisy 5.6: 2.0: OoR OoR 5.8: 2.7: 1.0: 0.5:C ii ii ]+ 7325 OoR OoR 31.2 7.7(40) OoR OoR OoR OoR 3.3 1.5(30) 1.8 0.8(25)F(H β ) c β ) 1.15 ± ± ± ± ± ± a The errors in flux are about 2-3% when F( λ )/F(H β ) ≥
1, about 10 % when F( λ )/F(H β ) ∼ λ )/F(H β ) ≤ a lines marked with : are very uncertain. b The errors in I( λ ), taking into account the uncertainties in c(H β ), are given in % enclosed in parenthesis. c The absolute flux in H β is in units of E-14 erg cm − s − and it is only indicative as it depends on the extraction window, the observing conditions, etc.MNRAS000
1, about 10 % when F( λ )/F(H β ) ∼ λ )/F(H β ) ≤ a lines marked with : are very uncertain. b The errors in I( λ ), taking into account the uncertainties in c(H β ), are given in % enclosed in parenthesis. c The absolute flux in H β is in units of E-14 erg cm − s − and it is only indicative as it depends on the extraction window, the observing conditions, etc.MNRAS000 , 1–15 (2015) M. Peña et al. T a b l e . P h y s i c a l c o nd i t i o n s ,i o n i c a nd t o t a l a bund a n ce s f o r t h ece n t r a l a nd o u t e r n e bu l a e , o b t a i n e d w i t h P y N e b , f o r a ll t h ee p o c h s D a t e ( yy - mm ) - - - - - - - - - - - - Z o n e I nn e r I nn e r I nn e r I nn e r I nn e r I nn e r O u t e r O u t e r O u t e r O u t e r O u t e r O u t e r n e [ S ii ] ( c m − ) ± ± ± ± ± ± ± ± ± T e [ N ii ] ( K ) . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . T e [ S ii ] ( K ) ————— . ± . —————— T e [ O ii ] ( K ) — . ± . —— . ± . . ± . — . ± . ——— . ± . H e + / H + —— . ± . . ± . . ± . . ± . —— . ± . . ± . . ± . . ± . H e ++ / H + ( − ) —— . ± . . ± . . ± . . ± . —————— I C F ( H e + + H e ++ ) —— . . . . —— . . . . O + / H + ( − ) — . ± . —— . ± . . ± . — . ± . ——— . ± . O ++ / H + ( − ) . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . I C F ( O + +O ++ ) — . —— . ± . . ± . — . ——— . N + / H + ( − ) . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± . I C F ( N + ) — . ± . —— . ± . . ± . — . ± . ——— . ± . S + / H + ( − ) — . ± . . ± . — . ± . . ± . . ± . . ± . . ± . — . ± . . ± . I C F ( S + + S ++ ) — . ± . —— . ± . . ± . — . ± . ——— . ± . N e ++ / H + ( − ) ————— . ± . —————— I C F ( N e ++ ) ————— . ± . —————— A r ++ / H + ( − ) —— . ± . — . ± . . ± . —— . ± . — . ± . . ± . I C F ( A r ++ ) —— . — . . —— . — . . H e / H —— . ± . . ± . . ± . . ± . —— . ± . . ± . . ± . . ± . O / H — . ± . —— . ± . . ± . — . ± . ——— . ± . N / H — . ± . —— . ± . . ± . — . ± . ——— . ± . N e / H ————— . ± . —————— S / H — . ± . —— . ± . . ± . — . ± . ——— . ± . A r / H —— . ± . — . ± . . ± . —— . ± . — . ± . . ± . N / O — − . ± . —— − . ± . − . ± . — − . ± . ——— − . ± . N e / O ————— . ± . —————— S / O — − . ± . —— − . ± . − . ± . — − . ± . ——— − . ± . A r / O ———— − . ± . − . ± . ————— − . ± . ∗ A d o p t e d v a l u e s . a T o t a l a bund a n ce s i n + l og X / H s c a l e . b S o l a r a bund a n ce s a d o p t e d f r o m ( A s p l und e t a l. ) a r e + l og H e / H = . , + l og O / H = . , + l og C / H = . , + l og N / H = . , + l og N e / H = . , + l og A r / H = . , + l og S / H = . MNRAS , 1–15 (2015) Table A1. Atomic parameters used in PyNeb calculations
Ion Transition probabilities Collisional strenghtsN + Froese Fisher & Tachiev (2004) Tayal (2011)O + Froese Fisher & Tachiev (2004) Kisielius et al. (2009)O ++ Froese Fisher & Tachiev (2004) Storey et al. (2014)Storey & Zeippen (2000)Ne ++ Galavís et al. (1995) McLaughlin & Bell (2000)S + Podobedova et al. (2009) Tayal & Zatsarinny (2010)Ar ++ Mendoza (1983) Galavís et al. (1995)Kaufman & Sugar (1986).Ion Effective recombination coefficientsH + Storey & Hummer (1995)He + Porter et al. (2012, 2013)He ++ Storey & Hummer (1995)
APPENDIX A: ATOMIC PARAMETERS USEDIN PYNEB CALCULATIONS AND THEIONISATION CORRECTION FACTORS
ICFs used for the total abundances calculation are listednext. Those marked with KB94 correspond to the expres-sions from Kingsburgh & Barlow (1994) and those markedDI14 correspond to the expressions given by Delgado-Inglada et al. (2014). • HeH = He + H + . If He ++ is detected HeH = He + +He ++ H + . • OH = ICF(O) × O + +O ++ H + . ICF(O) = 1 if no He ++ is de-tected.Otherwise, ICF(O) is given by the equation (12) in DI14. • NH = ICF(N) × N + H + . ICF(N) = OO + (KB94). • ArH = ICF(Ar) × Ar ++ H + . ICF(Ar) = 1 . (KB94). • NeH = ICF(Ne) × Ne ++ H + . ICF(Ne) = OO ++ (KB94). • SH = ICF(S) × S + +S ++ H + . ICF(S) = (cid:20) − (cid:16) − O + O (cid:17) (cid:21) − / .As no S +2 is detected, it can be estimated through theexpression: S ++ S + = 4 .
677 + (cid:16) O ++ O + (cid:17) . (KB94). This paper has been typeset from a TEX/L A TEX file prepared bythe author.MNRAS000