The Origin and Evolution of the Halo PN BoBn 1: From a Viewpoint of Chemical Abundances Based on Multiwavelength Spectra
We have performed a comprehensive chemical abundance analysis of the extremely metal-poor ([Ar/H]–2) halo planetary nebula (PN) BoBn 1 based on IUE archive data, Subaru/HDS spectra, VLT/UVES archive data, and Spitzer/IRS spectra. We have detected over 600 lines in total and calculated ionic and elemental abundances of 13 elements using detected optical recombination lines (ORLs) and collisionally excited lines (CELs). The estimations of C, N, O, and Ne abundances from the ORLs and Kr, Xe, and Ba from the CELs are done the first for this nebula, empirically and theoretically. The C, N, O, and Ne abundances from ORLs are systematically larger than those from CELs. The abundance discrepancies apart from O could be explained by a temperature fluctuation model, and that of O might be by a hydrogen deficient cold component model. We have detected 5 fluorine and several slow neutron capture elements (the s-process). The amounts of [F/H], [Kr/H], and [Xe/H] suggest that BoBn 1 is the most F-rich among F detected PNe and is a heavy s-process element rich PN. We have confirmed dust in the nebula that is composed of amorphous carbon and PAHs with a total mass of 5.810 . The photo-ionization models built with non-LTE theoretical stellar atmospheres indicate that the progenitor was a 1-1.5 star that would evolve into a white dwarf with an 0.62 core mass and 0.09 ionized nebula. We have measured a heliocentric radial velocity of +191.6 1.3 km s and expansion velocity 2 of 40.5 3.3 km s from an average over 300 lines. The derived elemental abundances have been reviewed from the standpoint of theoretical nucleosynthesis models. It is likely that the elemental abundances except N could be explained either by a 1.5 single star model or by a binary model composed of 0.75 + 1.5 stars. Careful examination implies that BoBn 1 has evolved from a 0.75 + 1.5 binary and experienced coalescence during the evolution to become a visible PN, similar to the other extremely metal-poor halo PN, K 648 in M 15.
Subject headings:ISM: planetary nebulae: individual (BoBn 1, K 648), ISM: abundances, ISM: dust, Stars: Population II
Planetary Nebulae (PNe) represent a stage in the evolution of low- to intermediate-mass stars with initial masses of 1-8 . At the end of their life, a star of such mass evolves first into a red giant branch (RGB) star, then an asymptotic giant branch (AGB) star, next a PN, and finally a white dwarf. During their evolution, such stars eject a large amount of their mass. The investigation of chemical abundances in PNe enables the determination of how much of a progenitor’s mass becomes a PN, when and how elements synthesized in the progenitor were brought to the surface, and how chemically rich the Galaxy was when the progenitors were born.
Currently, over 1,000 objects are regarded as PNe in the Galaxy (Acker et al. 1991). Of these, about 14 objects have been identified as halo members from their location and kinematics since the PN K 648 was discovered in M 15 (Pease 1928). Halo PNe are interesting objects as they provide direct insight into the final evolution of old, low-mass halo stars, and they are able to convey important information for the study of low-mass star evolution and the early chemical conditions of the Galaxy. However, in extremely metal-poor and C- and N-rich ([C,N/O]0, [Ar/H]–2) halo PNe, there are unresolved issues on chemical abundances and evolution time scales. BoBn 1 (PN G108.4–76.1) is one of the C- and N-rich and extremely metal-poor halo PNe ([C, N/O]1, [Ar/H]=–2.220.09, [Fe/H]=–2.390.14; this work), which composes a class of PN together with K 648 (Otsuka 2007, see Table 20) and H 4-1 (Otsuka et al. 2003).
The progenitors of halo PNe are generally thought to be 0.8 stars, which is the typical mass of a halo star. Above mentioned three metal-poor C- and N-rich halo PNe, however, show signatures that they have evolved from massive progenitors. For example, they would become N-rich, but would not C-rich if they have evolved from 0.8 single stars with [Fe/H]–2.3 (=10), according to the current stellar evolution models (e.g., Fujimoto et al. 2000). To become C-rich PNe, the third dredge-up (TDU) must take place in the late AGB phase. The efficiency of the TDU depends on the initial mass and composition, with increasing efficiency in models of increasing mass, or decreasing metallicity. At halo metallicities, it is predicted that the TDU is efficient in stars with initial masses greater than 1 (Karakas 2010; Stancliffe 2010). Also, current stellar evolutionary models predict that the post-AGB evolution of a star with an initial mass 0.8 proceeds too slowly for a visible PN to be formed. The origin and evolution of halo PNe are still one of the unresolved big problems in this research field.
How did these progenitor stars become visible C- and N-rich halo PNe? To answer this key question would deepen understanding of low-mass star evolution, in particular, extremely metal-poor C-rich stars found in the Galactic halo, and Galactic chemical evolution at early phases. If we can accurately estimate elemental abundances and ejected masses, then we can directly estimate elemental yields synthesized by PNe progenitors which might provide a constraint to the growth-rate of core mass, the number of thermal pulses and dredge-up mass. Hence we can build realistic stellar evolution models and Galactic chemical evolution models. We have observed these extremely metal-poor C- and N-rich halo PNe using the Subaru/High-Dispersion Spectrograph (HDS) and we also utilized collecting archival data carefully in order to revise our picture of these objects. In this paper, we focus on BoBn 1.
The known nebular and stellar parameters of BoBn 1 are listed in Table 1. Zijlstra et al. (2006) have associated BoBn 1 with the leading tail of the Sagittarius (Sgr) Dwarf Spheroidal Galaxy, which traces several halo globular clusters. The heliocentric distance to the Sgr dwarf Galaxy is 24.8 kpc (Kunder & Chaboyer 2009), while the distance to this object is between 16.5 (Henry et al. 2004) and 29 kpc (Kingsburgh & Barlow 1992).
BoBn 1 is a unique PN in that it might possess information about the chemical building-up history of the Galactic halo. Otsuka et al. (2008a) found that the [C/Fe] and [N/Fe] abundances of BoBn 1 are compatible with those of carbon-enhanced metal-poor (CEMP) stars. The C and N overabundances of CEMP can be explained by theoretical binary interaction models (e.g., Komiya et al. 2007; Lau et al. 2007). Otsuka et al. (2008a) detected two fluorine (F) lines and found that BoBn 1 is the most F-enhanced and metal-poor PN among F-detected PNe. They found that the C, N, and F overabundances of BoBn1 are comparable to those of the CEMP star HE1305+0132 (Schuler et al. 2007). Through a comparison between the observed enhancements of C, N, and F with the theoretical binary nucleosynthesis model by Lugaro et al. (2008), they concluded that BoBn 1 might share its origin and evolution with CEMP-s stars such as HE1305+0132, and if that is the case the slow neutron capture process (the s-process) should be considered.
According to current evolutionary models of low- to intermediate-mass stars, the s-process elements are synthesized by slowly capturing neutrons during the thermal pulse AGB phase. The s-process elements together with carbon are brought to the stellar surface by the TDU. If we could find signatures that BoBn 1 has experienced binary evolutions such as mass transfer from a massive companion and coalescence, the issues on chemical abundances and evolutionary time scale would be simultaneously resolved. It would also be of great significance to reveal the origin of these elements in the early Galaxy through the study of metal-poor objects such as BoBn 1. We will search s-process elements and investigate their enhancement in BoBn 1.
In this paper, we present a chemical abundance analysis of BoBn 1 using the newly obtained Subaru/HDS spectra, ESO VLT/UVES, Spitzer/IRS and IUE archive data. We detect several candidate collisional excited lines (CELs) of s-process elements and optical recombination lines (ORLs) of N, O, and Ne. We determine ionic and chemical abundances of 13 elements using ORLs and CELs. We construct a detailed photo-ionization model to derive the properties of the central star, ionized nebula, and dust. We also check consistency between our abundance estimations and the model. Finally, we compare the empirically derived abundances with the theoretical nucleosynthesis model values and discuss evolutionary scenarios for BoBn 1.
|Name||BoBn 1 (PN G108.4–76.1)||discovered by Boeshaar & Bond (1977)|
|Position (J2000.0)||=00:37:16.03 =–13:42:58.48|
|Distance (kpc)||22.5;29||Hawley & Miller (1978);Kingsburgh & Barlow (1992)|
|18.2;16.5||Mal’kov (1997);Henry et al. (2004)|
|24.8||Kunder & Chaboyer (2009)|
|Size (arcsec)||2 (diameter)||This work (see Fig. 1)|
|(erg cm s)||–12.54;–12.43||Cuisinier et al. (1996);Wright et al. (2005)|
|–12.38;–12.53(observed);–12.44(de-redden)||Kwitter et al. (2003);This work;This work|
|0.18;0.0;0.09||Cahn et al. (1992);Kwitter et al. (2003);This work|
|Rad. Velocity (km s)||191.6 (heliocentric)||This work|
|Exp. Velocity ()||See Table 6||This work|
|(K)||See Table 8||This work|
|(cm)||See Table 8||This work|
|Abundances||See Table 19||therein Table 19|
|3.57;3.72;3.07||Mal’kov (1997);Zijlstra et al. (2006);This work|
|(cm s)||5.52;6.5||Mal’kov (1997);This work|
|(M)||0.575;0.62||Mal’kov (1997);This work|
|(K)||125 000;96 300||Howard et al. (1997);Mal’kov (1997)|
|125 260||This work|
|Magnitude||16(B),14.6(R),16.13(J),15.62(H),15.18(K)||Simbad data base|
2. Data & Reductions
2.1. Subaru/HDS observations
The spectra of BoBn 1 were taken using the High-Dispersion Spectrograph (HDS; Noguchi et al. 2002) attached to one of the two Nasmyth foci of the 8.2-m Subaru telescope atop Mauna Kea, Hawaii on October 6th 2008 (program ID: S08B-110, PI: M.Otsuka). In Fig. 1, we present the optical image of BoBn 1 taken by the HDS silt viewer camera ( pixel, no filters) during the HDS observation. The sky condition was clear and stable, and the seeing was between and . The FWHM of the image is 1. BoBn 1 shows a small protrusion toward the southeast.
Spectra were taken for two wavelength ranges, 3600-5400 Å (hereafter, the blue region spectra) and 4600-7500 Å (the red region spectra). An atmospheric dispersion corrector (ADC) was used to minimize the differential atmospheric dispersion through the broad wavelength region. In these spectra, there are many recombination lines of hydrogen, helium, & metals and collisionally excited lines (CELs). These numerous spectral lines allowed us to derive reliable chemical compositions. We used a slit width of (0.6 mm) and a 22 on-chip binning, which enabled us to achieve a nominal spectral resolving power of =30 000 with a 4.3 binned pixel sampling. The slit length was set to avoid overlap of the echelle diffraction orders at the shortest wavelength portion of the observing wavelength range in each setup. This corresponds to (4.0mm), in which the nebula fits well and can allow us to directly subtract sky background from the object frames. The CCD sampling pitch along the slit length projected on the sky is per binned pixel. The achieved S/N is 40 at the nebular continuum level even in both ends of each Echelle order. The resulting resolving power is around 33 000, which derived from the mean of the full width at half maximum (FWHM) of narrow Th-Ar and night sky lines. All the data were taken as a series of 1800 sec exposure for weak emission-lines and 300 sec exposures for strong emission-lines. The total exposure times were 16 200 sec for red region spectra and 7200 sec for blue region spectra. During the observation, we took several bias, instrumental flat lamp, and Th-Ar comparison lamp frames, which were necessary for data reduction. For the flux calibration, blaze function correction, and airmass correction, we observed a standard star HR9087 at three different airmass.
2.2. VLT/UVES archive data
We also used archival high-dispersion spectra of BoBn 1, which are available from the European Southern Observatory (ESO) archive. These spectra were observed on August 2002 (program ID: 069.D-0413, PI: M.Perinotto) and June 2007 (program ID: 079.D-0788, PI: A.Zijlstra), using the Ultraviolet Visual Echelle Spectrograph (UVES; Dekker et al. 2000) at the Nasmyth B focus of KUEYEN, the second of the four 8.2-m telescopes of the ESO Very Large Telescope (VLT) at Paranal, Chile. We call the August 2002 data “UVES1” and the June 2007 data “UVES2” hereafter. We used these data to compensate for unobserved spectral regions and order gaps in the HDS spectra. We normalized these data to the HDS spectra using the intensities of detected lines in the overlapped regions between HDS and UVES1 & 2.
These archive spectra covered the wavelength range of 3300-6600 Å in UVES1 and 3300-9500 Å in UVES2. The entrance slit size in both observations was 11 in length and 1 in width, giving 30 000 derived from Th-Ar and sky lines. The CCDs used in UVES have 15 m pixel sizes. For UVES1 an 11 binning CCD pattern was chosen. For UVES2 a 22 on-chip binning pattern was chosen. The sampling pitch along the wavelength dispersion was 0.015-0.02 Å pixel for UVES1 and 0.03-0.04 Å pixel for UVES2. The exposure time for UVES1 was 2700 sec 4 frames, 10 800 sec in total. The exposure time for UVES2 was 1500 sec 2 frames, 3000 sec in total. The standard star Feige 110 was observed for flux calibration.
In Fig. 2, we present the combined HDS and UVES spectrum of BoBn 1 normalized to the H flux. The spectrum for the wavelength region of 3650-7500 Å is from the HDS data and that of 3450-3650 Å & 7500 Å is from the UVES data.
The observation logs are summarized in Table 2. The detected lines in the Subaru/HDS and VLT/UVES spectra are listed in Appendix A.
2.3. Iue archive data
We complemented optical spectra with UV spectra obtained by the International Ultraviolet Explorer (IUE) to derive C, C, N, and N abundances from semi-forbidden lines C ii], C iii], N iii], and N iv], since these emission lines cannot be observed in the optical region. These IUE spectra were retrieved from the Multi-mission Archive at the STScI (MAST). We collected the high- and low-resolution IUE spectra taken by the Short Wavelength Prime (SWP) and Long Wavelength Prime/Long Wavelength Redundant (LWP/LWR) cameras. Our used data set is listed in Table 3, and the wavelength dispersion mode is indicated in the column 3. All the IUE observations were made using the large aperture (10.323 arcsec). SWP and LWP/LWR spectra cover the wavelength range of 1150-1980 Å and 1850-3350 Å, respectively. For each SWP and LWP/LWR spectra, we did median combine to improve the S/N. The combined short wavelength spectrum was used to measure fluxes of emission-lines in 1910 Å because this allowed us to separate C iii1906/08 and C iv 1548/51 lines. C iii1906/08 are important as a density diagnostic. The combined long wavelength spectrum was for measurements of emission-line fluxes in 2000 Å. The measured line fluxes were normalized to the H flux using theoretical ratios of He ii (1640)/(4686) for the short wavelength spectrum and (2512)/(4686) for the long wavelength spectrum, respectively, adopting an electron temperature = 8840 K and density = 10 cm as given by Storey & Hummer (1995), then normalized to the H flux. The interstellar extinction correction was made using equation (1) (see section 3.1). The observed and normalized fluxes of detected lines are listed in the columns 4 and 5 of Table 4, respectively.
2.4. Spitzer archive data
We used two data sets (program IDs: P30333 PI: A.Zijlstra; P30652 PI: J.Bernard-Salas) taken by the Spitzer space telescope in December 2006. The data were taken by the Infrared Spectrograph (IRS, Houck et al. 2004) with the SH (9.5-19.5 m), LH (5.4-37m), SL (5.2-14.5 m) and LL (14-38m) modules. In Fig. 3 we present the Spitzer spectra of BoBn 1. We downloaded these data using Leopard provided by the Spitzer Science Center. The one-dimensional spectra were extracted using Spice version c15.0A. We extracted a region within 1 from the center of each spectral order summed up along the spatial direction. For SH and LH spectra, we subtracted sky background using off-set spectra. We normalized the SL and LL data to the SH and LH using the measured fluxes of [S iv] 10.5m, H i 12.4 m, [Ne ii] 12.8m, [Ne iii] 15.6m, and [Ne iii] 36.0m. Finally, the measured line fluxes were normalized to the H flux. The observed line ratio H i (11.2m)/(4861) (3.110) is consistent with the theoretical value (3.1510) for = 8840 K and = 10 cm as given by Storey & Hummer (1995). We did not therefore perform interstellar extinction correction.
The observed and normalized fluxes of detected lines are listed in Table 5. In addition to the ionized gas emissions, the amorphous carbon dust continuum and the polycyclic aromatic hydrocarbons (PAHs) feature around 6.2, 7.7, 8.7, and 11.2 m are found for the first time. In Fig. 4, we present these PAH features. The 11.2 m emission line is a complex of the narrow width H i 11.2 m and the broad PAH 11.2 m. The 6.2, 7.7, and 8.7 m bands emit strongly in ionized PAHs, while the 11.2 m does in neutral PAHs (Bernard-Salas et al. 2009). According to the PAH line profile classifications by Peeters et al. (2002) and van Diedenhoven et al. (2004), BoBn 1’s PAH line-profiles belong to class B. Bernard-Salas et al. (2009) classified 10 of 14 Magellanic Clouds (MCs) PNe into class B based on Spitzer spectra. In measuring PAH band fluxes, we used local continuum subtracted spectrum by a spline function fitting. We followed Bernard-Salas et al. (2009) and measured integrated fluxes between 6.1 and 6.6 m for the 6.2 m PAH band, 7.2-8.3 m for the 7.7 m PAH, 8.3-8.9 m for the 8.6 m PAH, and 11.1-11.7 m for the 11.2 m. The observed PAH flux ratios (6.2m)/(11.2m) and (7.7m)/(11.2m) follow a correlation among MCs PNe, shown in Fig. 2 of Bernard-Salas et al. (2009). BoBn 1 has a hot central star (10 K), so that ionized PAH might be dominant. However, these line ratios of BoBn 1 are somewhat lower than those of excited MCs PNe. One must take a look at a part of the neutral PAH emissions in a photodissociation region (PDR), too.
We found a plateau between 10 and 14 m, which are believed to be related to PAH clusters (Bernard-Salas et al. 2009). Meanwhile, MgS feature around 30 m sometimes observed in C-rich PNe, was unseen in BoBn 1.
2.5. Data reduction
Data reduction and emission line analysis were performed mainly with
a long-slit reduction package noao.twodspec in
First, we made a zero-intensity level correction to all frames
including flat lamp, object, and Th-Ar comparison frames
using the overscan region of each frame and the mean bias frames.
We also removed cosmic ray events and hot pixels from the object frames.
Second, we trimmed the overscan region and removed scattered light
from the flat lamp and object frames. Third, we made a CCD sensitivity
correction to the object frames using the median flat frames.
Fourth, we extracted a two-dimensional spectrum from each echelle
diffraction order of each object frame and made a wavelength calibration
using at least two Th-Ar frames taken before and after the object frame.
We referred to the Subaru/HDS comparison
In measuring emission line fluxes, we assumed that the line profiles were all Gaussian and we applied multiple Gaussian fitting techniques.
|(Å)||(erg s cm)||(H)=100|
|1485||N iv]||1.306||1.41(–13) 5.47(–14)||45.81 17.84|
|1548||C iv||1.239||3.27(–12) 5.04(–14)||1052.6 20.22|
|1551||C iv||1.237||1.62(–12) 4.27(–14)||519.3 14.96|
|1640||He ii||1.177||5.13(–13) 7.40(–14)||162.97 23.57|
|1750||N iii]||1.154||1.52(–13) 4.16(–14)||48.06 13.16|
|1906||C iii]||1.255||2.60(–12) 2.50(–14)||838.88 12.65|
|1908||C iii]||1.258||1.87(–12) 2.33(–14)||602.72 10.28|
|2324||C ii]||1.388||2.87(–13) 3.60(–14)||36.64 4.63|
|2424||[Ne iv]||1.134||1.41(–13) 1.46(–14)||17.12 1.78|
|2512||He ii||0.969||2.66(–14) 1.29(–14)||3.13 1.52|
Note. – () stands for 10. (H) = 3.63(–13) 6.47(–14) (see section 3.1).
|(m)||(erg s cm)||(H)=100|
|6.2||PAH||2.62(–14) 1.52(–15)||7.22 1.35|
|7.7||PAH||8.13(–14) 2.31(–15)||22.39 4.04|
|8.6||PAH||2.02(–14) 1.23(–15)||5.57 1.05|
|10.5||S iv||7.36(–15) 1.93(–16)||1.92 0.05|
|11.3||H i||1.18(–15) 1.81(–16)||0.31 0.05|
|11.3||PAH||4.97(–14) 9.65(–15)||13.68 3.61|
|12.4||H i||3.89(–15) 3.01(–16)||1.02 0.08|
|12.5||He i||1.72(–15) 4.17(–16)||0.45 0.11|
|12.8||Ne ii||9.53(–15) 3.13(–16)||2.49 0.08|
|14.6||He i||5.85(–16) 2.40(–16)||0.15 0.06|
|15.6||Ne iii||6.17(–13) 4.98(–15)||161.13 1.30|
|16.4||He i||2.19(–15) 2.65(–16)||0.57 0.07|
|18.7||S iii||2.65(–15) 1.79(–16)||0.69 0.05|
|19.1||H i||1.14(–15) 2.64(–16)||0.30 0.07|
|25.9||O iv||4.77(–14) 5.20(–16)||12.46 0.14|
|36.0||Ne iii||5.10(–14) 7.04(–15)||13.32 1.84|
Note. – () stands for 10. (H) = 3.63(–13) 6.47(–14) (see section 3.1).
3.1. Interstellar reddening correction
We have detected over 600 emission lines in total. Before proceeding to the chemical abundance analysis, it is necessary to correct the spectra for the effects of absorption due to the Earth’s atmosphere and interstellar reddening. The former was performed using experimental functions measured at the Keck observatories and the ESO/VLT. The interstellar reddening correction was made by determining the reddening coefficient at H, (H). We fitted the observed intensity ratio of H to H with the theoretical ratios computed by Storey & Hummer (1995). Two different situations are assumed for some lines: Case A assumes that the nebula is transparent to the lines of all series of hydrogen; Case B assumes the nebula is partially opaque to the lines of the Lyman series but is transparent for the Balmer series of hydrogen. Initially, we assumed that = 10 K and = 10 cm in Case B, and we estimated (H) = 0.087 0.004 from the HDS spectra. That is an intermediate value between Cahn et al. (1992) and Kwitter et al. (2003) (see Table 1). For the UVES spectra, we estimated (H) = 0.066 by the same manner. From Seaton’s (1979) relation (H) = 1.47 one obtains = 0.06 for the HDS spectra and 0.04 for UVES spectra, which are comparable to the Galactic value (0.02) to the direction to BoBn 1 measured by the Galactic extinction model of Schlegel et al. (1998).
All of the line intensities were then de-reddened using the formula:
where is the de-reddened line flux; is the observed line flux; and is the interstellar extinction at . We adopted the reddening law of Cardelli et al. (1989) with the standard value of = 3.1 for . We observed = 2.5610 2.1310 erg s cm (() stands for 10, hereafter) within the slit in the HDS observation. We estimated captured light from BoBn 1 using the image presented in Fig. 1, to be about 86.8 of the light from BoBn 1 in the HDS observation. The intrinsic observed H flux is 2.95(–13) 3.45(–16) erg s cm and the de-reddened H flux is 3.63(–13) 6.47(–14) erg s cm including the error of (H).
3.2. Radial and expansion velocities
We present the line-profiles of selected ions in Fig. 5. The observed wavelength at the time of observation was corrected to the averaged line-of-sight heliocentric radial velocity of +191.60 1.25 km s among over 300 lines detected in the HDS spectra. The line-profiles can be represented by a single Gaussian for weak forbidden lines such as [Ne v] 3426 and metal recombination lines such as O ii 4642. For the others, the profiles can be represented by the sum of two or three Gaussian components.
Most of the detected lines are asymmetric profile, in particular the profiles of low-ionization potential ions show strong asymmetry. The asymmetric line-profiles are sometimes observed in bipolar PNe having an equatorial disk structure. The similar line-profiles are also observed in the halo PN H4-1 (Otsuka et al. 2006). H4-1 has an equatorial disk structure and multi-polar nebulae. The elongated nebular shape of BoBn 1 (Fig. 1) might indicate the presence of such an equatorial disk. The receding ionized gas (especially, low-ionization potential ions) from the observers around the central star would be strongly weakened by the equatorial disk. In contrast, the relatively large extent bipolar flows perpendicular to the equatorial disk might be unaffected by the disk. Due to such a geometry, we observe asymmetric line-profiles.
In Table 6, we present twice the expansion velocity 2 measured from selected lines. When we fit line-profile with two or three Gaussian components, we define that 2 corresponds to the difference between the positions of the red and blue shifted Gaussian peak components. We call expansion velocity measured by this method ’2(a)’.
When we can fit line-profile with single Gaussian, we determine 2 from following equation,
where is the velocity FWHM of each Gaussian component. and (9 km s) are the thermal broadening and the instrumental broadening, respectively (Robinson et al. 1982). is represented by 21.4, where is the electron temperature (in units of 10 K) and is the atomic weight of the target ion. For CELs, we adopted listed in Table 9. For ORLs, we adopted = 0.88. We call twice the expansion velocity measured by Equation (2) ’2(b)’. We converted 2(a) into 2(b) using the relation 2(b) = (1.74 0.12) 2(a) for 2(b) 16 km s, which can be applied only to BoBn 1. We adopted 2(b) as twice the expansion velocities for BoBn 1. The averaged (b) is 40.5 3.28 km s among selected lines listed in Table 6 and 33.04 2.61 km s among over 300 detected lines in the HDS spectra.
In Fig. 6 we present relation between 2(b) and ionization potential (I.P.). When we assume that BoBn 1 has a standard ionized structure (i.e., high I.P. lines are emitted from close regions to the central star and low I.P. lines are from far regions), expansion velocity of BoBn 1 seems to be proportional to the distance from the central star. BoBn 1 might have Hubble type flows. We found that 2(b) values from ORLs are slightly smaller than CELs with the same I.P., for example, O ii & [O iii] (35.5 eV) and Ne ii & [Ne iii] (41.0 eV). O ii and Ne ii might be emitted from colder regions than [O iii] and [Ne iii] do.
|(Å)||(eV)||(km s)||(km s)|
Note. – The probable error of (b) is within 5 km s. We assume 2(b) as twice the expansion velocity of BoBn 1 (see text).
|Line||transition probabilities||collisional strength|
|C i||Froese-Fischer & Saha (1985)||Johnson et al. (1987); Péquignot & Aldrovandi (1976)|
|C ii||Nussbaumer & Storey (1981); Froese-Fischer (1994)||Blum & Pradhan (1992)|
|C iii||Wiese et al. (1996)||Berrington et al. (1985)|
|C iv||Wiese et al. (1996)||Badnell & Pindzola (2000); Martin et al. (1993)|
|[N i]||Wiese et al. (1996)||Péquignot & Aldrovandi (1976); Dopita et al. (1976)|
|[N ii]||Wiese et al. (1996)||Lennon & Burke (1994)|
|N iii||Brage et al. (1995); Froese-Fischer (1983)||Blum & Pradhan (1992)|
|N iv||Wiese et al. (1996)||Ramsbottom et al. (1994)|
|[O i]||Wiese et al. (1996)||Bhatia & Kastner (1995)|
|[O ii]||Wiese et al. (1996)||McLaughlin & Bell (1993); Pradhan (1976)|
|[O iii]||Wiese et al. (1996)||Lennon & Burke (1994)|
|[O iv]||Wiese et al. (1996)||Blum & Pradhan (1992)|
|F ii||Storey & Zeippen (2000); Baluja & Zeippen (1988)||Butler & Zeippen (1994)|
|F iii||Naqvi (1951)||See Text|
|F iv||Garstang (1951); Storey & Zeippen (2000)||Lennon & Burke (1994)|
|Ne ii||Saraph & Tully (1994)||Saraph & Tully (1994)|
|[Ne iii]||Mendoza (1983); Kaufman & Sugar (1986)||McLaughlin & Bell(2000)|
|[Ne iv]||Becker et al. (1989); Bhatia & Kastner (1988)||Ramsbottom et al. (1998)|
|[Ne v]||Kaufman & Sugar (1986); Bhatia & Doschek (1993)||Lennon & Burke (1994)|
|[S ii]||Verner et al. (1996); Keenan et al. (1993)||Ramsbottom et al. (1996)|
|[S iii]||Tayal & Gupta (1999)||Froese Fischer et al. (2006)|
|[S iv]||Johnson et al. (1986); Dufton et al. (1982); Verner et al. (1996)||Dufton et al. (1982)|
|[Cl iii]||Mendoza & Zippen (1982a); Kaufman & Sugar (1986)||Ramsbottom et al. (2001)|
|[Cl iv]||Mendoza & Zippen (1982b); Ellis & Martinson (1984)||Galavis et al. (1995)|
|Kaufman & Sugar (1986)|
|[Ar iii]||Mendoza (1983); Kaufman & Sugar (1986)||Galavis et al. (1995)|
|[Ar iv]||Mendoza & Zippen (1982a); Kaufman & Sugar (1986)||Zeippen et al. (1987)|
|[Fe iii]||Garstang (1957); Nahar & Pradhan (1996)||Zhang (1996)|
|[Fe iv]||Froese-Fischer & Rubin (1998); Garstang (1958)||Zhang & Pradhan (1997)|
|[Kr iv]||Biémont & Hansen (1986)||Schöning (1997)|
|[Kr v]||Biémont & Hansen (1986)||Schöning (1997)|
|Rb v||Persson et al. (1984)|
|[Xe iii]||Biémont et al. (1995)||Schöning & Butler (1998)|
|Ba ii||Klose et al. (2002)||Schöning & Butler (1998)|
3.3. Plasma diagnostics
We have detected a large number of collisionally excited lines (CEL), useful for estimations of the temperatures () and densities (). The electron temperature and density diagnostic lines analyzed here arise from various ions, which have a wide variety of ionization potentials ranging from 0 ([N i] & [O i]) to 63.5 eV ([Ne iv]). We have examined the electron temperature and density structure within the nebula of BoBn 1 using 17 diagnostic line ratios. The [O i], [Ne iii], [Ne iv], [S ii], and [S iii] zone electron temperatures and [N i], C iii] and [Ne iii] zone electron densities are estimated for the first time. Electron temperatures and densities were derived from each diagnostic ratio for each line by solving level populations for a multi-level ( 5 for almost all the ions) atomic model using the collision strengths () and spontaneous transition probabilities for each ion from the references given in Table 7.
The derived electron temperatures and densities are listed in Table 8. Fig. 7 is the diagnostic diagram that plots the loci of the observed diagnostic line ratios on the – plane. This diagram shows that most CELs in BoBn 1 are emitted from 12 000–16 000 K and 3.5 cm ionized gas.
First, we calculated electron densities assuming a constant electron temperature of 12 800 K. Estimated electron densities range from 1030 ([N i]) to 5740 cm ([S ii]). Although [S ii] and [O ii] have similar ionization potentials, a large discrepancy between their electron densities is found (see Table 8). Kniazev et al. (2008) and Kwitter et al. (2003) estimated [S ii] electron densities as large as 9600 and 7100 cm, respectively. Stanghellini & Kaler (1989), Copetti & Writzl (2002), and Wang et al. (2004) found that the [S ii] density is systematically larger than the [O ii] density in a large number of samples. The curve yielded by the [S ii]6716/31 ratio in the vs plane indicates higher electron density than critical density of these lines, 1600 4100 cm at = 12 800 K for 6716 and 6731, respectively (cf. 5740 cm in Fig. 7). This density discrepancy is not due to the errors in the [O ii] atomic data. Wang et al. (2004) also found the density discrepancy between [S ii] and [O ii] that might be likely caused by errors in the transition probabilities of [O ii] given by Wiese et al. (1996). In the case of BoBn 1, this possibility can be ruled out because we obtained similar [O ii] electron densities even when with the other transition probabilities. The high [S ii] density might be due to high-density blobs in the outer nebula. This could have contributed to producing the small [S ii]6717/6731 ratio and give rise to an apparently high density. Zhang et al. (2005a) pointed out the possibility that a dynamical plow by the ionization front effects yields large density of [S ii] because the ionization potential of S is close to the H edge. Since the estimated upper limit to the [S ii] density is close to the H density derived from the Balmer decrement (see below), this explanation might be plausible. In BoBn 1, caution is necessary when using the [S ii] electron density.
Next, we calculated the electron temperature. An average electron density of 3370 cm, which excluded the ([N i]) and ([S ii]), was adopted when estimating electron temperatures except the ([O i]). [N i] and [O i] are representative of the very outer part of the nebula and probably do not coexist with most of the other ions. ([O i]) was, therefore, estimated using ([N i]).
To obtain the [N ii], [O ii] and [O iii] temperatures it is necessary to subtract the recombination contamination to the [N ii] 5755, [O ii] 7320/30, and [O iii] 4363 lines, respectively. For [N ii] 5755, Liu et al. (2000) estimated the contamination to [N ii] 5755, ([N ii]5755) in the range 5000 20 000 K as
where N/H is the doubly ionized nitrogen abundance. Adopting the value derived from the ORL analysis (see Section 3.6) and using Equation (3), we estimated ([N ii]5755) 0.1, which is approximately 7 of the observed value. Given the corrected [N ii] 5755 intensity, the [N ii] temperature is 12 000 K, which is 400 K lower than that obtained without taking into account the recombination effect.
The same effect also exists for the [O ii] 7320/30 lines. We estimated the recombination contribution using the doubly ionized oxygen abundance derived from O ii lines and the equation of Liu et al. (2000) for these lines in the range 5000 10 000 K,
Using Equation (4), we estimated a contribution of 7 of the observed value and obtained 12 100 K, lower by 700 K than that without the recombination contribution. For [O iii] 4363, we estimated the recombination contribution using the O abundance derived from the fine-structure line [O iv] 25.9m adopting ([Ne iv]) and ([Ar iv]) and the equation of Liu et al. (2000). We chose the value from this line because O iii lines could be affected by star light excitation and the abundance derived from them could be erroneous. Assuming the ratio of O(ORLs)/O(CELs) = 10, we estimated the recombination contribution to [O iii] 4363 less than 1 of the observed value, which has a negligible effect on the ([O iii]) derivation.
The electron temperature in BoBn 1 ranges from 9520 ([O i]) to 14 920 K ([Ne iv]). Our estimated electron temperatures except for [O ii] are comparable to those of Kwitter et al. (2003) and Kniazev et al. (2008). Their estimated temperatures are 12 400-13 720 K for [O iii], 11 320-11 700 K for [N ii], and 13 250 K for [Ar iii]. Note that the [O ii] electron temperature of 8000 K of Kwitter et al. (2003) was estimated adopting the [S ii] density of 7100 cm. Our – plane predicts that the [O ii] temperature is 10 000 K when adopting [S ii] density.
The ionic abundances derived from the CELs depend strongly on the electron temperature. In the case of [O iii]5007, for example, only 500 K change makes a difference of over 10 for O abundance. It is therefore essential to find the proper electron temperature for each ionized stage of each ion. To that end, we examined the behavior of the electron temperature and density as a function of I.P. The upper panel of Fig. 8 shows that is increasing proportional to I.P. The observed behavior of is consistent with the schematic picture of stratified physical conditions in ionized nebula, where the electron temperature of ions in the inner part should be hotter than that in the outer part. is simply monotonically increasing up to 40 eV as I.P., except for [S ii]. The [S ii] might be emitted in high-density blobs in the outer nebula as we mentioned above.
To minimize the estimated error for ionic abundances due to electron temperature, we have assumed a 7-zone model for BoBn 1 by reference to Fig. 8. Adopted and for each ion are presented in Table 9. ([O i]) and ([N i]) are adopted for ions in zone 0, which have 10 eV. ([N ii]) and ([S ii]) are adopted for ions in zone 1 (I.P. 11.3 eV). ([N ii]) and ([O ii]) are for zone 2 (11.3-20 eV). ([S iii]) and (C iii]) are for zone 3 (20-25 eV). For zone 4 (25-41 eV), 5 (41-63.5 eV), and zone 6 ( 63.5 eV), we adopted ([Ar iv]) and ([O iii]), the averaged value from ([O iii]) and ([Ne iv]), and ([Ne iv]), respectively.
|(1)||[Ne iv] (2422+2425)/(4715/16/25/26)||100.7810.84||14 920810|
|(K)||(2)||[O iii] (4959+5007)/(4363)||85.014.32||13 650290|
|(3)||[Ar iii] (7135)/(5192)||85.7035.63||13 3303 310|
|(4)||[Ne iii] (15.5m)/(3869+3967)||0.200.01||13 050140|
|(5)||[Ne iii] (3869+3967)/(3344)||333.6714.63||12 870170|
|(6)||[N ii] (6548+6583)/(5755)||57.551.77||12 000190|
|(7)||[S iii] (9069)/(6312)||7.420.53||12 460490|
|(8)||[O i] (6300+6363)/(5577)||68.6510.76||9520550|
|(9)||[O ii] (3726+3729)/(7320+7330)||10.470.21||12 100180|
|(10)||[S ii] (6716+6731)/(4069+4076)||13.845.17||12 420|
|He i (7281)/(6678)||0.210.01||9430310|
|He i (7281)/(5876)||0.050.01||7340110|
|He i (6678)/(4471)||0.830.02||7400|
|He i (6678)/(5876)||0.270.01||9920310|
|(Balmer Jump)/(H 11)||8840210|
|(11)||[N i] (5198)/(5200)||1.430.03||1030130|
|(cm)||(12)||[O ii] (3726)/(3729)||1.650.03||151060|
|(13)||C iii] (1906)/(1909)||1.390.03||35901000|
|(14)||[Ar iv] (4711)/(4740)||1.050.07||39601090|
|(15)||[Ne iii] (15.5m)/(36.0m)||12.111.69||4400|
|(16)||[S ii] (6716)/(6731)||8.570.03||57401 310|
|Balmer decrement||5000-10 000|
We detected a large number of optical recombination lines (ORL). C iii,iv, O ii,iii,iv, N ii,iii, and Ne ii are for the first time detected. To calculate ORL abundances, the electron temperature and density derived from ORLs are needed. We estimated the electron temperature using the Balmer discontinuity and He i line ratios, and the electron density using the Balmer decrement. The results are listed in Table 8.
The ratio of the jump of continuum emission at the Balmer limit at 3646Å (BJ) to a given hydrogen emission line depends on the electron temperature. Following Liu et al. (2001), we use this ratio to determine the electron temperature. This temperature, (BJ), is used to deduce ionic abundances from ORLs. Defining BJ as , and taking the emissivities of H i Balmer lines and H i, He i, and He ii continuum emissivities, Liu et al. (2001) gave the following equation:
[ is in units of Å. (BJ) is valid over a range from 4000 to 20 000 K. The process was repeated until self-consistent values for the N(He)/N(H), N(He)/N(H) and (BJ) were reached, we estimated (BJ) of 8840 K.
We estimated the He i electron temperature (He i) from the ratios of He i7281/6678, 7281/5876, 6678/4471, and 6678/5876 assuming a constant electron density = 10 cm, estimated from the Balmer decrement as described below. All the He i line ratios we chose here are insensitive to the electron density. We adopted the emissivities of He i from Benjamin et al. (1999). We estimated (He i) values as 7340–9920 K. The (He i) from 7281/6678 ratio appears to be the most reliable value because (i) He i 6678 and 7281 levels have the same spin as the ground state and the Case B recombination coefficients for these lines by Benjamin et al. (1999) are more reliable than the other He i 4471 and 5876; (ii) the effect of interstellar extinction is less due to the close wavelengths. We adopted 9430 K as (He i). Note that in Fig. 7 the electron temperatures and densities derived from the H i and He i are not presented.
The intensity ratios of the high-order Balmer lines Hn (n 10, n: the principal quantum number of the upper level) to a lower Balmer line, e.g., H, are also sensitive to the electron density. In Fig. 9, we plot the ratio of higher-order Balmer lines to H with the theoretical values by Storey & Hummer (1995) for the cases of electron temperature of 8840 K (= (BJ)) & electron densities of 1000, 5000, 10, and 10 cm. This diagram indicates that the electron density in the ORL emitting region is between 5000 and 10 cm, which is fairly compatible with the CEL electron densities. Zhang et al. (2005b) estimated (He i) from the ratio of He i for 48 PNe and found that high-density blobs (10–10 cm) might be present in nebula if (He i) (BJ). Fig. 9 indicates that such components do not coexist in BoBn 1.
3.4. Ionic abundances from CELs
The derived CEL ionic abundances X/H are listed in Table 10. X and H are the number densities of an times ionized ion and ionized hydrogen, respectively. To estimate ionic abundances, we solved level populations for a multi-level atomic model. In the last one of the line series of each ion, we present the adopted ionic abundances in bold face characters. These values are estimated from the line intensity-weighted mean or average if there are two or more available lines. Over 10 ionic abundances of some elements are estimated for the first time. These newly estimated ionic abundance would reduce the uncertainty of estimation of each elemental abundance, in particular, N, O, F, Ne, S, Fe, and some s-process elements, which are key elements to the nucleosynthesis in low-mass stars and chemical evolution in galaxies.
Ne (zone 4 ion) and S (zone 3) abundances are derived from CELs seen in both the UV-optical and mid-infrared regions. CELs in the mid-infrared, namely, fine-structure lines, have an advantage in derivations of ionic abundances. Since the excitation energy of the fine-structure lines is much lower than that of the other transition lines, ionic abundances derived from these lines are nearly independent of the electron temperature or temperature fluctuation in the nebula. Note that these ionic abundances derived from fine-structure lines are almost consistent with those from other transition lines. This means that the adopted electron temperature and density for the ions in zones 3 and 4 are appropriate at least.
Followings are short comments on derivations of ionic abundances. We subtract the [O iii] 2322 contamination from C ii 2324 using the theoretical intensity ratio [O iii] (2322)/(4363) of 0.24 and then estimate the C abundance. The N and O abundances are derived only from [N ii] 6548/83 and [O ii] 3726/29 to avoid recombination contamination, respectively, while The Ne abundance is calculated from [Ne ii] 12.8 m by solving a two level atomic model.
We have detected two fluorine lines [F iv] 3996,4060 and estimate a F abundance of 1.50(–8). Otsuka et al. (2008a) detected these fluorine lines and estimated a F abundance of 5.32(–8). The F abundance discrepancy between Otsuka et al. (2008a) and the present work is due to different adopted electron temperature and (H) values. We have detected candidates of [F ii] 4790,4868 and [F iii] 5721/33. In the previous section, we confirmed that BoBn 1 has no high-density components, larger than the critical density of these [F iii] lines. The critical density of [F ii] 4790/4868 is 210 cm, and that of [F iii] 5721/33 is 810 cm. Therefore, the effect of collisional de-excitation is negligibly small. Accordingly, the ratios of [F ii] (4790)/(4868) and [F iii] (5733)/(5721) depend on their transition probabilities. When adopting the transition probabilities by Baluja & Zeippen (1988) for [F ii] and Naqvi (1951) for [F iii], the theoretical intensity ratios of [F ii] (4790)/(4868) and [F iii] (5721)/(5733) are 3.2 and 1, which are in agreement with our measurements (4.21.0 for [F ii] and 1.00.3 for [F iii]). Hence, these four emission lines can be identified as [F ii] 4790,4868 and [F iii] 5721/33. The F and F abundances are estimated from the each detected line by solving the statistical equilibrium equations for the lowest five energy levels. For [F iii] lines the relevant collision strength has not been calculated. However, since F is isoelectronic with Ne, and collision strengths for the same levels along an isoelectronic sequence tend to vary with effective nuclear charge (Seaton 1958). We therefore assume that the collision strengths for [F iii] are 22 smaller than those for [Ne iv] and estimate F abundances. Otsuka et al. (2008a) showed a correlation between [Ne/Ar] and [F/Ar] in PNe, suggesting that Ne and F were synthesized in the same layer and carried to the surface by the third dredge-up. If this is the case, the ionic abundance ratios of F (I.P. = 35 eV) and F (62.7 eV) to F (17.4 eV) should be comparable to those of Ne (41 eV) and Ne (63.5 eV) to Ne (21.6 eV). Indeed, these ionic abundance ratios follow our prediction; F:F:F 1:34:1 and Ne:Ne:Ne 1:28:1. This means that our identifications of two [F ii] and [F iii] lines and the estimated ionic abundances are reliable. So far, fluorine has found in only a handful of PNe (see Zhang & Liu 2005; Otsuka et al. 2008a). Among them BoBn 1 appears to be the most F-rich PN.
We subtract the contribution to [Cl iii] 8500.2 due to C iii 8500.32 using the C ORL abundance and give upper limit of Cl abundance from this line. The adopted Cl abundance is from [Cl iii] 5517 only. The Ar abundance is from [Ar iv] 4717/40. We adopt the S abundance based on [S ii] lines except [S ii] 4068, because [S ii] 4068 could be partially blended with C iii 4068. To estimate Fe and Fe abundances, we solved a 33 level model (from to b) for [Fe iii] and a 18 level model (from to ) for [Fe iv]. We adopt the transition probabilities of [Fe iv] recommended by Froese-Fischer & Rubin (1998). For those not considered by Froese-Fischer & Rubin (1998), the values by Garstang (1958) were adopted.
3.5. Ionic abundances of heavy elements (Z 30)
We have detected 10 emission line candidates of krypton (Kr), rubidium (Rb), xenon (Xe), and barium (Ba). Kr and Rb are light- s-process elements (30 Z 40, Z: atomic number), Xe and Ba are heavy s-process (Z 41). Kr has been detected in over 100 PNe, while the latter three s-process elements have been detected in only a handful of PNe (Sharpee et al. 2007). Selected line profiles of these candidates are presented in Fig. 10. The Kr, Kr, Xe, and Ba abundances in this object are estimated for the first time.
We have detected two nebular lines of [Kr iv] 5346.7,5867.7 ( and , respectively). For [Kr iv] 5346.7, the possibility of blending with C iii 5345.85 (multiplet V13.01) is low, and the contribution to this [Kr iv] line is probably negligible because other V13.01 C iii lines are not detected. For [Kr iv] 5867.7, we estimated the contamination from He ii 5867.7 using theoretical ratios of He ii (5867.7) to (5828.4), (5836.5), (5857.3), (5882.12), and (5896.8) given by Storey & Hummer (1995) assuming = 8840 K and = 10 cm, from which the contribution from He ii 5867.7 is estimated to be 64 . In Fig. 10, we present the isolated [Kr iv] 5867.7 profile. The observed ratio of [Kr iv] (5867.7)/(5346.7) (1.8) is comparable to the theoretical value (1.5) if the population at each energy level does not exceed the critical densities of 1.5(+6) cm () and 2(+7) cm (). We therefore identified these two lines as [Kr iv] 5346.7,5867.7.
[Kr v] 8243.4 () is likely to be blended with the blue wing of the Paschen series H i 8243.7 (n=3-43). Using theoretical ratios of H i (8243.7) to (8247.7) and (8245.6) given by Storey & Hummer (1995), we subtracted the contribution of H i 8243.7, then estimated the intensity of [Kr v] 8243.4. We found another nebular line [Kr v] 6256.1 (). The theoretical intensity ratio of (6256.1)/(8243.7) (1.1) is in good agreement with ours (1.2).
[Xe iii] 5846.8 () appears to be blended with He ii 5846.7. We subtracted the He ii 5846.7 contribution from it using the theoretical ratios of (5846.7) to (5828.4), (5836.5), (5857.3), (5882.12), and (5896.8) given by Storey & Hummer (1995), then obtained an upper limit to the intensity of [Xe iii] 5846.8. In Fig. 10, we present the isolated [Xe iii] 5846.8 profile.
Two Ba ii recombination lines 4934,6141.7 ( and ) are detected. Following Sharpee et al. (2007), we estimated the Ba abundances adopting transition electron temperature and density (zone 0). The Ba abundances from these lines are in good agreement each other.
We have detected a candidate [Rb v] 5363.6 (auroral line; ). Rb is one of the important elements as tracers of the neutron density. In the case of NGC 7027, Sharpee et al. (2007) argued the possibility that this line is O ii 5363.8 (). They also suggested that the intensity of O ii 5363.8 is comparable to O ii 4609.4 (), arising from the lower level of O ii 5363.8. In BoBn 1, based on the ORL O abundance of 1.45(–4) from the 3d-4f O ii lines (see next Section) the expected intensity of O ii 4609.4 is 6.8(–5)(H), which is lower than the observed intensity of the [Rb v] 5363.6 candidate. No O ii lines are detected in BoBn 1. Therefore, we consider that the detected line is [Rb v] 5363.6. Since there are no available collision strengths for this line at present, we do not estimate a Rb abundance.
|C||8727.12||7.93(–2) 4.73(–3)||9520||1030||4.74(–7) 9.34(–8)|
|C||2324||3.51(+1) 4.46(0)||12000||5740||2.30(–5) 3.48(–6)|
|C||1906||8.39(+2) 1.27(+1)||12460||3590||7.71(–4) 1.59(–4)|
|1908||6.03(+2) 1.03(+1)||7.70(–4) 1.59(–4)|
|C||1548||1.05(+3) 2.02(+1)||14290||3960||2.42(–4) 4.47(–5)|
|1551||5.19(+2) 1.50(+1)||2.36(–4) 4.38(–5)|
|N||5197.90||2.73(–1) 7.90(–3)||9520||1030||4.81(–7) 1.01(–7)|
|5200.26||1.91(–1) 5.47(–3)||4.82(–7) 9.67(–8)|
|N||5754.64||1.23(0) 1.40(–2)||12000||1510||7.20(–6) 4.74(–7)|
|6548.04||1.56(+1) 8.34(–1)||5.84(–6) 3.75(–7)|
|6583.46||5.04(+1) 1.28(0)||6.40(–6) 2.79(–7)|
|N||1750||4.81(+1) 1.32(+1)||13650||3960||6.24(–5) 1.88(–5)|
|N||1485||4.58(+1) 1.78(+1)||14290||3960||3.83(–5) 1.65(–5)|
|O||5577.34||1.70(–2) 2.64(–3)||9520||1030||2.06(–6) 7.06(–7)|
|6300.30||8.72(–1) 1.55(–2)||2.04(–6) 3.97(–7)|
|6363.78||2.91(–1) 9.91(–3)||2.13(–6) 4.20(–7)|
|O||3726.03||1.09(+1) 6.85(–2)||12000||1510||4.00(–6) 2.23(–7)|
|3728.81||6.61(0) 9.99(–2)||4.03(–6) 2.42(–7)|
|7319||1.02(0) 2.00(–2)||7.13(–6) 5.67(–7)|
|7330||7.81(–1) 1.38(–2)||6.74(–6) 5.30(–7)|
|O||4363.21||5.57(0) 9.73(–2)||13650||3960||4.76(–5) 4.87(–6)|
|4931.23||4.36(–2) 3.88(–3)||4.37(–5) 4.58(–6)|
|4958.91||1.22(+2) 5.96(0)||4.80(–5) 3.52(–6)|
|5006.84||3.51(+2) 2.18(+1)||4.76(–5) 3.94(–6)|
|O||25.9||1.25(+1) 1.36(–1)||14290||3960||3.41(–6) 8.21(–8)|
|F||4789.45||5.58(–2) 5.20(–3)||12000||1510||2.16(–8) 2.23(–9)|
|4868.99||1.34(–2) 3.00(–3)||1.66(–8) 3.79(–9)|
|F||5721.20||2.70(–2) 2.93(–3)||13650||3960||6.59(–7) 1.03(–7)|
|5733.05||2.68(–2) 5.43(–3)||6.70(–7) 1.55(–7)|
|F||3996.92||4.09(–2) 2.53(–3)||14920||3960||1.47(–8) 2.21(–9)|
|4059.90||1.19(–1) 3.72(–3)||1.51(–8) 2.12(–9)|
|Ne||12.8||2.49(0) 8.17(–2)||12460||3590||2.97(–6) 1.14(–7)|
|Ne||3342.42||8.47(–1) 2.75(–2)||13650||3960||6.40(–5) 8.10(–6)|
|3868.77||2.17(+2) 1.05(+1)||8.41(–5) 6.75(–6)|
|3967.46||6.39(+1) 4.13(–1)||5.94(–5) 3.82(–6)|
|4011.60||1.48(–2) 3.92(–3)||9.71(–5) 2.65(–5)|
|15.6||1.61(+2) 1.30(0)||9.15(–5) 1.27(–6)|
|36||1.33(+1) 1.84(0)||9.07(–5) 1.26(–5)|
|Ne||2423.50||1.71(+1) 1.78(0)||14920||3960||3.97(–6) 8.99(–7)|
|4714.25||5.52(–2) 2.76(–3)||4.35(–6) 1.23(–7)|
|4715.80||1.87(–2) 2.12(–3)||5.05(–6) 1.52(–6)|
|4724.15||5.08(–2) 1.88(–3)||3.60(–6) 1.01(–6)|
|4725.62||4.52(–2) 2.58(–3)||3.43(–6) 9.78(–7)|
|Ne||3345.83||3.22(–1) 1.98(–2)||14920||3960||1.99(–7) 3.40(–8)|
|3425.87||8.71(–1) 8.29(–3)||1.97(–7) 3.15(–8)|
|S||4068.60||3.95(–1) 8.37(–3)||12000||5740||4.05(–8) 1.96(–9)|
|4076.35||2.45(–2) 9.13(–3)||7.44(–9) 2.79(–9)|
|6716.44||1.23(–1) 4.55(–3)||1.03(–8) 5.12(–10)|
|6730.81||2.16(–1) 4.88(–3)||1.03(–8) 4.01(–10)|
|S||6312.10||4.77(–2) 4.75(–3)||12460||3590||6.87(–8) 1.07(–8)|
|9068.60||3.78(–1) 9.97(–3)||7.34(–8) 4.79(–9)|
|18.7||6.92(–1) 4.67(–2)||6.81(–8) 4.82(–9)|
|S||10.5||1.92(0) 5.03(–2)||13650||3960||5.28(–8) 1.51(–9)|
|Cl||5517.66||1.81(–2) 2.82(–3)||12460||3590||1.36(–9) 2.42(–10)|
|Cl||8046.30||2.05(–2) 2.60(–3)||13650||3960||7.82(–10) 1.04(–10)|
|Ar||5191.82||3.90(–3) 1.61(–3)||13650||3960||1.23(–8) 5.17(–9)|
|7135.80||2.73(–1) 1.14(–2)||1.30(–8) 7.41(–10)|
|7751.10||6.11(–2) 2.51(–3)||1.22(–8) 6.83(–10)|
|Ar||4711.37||9.40(–2) 5.23(–3)||13650||3960||7.53(–9) 5.46(–10)|
|4740.17||8.99(–2) 3.71(–3)||7.58(–9) 4.60(–10)|
|7170.50||6.53(–3) 7.75(–4)||4.32(–8) 6.13(–9)|
|7262.70||4.58(–3) 3.99(–3)||3.52(–8) 3.08(–8)|
|Fe||4881.00||2.14(–2) 4.84(–3)||12000||1510||1.16(–8) 2.80(–9)|
|5270.40||2.19(–2) 3.69(–3)||1.04(–8) 1.79(–9)|
|Fe||6740.63||1.58(–2) 4.81(–3)||13650||3960||1.02(–7) 3.26(–8)|
|Kr||5346.02||4.56(–3) 2.02(–3)||13650||3960||1.41(–10) 6.31(–11)|
|Ba||4934.08||6.42(–3) 1.80(–3)||9550||1030||1.90(–10) 6.45(–11)|
|6141.70||3.43(–3) 7.66(–4)||2.12(–10) 6.42(–11)|
3.6. Ionic abundances from ORLs
We have detected many optical recombination lines (ORLs) of helium, carbon, nitrogen, oxygen, and neon. To our knowledge, the nitrogen, oxygen, and neon ORLs are detected for the first time from this PN. These lines provide us with a new independent method to derive chemical abundances for BoBn 1. The recombination coefficient depends weakly on the electron temperature ( ). The ionic abundances are, therefore, insignificantly affected by small-scale fluctuations of electron temperature. This is the most important advantage of this determination method. The ionic abundances from recombination lines are robust against uncertainty in electron temperature estimation.
The ORL ionic abundances X/H are derived from
where is the recombination coefficient for the ion . For calculating ORL ionic abundances, we adopted of 8800 K and of 10 cm from the hydrogen recombination spectrum.
Effective recombination coefficients for the lines’ parent multiplets were taken from the references listed in Table 11. The recombination coefficients for each multiplet at a given electron density were calculated by fitting the polynomial functions of . The recombination coefficient of each line was obtained by a branching ratio, , which is the ratio of the recombination coefficient of the target line, to the total recombination coefficient, in a multiplet line. To calculate the branching ratio, we referred to Wiese et al. (1996) except for O ii 3p-3d and 3d-4f transitions and Ne ii. For O ii 3p-3d and 3d-4f transition lines, the branching ratios were provided by Liu et al. (1995) based on intermediate coupling. For Ne ii, Kisielius et al. (1998) provided the branching ratios based on -coupling.
The estimated ORL ionic abundances are listed in Tables 12 and 13. In general, a Case B assumption applies to the lines from levels having the same spin as the ground state, and a Case A assumption applies to lines of other multiplicities. In the last one of the line series of each ion, we present the adopted ionic abundance and the error, which are estimated from the line intensity-weighted mean.
|He ii||3-4, 4-6||(4)|
|C ii||3d-4f, 3s-3p, 3p-3d||(5)|
|4f-7g, 4d-6f, 4f-6g|
|C iv||2p-2s, 5fg-6gh||(4)|
|N ii||3s-3p, 3p-3d||(7)|
|N iii||3s-3p, 3p-3d, 4f-5g||(4)|
|Ne ii||3s-3p, 3p-3d, 3s-3p||(11)|
References. – (1) Aller (1984). (2) Storey & Hummer (1995). (3) Benjamin et al. (1999). (4) Péquignot et al. (1991). (5) Davey et al. (2000). (6) Nussbaumer & Storey (1984). (7) Kisielius & Storey (2002). (8) Escalate & Victor (1990). (9) Storey (1994). (10) Liu et al. (1995). (11) Kisielius et al. (1998).
The He abundances are estimated using electron density insensitive five He i lines to reduce intensity enhancement by collisional excitation from the He 2 level. The collisional excitation from the He 2 level enhances mainly the intensity of the triplet He i lines. We removed this contributions (1.4 for He i 4387; up to 7.4 for He i 5876) from the observed line intensities using the formulae given by Kingdon & Ferland (1995).
The He abundance is estimated from He ii 4686. Kniazev et al. (2008) estimated He = 8.52(–2) and He = 1.53(–2), which are close to or slightly smaller than our values.
We observed C ii lines which arose from different transitions. The ground state of C ii line is a doublet (2 ). The 3d-4f (multiplet V6), 4d-6f (V16.04), 4f-6g (V17.04), and 4f-7g (V17.06) lines, which have higher angular momentum as upper levels, are unaffected by both resonance fluorescence by starlight and recombination from excited and terms. Among these high angular momentum lines, the V6 lines are the most case-insensitive and reliable. Comparison of the C abundance derived from C ii 4267 with that of the other C ii lines indicates that the observed C ii lines are not populated by the intensity enhancement mechanisms discussed above. Therefore we can safely use all the C ii lines for the estimation of C abundance.
All the observed C iii lines are triplets. Since the ground state of C iii is singlet (2 ), we adopted Case A assumption. Unlike the case of C ii, C iii lines are relatively case insensitive. Our estimated C and C abundances (Table 12) are in good agreement with Kanizev et al. (2008); their C and C are 7.78(–4) and 5.62(–4), respectively.
We estimate the C abundance using multiplet V8 and V8.01 lines. Interestingly, we observed C iv 5811. C iv 5801/11 has been detected in PNe with Wolf Rayet type central stars, suggesting that the central star is very active. In the case of BoBn 1 C iv lines might be be nebular origin rather than the central star origin, because the 2 of C iv 5811 is 14.3 km s comparable with the value in close I.P. ions such as [Ne iv] and [F iv] (see Table 6).
All of the observed N ii lines are triplets. Since the ground level of N ii is a triplet (2 ), we adopted Case B assumption. The N ii resonance line 2 24 508.668 Å can be enhanced by the He i resonance line 1 s 18 508.697 Å. The cascade transition from 24 can enhance the line intensity of the multiplet V3 lines. But, this transition cannot enhance the line intensity of 3f-4d transition (multiplet V43b, V48a, V50a, and V55a) due to the lack of a direct resonance or cascade excitation path. Comparison of N abundances derived from the 3f-4d with those from the V3 lines implies that the fluorescence is negligible in BoBn 1.
The multiplet V1, V2, and V17 N iii lines are observed. We adopted Case B assumption except for the V17 multiplet. For the V17 line, we adopted Case A assumption. The intensity of the resonance N iii line 374.36 Å (2 3 ) may be enhanced by O iii resonance at 374.11 Å (2 3 ). The line intensity of the multiplet V1 and V2 lines might be enhanced by the O iii lines. The multiplet V17 line (4f-5g transition) does not appear to be enhanced. Therefore, we adopt the N abundance from this line.
We observed O ii doublet (3d-4f) and quadruplet lines (multiplet V1, V4, V10, V19). Most of the V 1 lines and all of the V 2 lines are observed. Since the ground level of O ii is a quadruplet, we adopted Case A for the doublet lines and Case B for quadruplet lines. It seems that the multiplet V1 and V10 lines give the most reliable value.
A number of O iii lines are observed. We consider Case B for the triplet lines (multiplet V2) and Case A for the singlet line (multiplet V5). There is a possibility that the multiplet V2 lines would be excited by the Bowen fluorescence mechanism or by the charge exchange of O and H instead of recombination and the multiplet V5 line could be excited by charge exchange. Therefore, we did not use O abundances in the estimation of a total oxygen abundance from ORLs.
The observed Ne ii lines are doublet (multiplets V9 and V21) and quartet lines (V1 and V2). We considered Case B for the doublet lines and Case A for the quartet lines. The multiplet V1 and V2 lines are insensitive to the case assumption and are pure recombination lines (Grandi 1976). Therefore, we adopted the Ne abundance derived from multiplet V 1 and V2 lines.
|V11||5876.62||18.1 0.14||9.93(–2) 1.26(–3)|
|V14||4471.47||4.82 0.04||9.63(–2) 2.89(–3)|
|V46||6678.15||4.02 0.11||9.59(–2) 3.80(–3)|
|V48||4921.93||1.32 0.02||9.66(–2) 3.02(–3)|
|V51||4387.93||0.59 0.01||9.54(–2) 3.57(–3)|
|3.4||4685.68||24.8 0.79||2.03(–2) 6.47(–4)|
|V2||6578.05||3.75(–1) 6.87(–3)||7.30(–4) 3.50(–5)|
|V6||4267.15||7.90(–1) 4.42(–2)||7.55(–4) 5.02(–5)|