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Shane B. Vickers, David J. Frew, Quentin A. Parker, Ivan S. Bojičić, New light on Galactic post-asymptotic giant branch stars – I. First distance catalogue, Monthly Notices of the Royal Astronomical Society, Volume 447, Issue 2, 21 February 2015, Pages 1673–1691, https://doi.org/10.1093/mnras/stu2383
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Abstract
We have commenced a detailed analysis of the known sample of Galactic post-asymptotic giant branch (PAGB) objects compiled in the Toruń catalogue of Szczerba et al., and present, for the first time, homogeneously derived distance determinations for the 209 likely and 87 possible catalogued PAGB stars from that compilation. Knowing distances are essential in determining meaningful physical characteristics for these sources and this has been difficult to determine for most objects previously. The distances were determined by modelling their spectral energy distributions (SEDs) with multiple blackbody curves, and integrating under the overall fit to determine the total distance-dependent flux. This approach was undertaken for consistency as precise spectral types, needed for more detailed fitting, were unknown in the majority of cases. The SED method works because the luminosity of these central stars is very nearly constant from the tip of the AGB phase to the beginning of the white dwarf cooling track. This then enables us to use a standard-candle luminosity to estimate the SED distances. For Galactic thin-disc PAGB objects, we use three luminosity bins based on typical observational characteristics, ranging between 3500 and 12 000 L⊙. We further adopt a default luminosity of 4000 L⊙ for bulge objects and 1700 L⊙ for the thick-disc and halo objects. We have also applied the above technique to a further sample of 54 related nebulae not in the current edition of the Toruń catalogue. In a follow-up paper, we will estimate distances to the subset of RV Tauri variables using empirical period–luminosity relations, and to the R CrB stars, allowing a population comparison of these objects with the other subclasses of PAGB stars for the first time.
1 INTRODUCTION
Pre-planetary nebulae (PPNe) are a very brief phase in the late-stage evolution of mid-mass stars (∼1–8 M⊙) between the asymptotic giant branch (AGB) and the planetary nebula (PN) phases (Kwok, Purton & Fitzgerald 1978; Kwok 1982; Balick & Frank 2002). The ejection of the tenuous envelope in the final superwind stage of AGB evolution (Renzini 1981) reaches rates of up to 10−4 M⊙ yr−1, and leads to an increase in effective temperature of the central star. This rate of temperature increase is a strong function of the core mass (Schönberner 1983; Vassiliadis & Wood 1994, hereafter VW94) and ultimately determines if the core reaches a temperature high enough to photoionize the ejected matter as a PN, before it disperses into the surrounding interstellar medium.
The relative scarcity of known Galactic post-asymptotic giant branch (PAGB) objects (∼450; Szczerba et al. 2007, 2012)1 stems from the brevity of the PAGB evolutionary stage (decades to a few thousand years), which for high core masses can be so brief that we are unlikely to observe these rapidly evolving objects (VW94; Blöcker 1995). The evolution of PPNe is typically characterized by a near-constant bolometric luminosity, and a double-peaked spectral energy distribution (SED), manifest as a large infrared (IR) excess.
Understanding these objects is dependent on accurate distances, which are not available for most of the more poorly quantified objects. Yet this phase is key to comprehending the shaping mechanisms of PNe (Balick & Frank 2002), as the dust shells around AGB stars, the precursors of PNe, have morphologies that are typically spherically symmetric (Corradi et al. 2003; Mauron & Huggins 2006; Cox et al. 2012; Mauron, Huggins & Cheung 2013), while imaging surveys of PNe show round morphologies to be in the minority (Balick 1987; Manchado et al. 1996; Górny et al. 1999; Parker et al. 2006; but see Jacoby et al. 2010). In order to better understand this conundrum, several imaging studies of PPNe have been undertaken over the last two decades, both at optical and IR wavelengths (Sahai & Trauger 1998; Su et al. 1998; Hrivnak et al. 1999; Meixner et al. 1999; Sahai et al. 1999; Kwok et al. 2000; Ueta, Meixner & Bobrowsky 2000; Hrivnak, Kwok & Su 2001; Gledhill 2005; Siódmiak et al. 2008; Lagadec et al. 2011a), with the goal of better understanding the shaping mechanisms of PPNe and PNe.
The central stars of dusty PPNe are invariably obscured making their identification difficult, and in contrast the dust shell is subordinate to the central star in many high-latitude objects. This variety of properties has led to a myriad of different identification schemes and criteria in their identification to date. These include the presence of an IR excess (Zuckerman 1978; Parthasarathy & Pottasch 1986; Pottasch & Parthasarathy 1988; Hrivnak, Kwok & Volk 1989), including the utilization of Infrared Astronomical Satellite (IRAS) colour–colour diagrams (e.g. Preite-Martinez 1988; van der Veen & Habing 1988; Manchado et al. 1989; Van der Veen, Habing & Geballe 1989; Oudmaijer et al. 1992; Hu et al. 1993; García-Lario et al. 1997; Sahai et al. 2007). Identifying PAGB stars on the basis of their mid-IR (MIR) spectra is also undertaken, including the presence of the distinctive 21 μm feature in carbon-rich PPNe (Kwok, Volk & Hrivnak 1989; Cerrigone et al. 2011). More recently, surveys for new PAGB stars using near-IR (NIR) photometric data (Ramos-Larios et al. 2009, 2012), and the related R CrB stars using either NIR and MIR colours (Tisserand et al. 2011; Tisserand 2012), or ASAS-3 optical light curves (Tisserand et al. 2013), have been undertaken.
An improved understanding of this evolutionary phase is dependent on determining accurate distances to a large sample of objects, which can be used to furnish meaningful physical characteristics. Unfortunately at present, reliable distances have so far been determined for only a small fraction of these objects. It is therefore imperative that an accurate method for calculating distances to PAGB objects is determined. Since the PAGB phase is characterized by a near constant bolometric luminosity for their central stars, from the AGB-tip to the beginning of the white dwarf (WD) cooling track (Paczyński 1971; Schönberner 1983; VW94; Blöcker 1995), we can use a standard-candle luminosity to estimate the distances to them. This is the main focus and legacy of this paper, described in detail in Section 3, below.
1.1 Nomenclature
For the benefit of the reader, we briefly describe the nomenclature that we have adopted in this paper. The generic term, PAGB star, includes all objects evolving from the AGB to the beginning of the WD cooling track, with or without a surrounding nebula, though in practice, PNe are defined separately (Kwok 1993, 2010; Frew & Parker 2010). With a few exceptions (e.g. Jacoby et al. 1997; Alves, Bond & Livio 2000), most Population II PAGB stars do not have extensive surrounding dust shells nor ionized PNe. In this paper, we use the term PPNe to describe the non-ionized, dusty nebulae that scatter the light of their embedded central stars, and which emit in the thermal-IR (Pottasch & Parthasarathy 1988; van der Veen & Habing 1988). Spectra of their central stars range from B-type at the hot end to late-K or even early M-type at the cool end (Volk & Kwok 1989; van Hoof, Oudmaijer & Waters 1997; Van Winckel 2003). Once the effective temperature of the central star reaches about 20 kK, the surrounding material is photoionized to produce a young PN. The term ‘transition object’ is sometimes used to describe objects just commencing this process of ionization (Suárez et al. 2006; Cerrigone et al. 2008; Frew, Bojičić & Parker 2013), graphically demonstrated by the recent evolution of CRL 618 (Tafoya et al. 2013).
Ueta et al. (2000) classified resolved PPNe into two groups: the star-obvious low-level elongated (SOLE) nebulae, which have a visible nebula around an obvious central star, and the dust-prominent longitudinally extended (DUPLEX) objects, which typically have faint or even invisible central stars obscured by a dusty torus. These differences extend to other observed properties. In general, the SOLE objects have bluer IR colours than the dustier DUPLEX nebulae (Ueta et al. 2000; Siódmiak et al. 2008). Similarly the SEDs are different: the SEDs of the SOLE objects are typically two peaked, dominated by the stellar photosphere and the dust component, while the DUPLEX sources have a very prominent dust peak, with little or no optical peak (Siódmiak et al. 2008).
Siódmiak et al. (2008) also noted the different distributions in Galactic latitude of the two classes, and along with Meixner et al. (2002), suggested that the DUPLEX nebulae derive from more massive progenitor stars and are the natural precursors to bipolar PNe, which have been shown to have a lower scaleheight than other PNe (Corradi & Schwarz 1995; Phillips 2001; Frew 2008). We will investigate this problem in more detail in a later paper in this series. It has become apparent that many PAGB stars have little or no dust around them, and these are generally thought to be of low mass (e.g. Alcolea & Bujarrabal 1991; Bujarrabal et al. 2013). Note that the ‘likely’ PAGB section of the Toruń catalogue includes the UU Her variables (e.g. Sasselov 1984), named after the prototype UU Herculis, whose own classification is debatable (Klochkova, Panchuk & Chentsov 1997). We consider these stars to be the lower luminosity halo analogues of the old-disc ‘89 Herculis’ stars (Gillett, Hyland & Stein 1970; Bujarrabal et al. 2007, and references therein).
Several other groups of uncommon stars are often classed as possible PAGB objects. These include the RV Tauri stars (Preston et al. 1963; Goldsmith et al. 1987; Van Winckel et al. 1999), pulsating yellow supergiants related to the type II Cepheids (Wallerstein 2002). The more luminous RV Tau stars are usually considered to be PAGB stars with low initial masses (Jura 1986; cf. Matsuura et al. 2002). For these stars, distances can be estimated using the period–luminosity (P–L) relation for Population II Cepheids (Alcock et al. 1998; Matsunaga et al. 2006; Soszyński et al. 2008; Matsunaga, Feast & Menzies 2009).2 A detailed study of their distances and space distribution will be the subject of the second paper in this series (Vickers et al., in preparation). The hydrogen-deficient R Coronae Borealis stars (Clayton 1996, 2012), and their hotter kin, the extreme helium stars (e.g. Pandey et al. 2001; Jeffery 2008), will also be evaluated in that work.
1.2 Scientific motivation
The compilation of the Toruń catalogue of Galactic PAGB and related objects (Szczerba et al. 2007, 2012) now provides a central repository of information for all currently identified Galactic PAGB stars, facilitating a wider study of these objects. To date, the general physical characteristics of PAGB objects have only been determined from relatively small samples, using well-studied PAGB objects with ample data available. Prior to the Toruń catalogue, it was necessary to collect scattered photometric and spectroscopic data to find candidate PAGB stars, i.e. a source displaying canonical PAGB colours (Pottasch & Parthasarathy 1988; van der Veen & Habing 1988). This effectively meant that that a large-scale investigation of the Galactic PAGB population was not feasible.
While some of the data sets available in the Toruń catalogue are limited in quality, more recent all-sky surveys such as the AKARI (Astro-F) survey (Ishihara et al. 2010) and Wide-field Infrared Survey Explorer (WISE; Wright et al. 2010) survey provide photometric data that is more sensitive and of higher resolution than those previously utilized.
Here we present, for the first time, a homogenized catalogue of distances of all known Galactic PAGB objects available in the Toruń catalogue. Distances have been calculated using the observed SEDs, generated using the photometric and spectroscopic data gathered in the Toruń catalogue, as well as additional photometric data from recent all-sky surveys. Due to the narrow distribution of WD masses (Vennes et al. 2002; Kleinman et al. 2013; and others), we have adopted assumed luminosities for specific subtypes rather than attempting to determine an individual luminosity for each object (see Section 3.1, below).
Our distance catalogue will allow new insights into this brief, poorly understood phase of late-stage stellar evolution by allowing a population study based on improved distance estimates. The paper will proceed as follows: in Section 2, we outline the material used including additional data sources taken from the literature, and in Section 3, we detail the method used for the SED derived distances. In Section 4, we provide the reader with a sample table of the SED calculated distances, and a comparison of the results with independent literature distances. In Section 6, we summarize our findings and give suggestions for future work. The resulting catalogue of distances as well as the fitted SEDs will be available in full as an online supplement.
2 THE TORUŃ CATALOGUE
The Toruń catalogue provides easy online access to processed photometric and spectroscopic data for the currently identified Galactic population of PAGB stars and related objects. With the advent of this compilation of all known such objects with associated flux data, our distance technique can be applied to them in their entirety, leading to a large-enough sample to exploit for scientific purposes. Prior to the Toruń catalogue, the Galactic PAGB population was only available in subsets of ‘candidate’ objects. With the advent of this compilation of all known objects and flux data, our distance technique can be applied to the known PAGB population, leading to a large-enough sample to exploit for scientific purposes.
The catalogue is divided into five categories: (i) very-likely PAGB stars, (ii) RV Tauri stars, (iii) R Coronae Borealis/extreme helium/late thermal pulse stars, (iv) possible PAGB stars and (v) unlikely PAGB objects. The data from the catalogue are summarized in Table 1. Hereafter, likely PAGB stars will be referred to simply as PAGB, R Coronae Borealis/extreme helium/late thermal pulse as R CrB/eHe/LTP, while the possible PAGB objects will be simply referred to as possible. We will present a distance catalogue of the R Tau and R CrB/eHe/LTP stars in a second paper (Vickers et al., in preparation), concentrating on the likely and possible PAGB objects in this work. Unlikely objects will in the main not be considered in this paper, except for some objects included in Section 5.1.
Category . | Number . |
---|---|
Likely PAGB | 209 |
Possible PAGB | 87 |
RV Tauri | 112 |
R CrB/eHe/LTP | 72 |
Unlikely PAGB | 72 |
Total | 480a |
Category . | Number . |
---|---|
Likely PAGB | 209 |
Possible PAGB | 87 |
RV Tauri | 112 |
R CrB/eHe/LTP | 72 |
Unlikely PAGB | 72 |
Total | 480a |
Note.aExcludes those objects categorized as unlikely PAGB stars.
Category . | Number . |
---|---|
Likely PAGB | 209 |
Possible PAGB | 87 |
RV Tauri | 112 |
R CrB/eHe/LTP | 72 |
Unlikely PAGB | 72 |
Total | 480a |
Category . | Number . |
---|---|
Likely PAGB | 209 |
Possible PAGB | 87 |
RV Tauri | 112 |
R CrB/eHe/LTP | 72 |
Unlikely PAGB | 72 |
Total | 480a |
Note.aExcludes those objects categorized as unlikely PAGB stars.
The Toruń catalogue also includes optical fluxes from the Tycho-2 and Guide Star Catalogues (GSC; Høg et al. 2000; Lasker et al. 2008), along with Deep Near Infrared Survey of the Southern Sky (DENIS) IJKs (Epchtein et al. 1999; Schmeja & Kimeswenger 2001) and Two Micron All Sky Survey (2MASS) JHKs photometry (Cutri 2003; Skrutskie et al. 2006). This is supplemented with MIR photometric data from the IRAS (Neugebauer et al. 1984) and the Midcourse Space Experiment 6C catalogues (MSX6C; Price et al. 2001).
For the 2MASS data, we have excluded magnitudes with problematic quality flags of X, U, F and E, which leaves valid data with flags A, B, C and D, where SNR ≥ 5 for A, B and C flags (Cutri 2003). For the MSX6C data, we have excluded data with quality flag 1 which removes all data with SNR ≤ 5 (Egan et al. 2003), while for the IRAS fluxes, we have removed all upper limits (FQUAL = 1) across all four wavebands (Beichman et al. 1988).
2.1 Supplemental data
Here, we describe a number of additional data sources which we have used to supplement the data presented in the Toruń catalogue. The additional data include data from several minor surveys, plus data that have been published since the most recent release of the Toruń catalogue (v 2.0; Szczerba et al. 2012). To gather much of these data, we interrogated the fifth edition of the Catalogue of Infrared Observations (Gezari et al.1993, and references therein). This is a valuable source of literature data, but the catalogue includes both line and continuum fluxes, and data obtained using different aperture diameters, so in order to remove problematic fluxes we needed to carefully vet the data, object by object. We also utilized more recent MIR flux data from the literature for individual sources if available (e.g. Smith & Gehrz 2005; Hora et al. 2008; Lagadec et al. 2011a). Table 2 gives a comparison of the wavelengths and the angular resolution of the major surveys and catalogues that we have utilized. For consistency, we have taken the zero-magnitude fluxes from the SVO filter profile service.3
Survey/catalogue . | Wavebands (λeff μm) . | Resolution . | Reference . |
---|---|---|---|
GALEX | FUV (0.15), NUV (0.23) | ∼4–6 arcsec | Morrissey et al. (2007) |
TD-1 | 0.157, 0.197, 0.237, 0.274 | ∼7 arcmin | Thompson et al. (1978) |
ANS | 0.155, 0.180, 0.220, 0.250, 0.330 | 2.5 arcmin | Wesselius et al. (1982) |
Tycho-2 | BT (0.44), VT (0.51) | ∼0.8 arcsec | Høg et al. (2000) |
APASS | B (0.44), g′ (0.47), V (0.54), r′ (0.62), i′ (0.75) | ∼10 arcsec | Henden et al. (2012) |
DENIS | I (0.82), J (1.25), Ks (2.15) | 1–3 arcsec | Epchtein et al. (1997) |
UKIDSS | Z (0.88), Y (1.03), J (1.25), H (1.66), Ks (2.15) | 1 arcsec | Lawrence et al. (2007) |
2MASS | J (1.24), H (1.66), Ks (2.16) | 2 arcsec | Skrutskie et al. (2006) |
WISE | W1 (3.4), W2 (4.6), W3 (12), W4 (22) | 6–12 arcsec | Wright et al. (2010) |
COBE/DIRBEa | 3.5, 4.9, 12, 25, 60 | 40 arcmin | Smith, Price & Baker (2004) |
Spitzer (IRAC) | IRAC1 (3.6), IRAC2 (4.5), IRAC3 (5.8), IRAC4 (8.0) | ≤2 arcsec | Fazio et al. (2004) |
RAFGL | 4.2, 11.0, 19.8, 27.4 | 3.5 arcmin | Price & Murdock (1983) |
MSX6C | A (8.3), C (12.1), D (14.7), E (21.3) | 18 arcsec | Price et al. (2001) |
AKARI (IRC) | S9W (9.0), L18W (18.0) | ∼2 arcsec | Ishihara et al. (2010) |
AKARI (FIS) | 65, 90, 140, 160 | 30–50 arcsec | Ishihara et al. (2010) |
IRAS | 12, 25, 60, 100 | 0.5–2 arcmin | Neugebauer et al. (1984) |
Spitzer (MIPS) | 24, 70, 160 | 6–40 arcsec | Rieke et al. (2004); Carey et al. (2009) |
Herschel (PACS) | Blue (70), red (160) | 5–35 arcsec | Pilbratt et al. (2010) |
Herschel (SPIRE) | PSW (250), PMW (350), PLW (500) | 5–35 arcsec | Pilbratt et al. (2010) |
SCUBA | 450, 850 | 8–14 arcsec | Holland et al. (1999) |
Planckb | 857 GHz (350), 545 GHz (550), 353 GHz (849) | 5–30 arcsec | Planck Collaboration VII (2011) |
Survey/catalogue . | Wavebands (λeff μm) . | Resolution . | Reference . |
---|---|---|---|
GALEX | FUV (0.15), NUV (0.23) | ∼4–6 arcsec | Morrissey et al. (2007) |
TD-1 | 0.157, 0.197, 0.237, 0.274 | ∼7 arcmin | Thompson et al. (1978) |
ANS | 0.155, 0.180, 0.220, 0.250, 0.330 | 2.5 arcmin | Wesselius et al. (1982) |
Tycho-2 | BT (0.44), VT (0.51) | ∼0.8 arcsec | Høg et al. (2000) |
APASS | B (0.44), g′ (0.47), V (0.54), r′ (0.62), i′ (0.75) | ∼10 arcsec | Henden et al. (2012) |
DENIS | I (0.82), J (1.25), Ks (2.15) | 1–3 arcsec | Epchtein et al. (1997) |
UKIDSS | Z (0.88), Y (1.03), J (1.25), H (1.66), Ks (2.15) | 1 arcsec | Lawrence et al. (2007) |
2MASS | J (1.24), H (1.66), Ks (2.16) | 2 arcsec | Skrutskie et al. (2006) |
WISE | W1 (3.4), W2 (4.6), W3 (12), W4 (22) | 6–12 arcsec | Wright et al. (2010) |
COBE/DIRBEa | 3.5, 4.9, 12, 25, 60 | 40 arcmin | Smith, Price & Baker (2004) |
Spitzer (IRAC) | IRAC1 (3.6), IRAC2 (4.5), IRAC3 (5.8), IRAC4 (8.0) | ≤2 arcsec | Fazio et al. (2004) |
RAFGL | 4.2, 11.0, 19.8, 27.4 | 3.5 arcmin | Price & Murdock (1983) |
MSX6C | A (8.3), C (12.1), D (14.7), E (21.3) | 18 arcsec | Price et al. (2001) |
AKARI (IRC) | S9W (9.0), L18W (18.0) | ∼2 arcsec | Ishihara et al. (2010) |
AKARI (FIS) | 65, 90, 140, 160 | 30–50 arcsec | Ishihara et al. (2010) |
IRAS | 12, 25, 60, 100 | 0.5–2 arcmin | Neugebauer et al. (1984) |
Spitzer (MIPS) | 24, 70, 160 | 6–40 arcsec | Rieke et al. (2004); Carey et al. (2009) |
Herschel (PACS) | Blue (70), red (160) | 5–35 arcsec | Pilbratt et al. (2010) |
Herschel (SPIRE) | PSW (250), PMW (350), PLW (500) | 5–35 arcsec | Pilbratt et al. (2010) |
SCUBA | 450, 850 | 8–14 arcsec | Holland et al. (1999) |
Planckb | 857 GHz (350), 545 GHz (550), 353 GHz (849) | 5–30 arcsec | Planck Collaboration VII (2011) |
Notes.aOther wavelengths have been excluded.
bWe have excluded data with a wavelength longer than 1 mm (217, 143, 100 GHz).
Survey/catalogue . | Wavebands (λeff μm) . | Resolution . | Reference . |
---|---|---|---|
GALEX | FUV (0.15), NUV (0.23) | ∼4–6 arcsec | Morrissey et al. (2007) |
TD-1 | 0.157, 0.197, 0.237, 0.274 | ∼7 arcmin | Thompson et al. (1978) |
ANS | 0.155, 0.180, 0.220, 0.250, 0.330 | 2.5 arcmin | Wesselius et al. (1982) |
Tycho-2 | BT (0.44), VT (0.51) | ∼0.8 arcsec | Høg et al. (2000) |
APASS | B (0.44), g′ (0.47), V (0.54), r′ (0.62), i′ (0.75) | ∼10 arcsec | Henden et al. (2012) |
DENIS | I (0.82), J (1.25), Ks (2.15) | 1–3 arcsec | Epchtein et al. (1997) |
UKIDSS | Z (0.88), Y (1.03), J (1.25), H (1.66), Ks (2.15) | 1 arcsec | Lawrence et al. (2007) |
2MASS | J (1.24), H (1.66), Ks (2.16) | 2 arcsec | Skrutskie et al. (2006) |
WISE | W1 (3.4), W2 (4.6), W3 (12), W4 (22) | 6–12 arcsec | Wright et al. (2010) |
COBE/DIRBEa | 3.5, 4.9, 12, 25, 60 | 40 arcmin | Smith, Price & Baker (2004) |
Spitzer (IRAC) | IRAC1 (3.6), IRAC2 (4.5), IRAC3 (5.8), IRAC4 (8.0) | ≤2 arcsec | Fazio et al. (2004) |
RAFGL | 4.2, 11.0, 19.8, 27.4 | 3.5 arcmin | Price & Murdock (1983) |
MSX6C | A (8.3), C (12.1), D (14.7), E (21.3) | 18 arcsec | Price et al. (2001) |
AKARI (IRC) | S9W (9.0), L18W (18.0) | ∼2 arcsec | Ishihara et al. (2010) |
AKARI (FIS) | 65, 90, 140, 160 | 30–50 arcsec | Ishihara et al. (2010) |
IRAS | 12, 25, 60, 100 | 0.5–2 arcmin | Neugebauer et al. (1984) |
Spitzer (MIPS) | 24, 70, 160 | 6–40 arcsec | Rieke et al. (2004); Carey et al. (2009) |
Herschel (PACS) | Blue (70), red (160) | 5–35 arcsec | Pilbratt et al. (2010) |
Herschel (SPIRE) | PSW (250), PMW (350), PLW (500) | 5–35 arcsec | Pilbratt et al. (2010) |
SCUBA | 450, 850 | 8–14 arcsec | Holland et al. (1999) |
Planckb | 857 GHz (350), 545 GHz (550), 353 GHz (849) | 5–30 arcsec | Planck Collaboration VII (2011) |
Survey/catalogue . | Wavebands (λeff μm) . | Resolution . | Reference . |
---|---|---|---|
GALEX | FUV (0.15), NUV (0.23) | ∼4–6 arcsec | Morrissey et al. (2007) |
TD-1 | 0.157, 0.197, 0.237, 0.274 | ∼7 arcmin | Thompson et al. (1978) |
ANS | 0.155, 0.180, 0.220, 0.250, 0.330 | 2.5 arcmin | Wesselius et al. (1982) |
Tycho-2 | BT (0.44), VT (0.51) | ∼0.8 arcsec | Høg et al. (2000) |
APASS | B (0.44), g′ (0.47), V (0.54), r′ (0.62), i′ (0.75) | ∼10 arcsec | Henden et al. (2012) |
DENIS | I (0.82), J (1.25), Ks (2.15) | 1–3 arcsec | Epchtein et al. (1997) |
UKIDSS | Z (0.88), Y (1.03), J (1.25), H (1.66), Ks (2.15) | 1 arcsec | Lawrence et al. (2007) |
2MASS | J (1.24), H (1.66), Ks (2.16) | 2 arcsec | Skrutskie et al. (2006) |
WISE | W1 (3.4), W2 (4.6), W3 (12), W4 (22) | 6–12 arcsec | Wright et al. (2010) |
COBE/DIRBEa | 3.5, 4.9, 12, 25, 60 | 40 arcmin | Smith, Price & Baker (2004) |
Spitzer (IRAC) | IRAC1 (3.6), IRAC2 (4.5), IRAC3 (5.8), IRAC4 (8.0) | ≤2 arcsec | Fazio et al. (2004) |
RAFGL | 4.2, 11.0, 19.8, 27.4 | 3.5 arcmin | Price & Murdock (1983) |
MSX6C | A (8.3), C (12.1), D (14.7), E (21.3) | 18 arcsec | Price et al. (2001) |
AKARI (IRC) | S9W (9.0), L18W (18.0) | ∼2 arcsec | Ishihara et al. (2010) |
AKARI (FIS) | 65, 90, 140, 160 | 30–50 arcsec | Ishihara et al. (2010) |
IRAS | 12, 25, 60, 100 | 0.5–2 arcmin | Neugebauer et al. (1984) |
Spitzer (MIPS) | 24, 70, 160 | 6–40 arcsec | Rieke et al. (2004); Carey et al. (2009) |
Herschel (PACS) | Blue (70), red (160) | 5–35 arcsec | Pilbratt et al. (2010) |
Herschel (SPIRE) | PSW (250), PMW (350), PLW (500) | 5–35 arcsec | Pilbratt et al. (2010) |
SCUBA | 450, 850 | 8–14 arcsec | Holland et al. (1999) |
Planckb | 857 GHz (350), 545 GHz (550), 353 GHz (849) | 5–30 arcsec | Planck Collaboration VII (2011) |
Notes.aOther wavelengths have been excluded.
bWe have excluded data with a wavelength longer than 1 mm (217, 143, 100 GHz).
2.1.1 Optical and ultraviolet photometry
From a perusal of the GSC magnitudes presented in the Torun catalogue, it is clear that there is significant confusion between stellar and nebular fluxes, such that the B-, V- and R-band data differ from independent data in some cases by more than an order of magnitude. Therefore, we have removed these data from the fitting process, substituting with UBVRI photometry extracted from the compilations of Mermilliod, Mermilliod & Hauck (1997) and Mermilliod (2006), supplemented with other data retrieved through the VizieR service4 if available. The recently available AAVSO Photometric All-Sky Survey (APASS; Henden et al. 2012) was particularly useful for stars in the 10th–17th visual magnitude range. A lack of optical data is to be expected for dust enshrouded PAGB objects where the central star is almost entirely obscured. Because of this, APASS photometry was only available for about 40 objects.
We have collected ubvy Strömgren photometry from the ubvy − β catalogue of Hauck & Mermilliod (1998). Here, we have assumed that the Strömgren y band is equivalent to the Johnson V, and we have removed the u-band photometry from the fitting procedure as it lies entirely below the Balmer discontinuity and does not lend itself to blackbody fitting (but would be useful for fitting model atmospheres).
We also utilized ultraviolet (UV) fluxes, which are needed to constrain the SEDs of the hotter PAGB stars. Our primary source of photometry is from the Galaxy Evolution Explorer (GALEX) mission (Morrissey et al. 2007), which covers most of the sky (excluding the Galactic plane) at a resolution of 4–6 arcsec. We have taken the GALEX magnitudes directly from the survey website,5 or from the compilations of Bianchi et al. (2011a,b), which adopted 5σ depths of mAB = 19.9 for the FUV and mAB = 20.8 for the NUV. Since saturation in both the FUV and NUV detectors begins around mAB = 15, we checked all GALEX data brighter than 15th magnitude on the AB scale (Morrissey et al. 2007). Data from the older TD-1 and Netherlands Astronomical Satellite (ANS) UV space missions (Thompson et al. 1978; Wesselius et al. 1982) were used to supplement the GALEX photometry. For the TD-1 data, we have only used data with an SNR ≥ 10, necessarily restricting its use to fairly bright stars. For a single scan, the limit of the system is about 9th visual magnitude for a B-type star (Boksenberg et al. 1973). For the ANS data, only those objects are in the catalogue that have in at least one channel an SNR > 4, or in at least three channels an SNR > 3, were utilized.
2.1.2 NIR and MIR photometry
From the Toruń catalogue, we have used the given DENIS and 2MASS NIR data and supplemented this with recent JHK observations of heavily obscured objects from Ramos-Larios et al. (2012). These newly acquired data are up to 7 mag deeper than DENIS (Ks = 13.5; Epchtein et al. 1997) and 2MASS (Ks = 14.3; Skrutskie et al. 2006) going to a depth of Ks ∼ 20.4, making these data useful for fainter Galactic PAGB stars. In addition, we also interrogated the UKIRT Infrared Deep Sky Survey (UKIDSS; Lawrence et al. 2007) to obtain ZYJHKs magnitudes for fainter objects. UKIDSS is much more sensitive than 2MASS, with a Ks depth of 18.4–21.0 mag. We used the latest public data releases (DR) from the various UKIDSS surveys; the Large Area Survey (DR9), Galactic Clusters Survey (DR9), Deep Extragalactic Survey (DR9) and the Galactic Plane Survey (DR6). In several cases, we noted that incorrect NIR data have been incorporated into the Toruń catalogue, especially for objects at low latitude in crowded fields. We corrected these data accordingly, after perusal of a range of multiwavelength images taken from our new Macquarie data base (Bojičić et al., in preparation).
While the UKIDSS data supersede 2MASS for faint objects, the saturation limits (∼12 mag; Lucas et al. 2008) are significantly fainter than for 2MASS (∼5 mag for a 51 ms exposure; Cutri 2003), so the UKIDSS data are not usable for many bright PAGB stars. Similarly, we have removed DENIS I, J and Ks data brighter than the nominal saturation limits of 10, 8 and 6.5 mag, respectively, where those data do not agree with other NIR data. Where the DENIS Ks magnitudes differ significantly from the 2MASS Ks magnitudes, we uniformly adopt the 2MASS data. As the WISE 3.4, 4.6 and 11.6 μm wavebands are also prone to saturation for brighter objects, we have chosen to fit the blackbody curves through those WISE data only if there exists agreement between the WISE and other MIR data.
We also have IRAC MIR fluxes from the Spitzer Space Telescope Galactic Legacy Infrared Mid-Plane Survey Extraordinaire (Fazio et al. 2004), as reported by Hora et al. (2008), Cerrigone et al. (2009) and others, plus archival data from the Revised Air Force Geophysics Laboratory (RAFGL) survey at wavelengths of 4.2, 11.0, 20.0 and 27.0 μm (Price & Murdock 1983); see also Kleinmann, Gillett & Joyce (1981) for a review. We have also utilized five of the ten IR bands of the Cosmic Microwave Background Explorer (COBE) Diffuse Infrared Background Experiment (DIRBE) point source catalogue (Smith et al. 2004), at 3.5, 4.9, 12, 25 and 60 μm. Due to the large (42 arcmin) beam size, we have restricted our use of the DIRBE data to bright (F12 ≥ 150 Jy), high-latitude (|b| ≥ 5°) objects (e.g. Smith 2003) leaving us with data for ∼10 objects.
While IRAS (Neugebauer et al. 1984) has been the most influential instrument in the discovery and identification of PAGB objects, the survey has been superseded in the MIR, first by MSX6C (Price et al. 2001) and then by the AKARI and WISE MIR surveys. The MSX survey, while of higher resolution and sensitivity than IRAS, only covered the Galactic plane. Compared to IRAS, the AKARI satellite operated over six NIR to far-infrared (FIR) passbands. The Infrared Camera (IRC) provides photometry in two wavebands (S9W and L18W), while the far-infrared surveyor (FIS) operated at wavelengths of 65, 90, 140 and 160 μm. AKARI is up to 10 times more sensitive than the IRAS 12 and 25 μm bands (Ishihara et al. 2010) and covers 90 per cent of the sky. AKARI surveyed the sky in the MIR with greater resolution than that of IRAS.
Using VizieR, we extracted AKARI IRC and FIS data whilst excluding data from unconfirmed sources or those with a flux measurement considered to be unreliable, i.e. FQUAL < 3. In addition to the IRAS fluxes collected in the Toruń catalogue, we have queried VizieR for IRAS data (IPAC 1986) for the 54 additional objects. WISE surveyed 99 per cent of the sky with greater sensitivity than all previous MIR surveys (Wright et al. 2010). While IRAS had two FIR bands, WISE has two MIR bands (3.4 and 4.6 μm) that IRAS does not. We have excluded those WISE data which are considered upper limits (SNR ≤ 2).
2.1.3 FIR photometry
In order to constrain the FIR during the SED fitting process, we have included sub-mm fluxes where available. Several data sets, including from the Kuiper Airborne Observatory, were extracted from Gezari et al. (1993, and references therein). We also utilized 450 and 850 μm fluxes (Holland et al. 1999) obtained with the Submillimetre Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope. We have also made use of publicly released data from the Herschel Space Observatory (Pilbratt et al. 2010), using 70 and 160 μm fluxes from the Photodetector Array Camera and Spectrometer (PACS; Poglitsch et al. 2010) and 250, 350 and 500 μm fluxes from the Spectral and Photometric Imaging Receiver (SPIRE; Griffin et al. 2010), largely taken from the Mass loss for Evolving StarS (MESS) programme (Groenewegen et al. 2011).
We also used 857, 545 and 353 GHz fluxes for a few objects from the Early Release Compact Source Catalogue (Planck Collaboration VII, 2011) of the Planck6 mission (Planck Collaboration I, 2011).
2.1.4 Spectroscopy
We supplemented the photometric data with spectroscopic data where available. However, these data were not used in the fitting process as the spectra are usually not normalized; the spectral data have been retained on the plots for illustrative purposes. UV spectra from the International Ultraviolet Explorer (IUE) were downloaded from the MAST website.7 We generally excluded the IUE high-dispersion spectra (due to bad data flags) and any other spectra with a low signal-to-noise ratio. We averaged the remaining long- and short-wavelength spectra before overplotting on the SED fits.
We also included both Infrared Space Observatory (ISO) and IRAS low-resolution spectrometer (LRS) spectra (see e.g. Kwok, Volk & Bidelman 1997), when available for the objects in our sample. The LRS spectra used in the Toruń catalogue were extracted from Kevin Volk's Home Page8 just as we have done for the additional objects used in the literature comparison (see later). The LRS spectra have been absolutely calibrated according to the procedure followed by Volk & Cohen (1989) and Cohen, Walker & Witteborn (1992). For objects with available 2–45 μm ISO short wave spectrometer spectra we have extracted and overplotted these data where appropriate.
3 METHODOLOGY
Here, we outline the SED-based procedure used to calculate a set of homogenized distances for the currently known Galactic PAGB objects. In general, the SED technique for distance determinations has only been applied since the 1980s, after the IR fluxes from the IRAS mission were published (Neugebauer et al. 1984; Beichman et al. 1988). Prior to that time, distance-dependent luminosities were determined for the bipolar nebulae CRL 2688 and CRL 618 (Ney et al. 1975; Westbrook et al. 1975), based on the limited thermal–IR data available, and before the nature of these objects was entirely clear (Cohen et al. 1975; Humphreys, Warner & Gallagher 1976; Zuckerman et al. 1976; Cohen & Kuhi 1977). The first real attempt at measuring SED-based distances for an ensemble of PPNe was by Zuckerman (1978). Now, the availability of flux data supplied by the Toruń catalogue, as well as additional recent flux data, provides us with the necessary tools to use the same SED modelling technique as for example demonstrated in van der Veen et al. (1989), Kwok et al. (1989), Su, Hrivnak & Kwok (2001), De Ruyter et al. (2005, 2006), Sahai et al. (2007) and Gielen et al. (2011). This method requires the assumption of a ‘standard-candle’ luminosity, statistically representative of PAGB stars, though with caveats, as we explain in detail below.
Most PPNe are seen to have a double-peaked SED comprised of a central star component and a cool dust component arising from the reprocessing of UV photons from the central star into IR photons from the detached circumstellar envelope, as in Kwok (1993). In some cases, the presence of either a cool stellar companion or a Keplerian dust disc can be seen by an increase in the NIR flux (Dominik et al. 2003; De Ruyter et al. 2005; Van Winckel et al. 2006, 2009; Hillen et al. 2013). We assumed in all cases that the observed SED can be expressed as a combination of one or more Planck functions. No object required more than three Planck functions to model the observed IR energy distribution.
Furthermore, due to the large amount of data on more than 200 catalogued stars, and the resulting statistical nature of the project, a number of simplifying assumptions were necessary in our approach. We used a combination of Planck functions at different temperatures rather than taking a more detailed radiative transfer approach, such as was done for the Red Rectangle by Men'shchikov et al. (2002). We only used broad-band continuum fluxes, ignoring spectroscopic features such as silicate absorption and emission bands and any fine-structure lines (e.g. Engelke, Kraemer & Price 2004). Finally, in order to calculate the total distance-dependent flux of each object, we have assumed an isotropic flux distribution, ignoring both the nebular morphology and orientation of each object which would inherently alter the flux distribution (Su et al. 2001).
3.1 Assumed luminosity
It has been shown several times in the literature (Vennes et al. 2002; Liebert, Bergeron & Holberg 2005; Kepler et al. 2007) that thin-disc WDs, which are the direct progeny of most PNe and their precursor PAGB stars, have a narrow mass distribution around ∼0.6 M⊙. Tremblay, Bergeron & Gianninas (2011) found a mean mass of DA WDs of 0.613 M⊙ from Sloan Digital Sky Survey (SDSS; Adelman-McCarthy et al. 2006) DR4 data. In a more recent survey of ∼2200 hot (Teff > 13 000 K) DA WDs found in the SDSS, Kleinman et al. (2013) found that the disc DA WD population can be divided into three distinct populations with mass peaks at 0.39, 0.59 and 0.82 M⊙, and population fractions of 6, 81 and 13 per cent, respectively. The lower mass peak can be excluded from further analysis here, on the grounds that single-star evolution cannot account for such low-mass remnants within a Hubble time. The higher mass WDs have also been suggested to be a product of binary evolution, and possibly represent WD mergers (Jeffery, Karakas & Saio 2011; Kleinman et al. 2013), though the WDs derived from single higher mass progenitors will also contribute to this mass bin. Independently, the empirical mass distribution of PN central stars has also been estimated by Stasińska, Gorny & Tylenda (1997) and Gesicki & Zijlstra (2007), who determine a range of 0.55–0.65 M⊙ and 0.61 ± 0.02 M⊙, respectively.
In this paper, we adopt a mass of 0.61 ± 0.02 M⊙ to represent the mean core mass during the PAGB evolutionary stage for the nearby intermediate-age population of the thin disc, which we expect produced the majority of local PNe (Frew & Parker 2006; Frew 2008). Using the core mass/luminosity relation of VW94, this translates to an approximate PAGB luminosity of L⋆ ≈ 6000 ± 1500 L⊙. This does not imply that all catalogued PPNe have this luminosity, just that they are statistically likely to fall in this luminosity range. Indeed, given the wide range of PN morphologies and ionized masses (Frew & Parker 2010), and the diversity of central star photospheric compositions observed in the solar neighbourhood (e.g. Frew 2008; Frew et al. 2014a), this may be a somewhat simplistic approach. Interestingly, the luminosity distribution function for PAGB stars observed in the Large Magellanic Cloud (LMC) by van Aarle et al. (2011) is much broader, between adopted extrema of ∼1000 and 35 000 L⊙. This result is possibly due to contamination issues, and uncertain selection biases at play manifest when comparing a local disc-dominated sample with a colour-selected and flux-limited LMC sample.
In light of the Kleinman et al. results, we endeavour to divide the Galactic population of PAGB stars into several subpopulations of different ages, to refine our distance estimates, especially for objects at low Galactic latitude. The high-mass objects are potentially the objects with the largest distance uncertainties. The extreme OH/IR stars are generally considered to be the descendants of the highest mass progenitors (e.g. Justtanont et al. 2013). These are evolved, dust-enshrouded AGB stars with very strong 1612 MHz OH maser emission (Johansson et al. 1977; te Lintel Hekkert et al. 1989; Sevenster 2002), and are thought to produce the small homogeneous class of OHPNe, after photoionization has commenced (Zijlstra et al. 1989; Uscanga et al. 2012).
García-Hernández et al. (2007) propose that heavily obscured OHPNe descend from such high-mass progenitor stars, which could represent a link between OH/IR stars with extreme outflows and collimated bipolar PN. García-Hernández et al. (2007) adopted L = 10 000 L⊙ for this class of stars. However, on the assumption that hot-bottom burning (HBB) of dredged-up carbon to nitrogen only occurs in stars heavier than 4.0–5.0 M⊙ (Boothroyd, Sackmann & Ahern 1993; Mazzitelli, D'Antona & Ventura 1999; Izzard et al. 2004; McSaveney et al. 2007; Karakas et al. 2009), then at a minimum, L = 20 000 L⊙ is more appropriate for the oxygen-rich AGB stars that are produced. On the other hand, there is some evidence that substantial nitrogen enrichment, at least at the amounts needed to make a type I PN (N/O > 0.8, following Kingsburgh & Barlow 1994), may occur at masses less than this, down to stars with initial masses of ∼3.0 M⊙, or even less (Karakas et al. 2009; Parker et al. 2011).
Accordingly, we divide the young and intermediate-age (carbon star) populations at an initial mass of 3.0 M⊙, corresponding to a turn-off age of ∼5× 109 yr (cf. Wood, Bessell & Fox 1983). However, recent observations in the inner disc of M31 (Boyer et al. 2013) suggest that there is a metallicity threshold above which carbon stars (with C/O > 1) cannot be produced through dredge-up processes. If the same threshold exists in the inner disc and bulge of our Galaxy, then the presence of oxygen-rich stars there cannot be used as an indicator of high-mass progenitors. In light of this, we use a high luminosity (>20 000L⊙) only for objects with an unambiguous isotopic signature of HBB, i.e. with low values of the 18O/17O and 12C/13C ratios (Imai et al. 2012; Edwards & Ziurys 2013; Justtanont et al. 2013), or a high N/O ratio in shock-ionized optical knots. Other criteria, such as nebular expansion velocities (e.g. Barnbaum, Zuckerman & Kastner 1991), need caution in their interpretation.
For the low-mass Galactic thick-disc objects we have identified (from kinematics and/or distances from the disc mid-plane), we adopt an approximate age of 10 Gyr (Ramírez & Allende Prieto 2011; Hansen et al. 2013) and a progenitor mass of 1.0 M⊙. Using an average of the WD initial-to-final mass relations (IFMRs) of Kalirai et al. (2008) and Renedo et al. (2010), we then estimate a core mass of 0.53 M⊙, leading to an assumed luminosity of 1700 L⊙ (following VW94). For the more metal poor Population II stars of the Galactic halo, we assume an age of 11–13 Gyr (Kalirai 2012; Hansen et al. 2013) and that the 0.8 M⊙ progenitor stars produced WDs with a mass of 0.53–0.55 M⊙ (Kalirai et al. 2009; Kalirai 2012). From this mass, we estimate a mean PAGB luminosity of 2000 L⊙ (VW94) for this population. This luminosity is consistent with observations of ‘yellow’ PAGB stars in globular clusters (Alves, Bond & Onken 2001).
For PAGB objects in the direction of the Galactic bulge, we initially assume potential membership for objects within 10 deg of the Galactic Centre. However, if the calculated distance is less than 6 kpc, then bulge membership is excluded and we re-calculate the distance on the assumption of disc membership (see above). When considering an appropriate luminosity to adopt for bulge PAGB objects, the situation is less secure due to the complex formation history of the bulge. Bensby et al. (2013) found a population of old (>10 Gyr) metal-poor stars with several groups of younger higher metallicity stars with a broad distribution of ages, some as young as 2 Gyr (Bensby et al. 2013), corresponding to a progenitor mass of ∼1.8 M⊙. The age distribution of the metal-rich component peaks around 4–5 Gyr, with a progenitor mass of ∼1.4 M⊙. On the assumption that the observed bulge population of PAGB stars and PNe derive from 1.6 ± 0.2 M⊙ stars, we infer an average core mass of 0.57 M⊙, and hence a PAGB luminosity of 4000 L⊙, again following VW94. Since the bulge sample of compact, optically thick PNe we have also analysed (Section 5.1) is dominated by helium-burning [WCL] stars, it would be useful to have an understanding of their mean mass. As far as we know, there is no literature determination of the mean mass of these stars, but Althaus et al. (2009) have measured the average mass of a sample of 37 PG 1159 stars, the possible descendants of the [WCL] stars, to be 0.573 M⊙, in good agreement with our assumed bulge mass. We hence assume a luminosity 4000 L⊙ for all [WCL] stars, including those that belong to the Galactic disc.9
The adopted parameters of the subpopulations of the Galaxy used in this study are summarized in Table 3. The ages and metallicities have been adopted (or derived) from Carollo et al. (2007, 2010), Fuhrmann (2011), Ramírez & Allende Prieto (2011), Kalirai (2012), Bensby et al. (2013) and Hansen et al. (2013). We also give the inferred luminosities derived from the progenitor masses via an averaged IFMR (see the above references) and the core-luminosity relation of VW94, and we attempt to map our populations to the classes of Peimbert (1978), expected when the objects evolve into PNe (see also Casassus & Roche 2001; Quireza, Rocha-Pinto & Maciel 2007). Owing to various assumptions that differ between these studies, this can only be done approximately.
Population . | Age . | Metallicity range . | Mass range . | Initial massa . | PAGB Mass . | PAGB Luminosity . | Classb . |
---|---|---|---|---|---|---|---|
. | (Gyr) . | [Fe/H] . | (M⊙) . | (M⊙) . | (M⊙) . | (L⊙) . | . |
Young thin disc | ≤ 0.7 | −0.2 to +0.5 | 3.0–8.0 | 3.0 | 0.70 | 12000 | I |
Intermediate thin disc | 0.7–3.0 | −0.2 to +0.5 | 1.6–3.0 | 2.0 | 0.61 | 6000 | IIa |
Old thin disc | 3.0–8.0 | −0.7 to +0.5 | 1.1–1.6 | 1.5 | 0.56 | 3500 | IIb |
Thick disc | 8–12 | −1.6 to −0.3 | 0.9–1.1 | 1.0 | 0.53 | 1700 | III |
Bulge | 2.0–12 | −1.9 to +0.6 | 1.0–1.8 | 1.6 | 0.57 | 4000 | V |
Inner halo | 11–13 | −2.4 to −0.6 | ∼0.8 | 0.8 | 0.54 | 1700 | IV |
Outer halo | 11–13 | ≤−1.5 | ∼0.8 | 0.8 | 0.53: | 1700 | IV |
Population . | Age . | Metallicity range . | Mass range . | Initial massa . | PAGB Mass . | PAGB Luminosity . | Classb . |
---|---|---|---|---|---|---|---|
. | (Gyr) . | [Fe/H] . | (M⊙) . | (M⊙) . | (M⊙) . | (L⊙) . | . |
Young thin disc | ≤ 0.7 | −0.2 to +0.5 | 3.0–8.0 | 3.0 | 0.70 | 12000 | I |
Intermediate thin disc | 0.7–3.0 | −0.2 to +0.5 | 1.6–3.0 | 2.0 | 0.61 | 6000 | IIa |
Old thin disc | 3.0–8.0 | −0.7 to +0.5 | 1.1–1.6 | 1.5 | 0.56 | 3500 | IIb |
Thick disc | 8–12 | −1.6 to −0.3 | 0.9–1.1 | 1.0 | 0.53 | 1700 | III |
Bulge | 2.0–12 | −1.9 to +0.6 | 1.0–1.8 | 1.6 | 0.57 | 4000 | V |
Inner halo | 11–13 | −2.4 to −0.6 | ∼0.8 | 0.8 | 0.54 | 1700 | IV |
Outer halo | 11–13 | ≤−1.5 | ∼0.8 | 0.8 | 0.53: | 1700 | IV |
Notes. aAdopted mass.
Population . | Age . | Metallicity range . | Mass range . | Initial massa . | PAGB Mass . | PAGB Luminosity . | Classb . |
---|---|---|---|---|---|---|---|
. | (Gyr) . | [Fe/H] . | (M⊙) . | (M⊙) . | (M⊙) . | (L⊙) . | . |
Young thin disc | ≤ 0.7 | −0.2 to +0.5 | 3.0–8.0 | 3.0 | 0.70 | 12000 | I |
Intermediate thin disc | 0.7–3.0 | −0.2 to +0.5 | 1.6–3.0 | 2.0 | 0.61 | 6000 | IIa |
Old thin disc | 3.0–8.0 | −0.7 to +0.5 | 1.1–1.6 | 1.5 | 0.56 | 3500 | IIb |
Thick disc | 8–12 | −1.6 to −0.3 | 0.9–1.1 | 1.0 | 0.53 | 1700 | III |
Bulge | 2.0–12 | −1.9 to +0.6 | 1.0–1.8 | 1.6 | 0.57 | 4000 | V |
Inner halo | 11–13 | −2.4 to −0.6 | ∼0.8 | 0.8 | 0.54 | 1700 | IV |
Outer halo | 11–13 | ≤−1.5 | ∼0.8 | 0.8 | 0.53: | 1700 | IV |
Population . | Age . | Metallicity range . | Mass range . | Initial massa . | PAGB Mass . | PAGB Luminosity . | Classb . |
---|---|---|---|---|---|---|---|
. | (Gyr) . | [Fe/H] . | (M⊙) . | (M⊙) . | (M⊙) . | (L⊙) . | . |
Young thin disc | ≤ 0.7 | −0.2 to +0.5 | 3.0–8.0 | 3.0 | 0.70 | 12000 | I |
Intermediate thin disc | 0.7–3.0 | −0.2 to +0.5 | 1.6–3.0 | 2.0 | 0.61 | 6000 | IIa |
Old thin disc | 3.0–8.0 | −0.7 to +0.5 | 1.1–1.6 | 1.5 | 0.56 | 3500 | IIb |
Thick disc | 8–12 | −1.6 to −0.3 | 0.9–1.1 | 1.0 | 0.53 | 1700 | III |
Bulge | 2.0–12 | −1.9 to +0.6 | 1.0–1.8 | 1.6 | 0.57 | 4000 | V |
Inner halo | 11–13 | −2.4 to −0.6 | ∼0.8 | 0.8 | 0.54 | 1700 | IV |
Outer halo | 11–13 | ≤−1.5 | ∼0.8 | 0.8 | 0.53: | 1700 | IV |
Notes. aAdopted mass.
Based on the observed mass function of PN central stars, we consider it likely that almost all observed disc PAGB stars in the Toruń catalogue have luminosities between 3500 and 10 000 L⊙, with only a few more luminous objects. Hence, at worst, the error on an individual distance is about 40 per cent, and significantly smaller if our classification into subgroups has been effective.
3.2 SED fitting and distance determination
Dusty PPNe typically exhibit a double-peaked SED (Kwok 1993) and so would require two Planck functions to model the total flux of the object. However, some objects, e.g. those objects with circumstellar discs (e.g. Waters et al. 1993; De Ruyter et al. 2006; Van Winckel et al. 2006), require additional functions. We have modelled the SEDs of all 209 likely PAGB objects in the Toruń catalogue, as well as a number of candidate PPNe and young compact PNe (see Appendix A1), by fitting Planck functions to the observational data described in Section 2.
For the young PNe, we have used the same method as above except for objects in the Galactic plane (|$|b| < 4\,\deg$|). For these objects, we have used the extinctions derived from the Balmer decrement, taken from Tylenda et al. (1992) and Frew et al. (2013). We have adopted uncertainties ranging from 10 to 30 per cent for the Tylenda et al. (1992) values, and adopt the reddening uncertainties unchanged from Frew et al. (2013). The SEDs have been dereddened using the extinction law of Cardelli, Clayton & Mathis (1989) with the updated NIR coefficients of O'Donnell (1994). Here, we have assumed RV = 3.09. Fig. 1 shows representative SEDs and our model fits for nine representative PAGB objects: IRAS 06176−1036 (the Red Rectangle), IRAS 10158−2844 (HR 4049), IRAS 10256−5628, IRAS 11339−6004, IRAS 11544−6408, IRAS 13416−6243, IRAS 18071−1727, IRAS 18450-0148 and IRAS 19480+2504. The SED plots for all 209 likely and 87 possible PAGB objects are given in the online supplement.
3.3 Other derived parameters
The dereddened photospheric temperature and the dust component temperatures were derived directly from the fitting process, as solutions to the least-squares problem defined in Section 3.2, along with their corresponding 1σ uncertainties, which in some cases are quite large due to the inability to constrain one side of the Planck function. For those objects with a Planck function clearly representative of the central star, i.e. with a blackbody temperature higher than an approximate sublimation temperature of Tsub ∼ 1300 K for silicate dust (Kama, Min & Dominik 2009), we have calculated the ratio of the IR/dust flux to that of the central star, FIR/F⋆. We define the FIR/F⋆ ratio as the ratio of the integrated flux from the photospheric Planck function(s), integrated from 1000 Å to infinity, to the integrated flux of the dust component(s) also integrated from 1000 Å to infinity. The ratio of IR to stellar flux was calculated after the application of the reddening correction.
For objects with a bipolar morphology, FIR/F⋆ can be used as an indication of the orientation of the object whether it be pole-on or edge-on (Su et al. 2001). We also determined the height |z| = Dsin b above the Galactic mid-plane for each star from our distance and the Galactic latitude, to help ascertain the kinematic population that each object belongs to. A full analysis of the various PAGB subsamples will be given in a future paper.
4 DISTANCE CATALOGUE
In this section, we provide a catalogue of distances to the PAGB objects listed in the Toruń catalogue. In the past, the distances to PAGB objects were difficult to determine, due to the near total obscuration of the central star by their surrounding nebula and the inherent difficulty in their classification. The purpose of this work is to present a more accurate and homogenized SED method of calculating distances to PAGB objects. For the SED method to be applicable (without reddening correction), the objects require a larger IR flux compared to the central star flux i.e. FIR/F⋆ > 1. For objects with no IR excess, the derived distances are questionable without applying an extinction correction to the stellar flux. In this section, we provide a catalogue of distances to the PAGB objects listed in the Toruń catalogue.
Table 4 contains the first 10 rows of the SED derived parameters for the 209 likely PAGB objects. The entire table is available in the associate online supplement. Columns 1 and 2 are the IRAS identifier and the other name listed in the Toruń catalogue based on an order of preference given in Szczerba et al. (2007). The longitude and latitude for each object are given in columns 3 and 4. The total integrated fluxes of each star are expressed in units of erg s−1 cm−2 and L⊙ kpc−2 (assuming an isotropic flux distribution) and are presented in columns 5 and 6. In column 7, we give the adopted luminosity of the source, and the estimated reddening for each object is given in column 8. The derived distance and calculated uncertainty are given in column 9. The ratio of IR to stellar flux after reddening correction is listed in column 10. Finally, we have listed the dust temperature and uncertainty if only one blackbody was applied, or a range of dust temperatures if multiple blackbodies were applied, up to the sublimation temperature of astrophysical silicates, Tsub ∼ 1300 K as given by Kama et al. (2009). We have also estimated distances in an identical fashion for the 87 possible PAGB objects in the Toruń catalogue, presented in Table 5. The meanings of the column headings are identical to those in Table 4.
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
17581−2926 | GLMP 688 | 1.293 | −3.199 | 1.46E−09 ± 2.60E−10 | 45 ± 8 | 4000 ± 1500 | 0.51 ± 0.05 | 9.38 ± 1.95 | 1.49 | 113±6 |
17291−2402 | GLMP 575 | 2.518 | 5.120 | 4.22E−09 ± 5.85E−10 | 131 ± 18 | 4000 ± 1500 | 1.16 ± 0.23 | 5.52 ± 1.10 | 8.99 | 130–853 |
17349−2444 | GLMP 593 | 2.652 | 3.637 | 2.17E−09 ± 3.81E−10 | 67 ± 12 | 4000 ± 1500 | 0.51 ± 0.05 | 7.70 ± 1.59 | 9.61 | 121±6 |
18371−3159 | LSE 63 | 2.918 | −11.818 | 1.92E−09 ± 3.99E−10 | 60 ± 12 | 1700 ± 750 | 0.13 ± 0.01 | 5.34 ± 1.30 | 0.77 | 134±7 |
17576−2653 | – | 3.472 | −1.853 | 2.57E−09 ± 3.66E−10 | 80 ± 11 | 4000 ± 1500 | 0.51 ± 0.05 | 7.07 ± 1.42 | 2.32 | 187±9 |
17516−2525 | GLMP 662 | 4.038 | 0.056 | 4.91E−08 ± 7.06E−09 | 1528 ± 219 | 6000 ± 1500 | 0.50 ± 0.05 | 1.98 ± 0.29 | 7.07 | 141–749 |
17074−1845 | LSE 3 | 4.100 | 12.263 | 4.07E−09 ± 9.59E−10 | 127 ± 30 | 6000 ± 1500 | 0.25 ± 0.02 | 6.88 ± 1.18 | 0.51 | 146±7 |
17441−2411 | Silkworm Nebula | 4.223 | 2.145 | 3.27E−08 ± 3.97E−09 | 1018 ± 123 | 12 000 ± 3000 | 1.65 ± 0.43 | 3.43 ± 0.48 | 17 | 208–962 |
17332−2215 | GLMP 588 | 4.542 | 5.295 | 2.52E−09 ± 3.92E−10 | 78 ± 12 | 4000 ± 1500 | 0.78 ± 0.16 | 7.14 ± 1.45 | 8.97 | 137–833 |
17360−2142 | GLMP 600 | 5.364 | 5.038 | 2.20E−09 ± 3.77E−10 | 69 ± 12 | 4000 ± 1500 | 0.72 ± 0.14 | 7.64 ± 1.57 | 4.12 | 141±7 |
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
17581−2926 | GLMP 688 | 1.293 | −3.199 | 1.46E−09 ± 2.60E−10 | 45 ± 8 | 4000 ± 1500 | 0.51 ± 0.05 | 9.38 ± 1.95 | 1.49 | 113±6 |
17291−2402 | GLMP 575 | 2.518 | 5.120 | 4.22E−09 ± 5.85E−10 | 131 ± 18 | 4000 ± 1500 | 1.16 ± 0.23 | 5.52 ± 1.10 | 8.99 | 130–853 |
17349−2444 | GLMP 593 | 2.652 | 3.637 | 2.17E−09 ± 3.81E−10 | 67 ± 12 | 4000 ± 1500 | 0.51 ± 0.05 | 7.70 ± 1.59 | 9.61 | 121±6 |
18371−3159 | LSE 63 | 2.918 | −11.818 | 1.92E−09 ± 3.99E−10 | 60 ± 12 | 1700 ± 750 | 0.13 ± 0.01 | 5.34 ± 1.30 | 0.77 | 134±7 |
17576−2653 | – | 3.472 | −1.853 | 2.57E−09 ± 3.66E−10 | 80 ± 11 | 4000 ± 1500 | 0.51 ± 0.05 | 7.07 ± 1.42 | 2.32 | 187±9 |
17516−2525 | GLMP 662 | 4.038 | 0.056 | 4.91E−08 ± 7.06E−09 | 1528 ± 219 | 6000 ± 1500 | 0.50 ± 0.05 | 1.98 ± 0.29 | 7.07 | 141–749 |
17074−1845 | LSE 3 | 4.100 | 12.263 | 4.07E−09 ± 9.59E−10 | 127 ± 30 | 6000 ± 1500 | 0.25 ± 0.02 | 6.88 ± 1.18 | 0.51 | 146±7 |
17441−2411 | Silkworm Nebula | 4.223 | 2.145 | 3.27E−08 ± 3.97E−09 | 1018 ± 123 | 12 000 ± 3000 | 1.65 ± 0.43 | 3.43 ± 0.48 | 17 | 208–962 |
17332−2215 | GLMP 588 | 4.542 | 5.295 | 2.52E−09 ± 3.92E−10 | 78 ± 12 | 4000 ± 1500 | 0.78 ± 0.16 | 7.14 ± 1.45 | 8.97 | 137–833 |
17360−2142 | GLMP 600 | 5.364 | 5.038 | 2.20E−09 ± 3.77E−10 | 69 ± 12 | 4000 ± 1500 | 0.72 ± 0.14 | 7.64 ± 1.57 | 4.12 | 141±7 |
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
17581−2926 | GLMP 688 | 1.293 | −3.199 | 1.46E−09 ± 2.60E−10 | 45 ± 8 | 4000 ± 1500 | 0.51 ± 0.05 | 9.38 ± 1.95 | 1.49 | 113±6 |
17291−2402 | GLMP 575 | 2.518 | 5.120 | 4.22E−09 ± 5.85E−10 | 131 ± 18 | 4000 ± 1500 | 1.16 ± 0.23 | 5.52 ± 1.10 | 8.99 | 130–853 |
17349−2444 | GLMP 593 | 2.652 | 3.637 | 2.17E−09 ± 3.81E−10 | 67 ± 12 | 4000 ± 1500 | 0.51 ± 0.05 | 7.70 ± 1.59 | 9.61 | 121±6 |
18371−3159 | LSE 63 | 2.918 | −11.818 | 1.92E−09 ± 3.99E−10 | 60 ± 12 | 1700 ± 750 | 0.13 ± 0.01 | 5.34 ± 1.30 | 0.77 | 134±7 |
17576−2653 | – | 3.472 | −1.853 | 2.57E−09 ± 3.66E−10 | 80 ± 11 | 4000 ± 1500 | 0.51 ± 0.05 | 7.07 ± 1.42 | 2.32 | 187±9 |
17516−2525 | GLMP 662 | 4.038 | 0.056 | 4.91E−08 ± 7.06E−09 | 1528 ± 219 | 6000 ± 1500 | 0.50 ± 0.05 | 1.98 ± 0.29 | 7.07 | 141–749 |
17074−1845 | LSE 3 | 4.100 | 12.263 | 4.07E−09 ± 9.59E−10 | 127 ± 30 | 6000 ± 1500 | 0.25 ± 0.02 | 6.88 ± 1.18 | 0.51 | 146±7 |
17441−2411 | Silkworm Nebula | 4.223 | 2.145 | 3.27E−08 ± 3.97E−09 | 1018 ± 123 | 12 000 ± 3000 | 1.65 ± 0.43 | 3.43 ± 0.48 | 17 | 208–962 |
17332−2215 | GLMP 588 | 4.542 | 5.295 | 2.52E−09 ± 3.92E−10 | 78 ± 12 | 4000 ± 1500 | 0.78 ± 0.16 | 7.14 ± 1.45 | 8.97 | 137–833 |
17360−2142 | GLMP 600 | 5.364 | 5.038 | 2.20E−09 ± 3.77E−10 | 69 ± 12 | 4000 ± 1500 | 0.72 ± 0.14 | 7.64 ± 1.57 | 4.12 | 141±7 |
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
17581−2926 | GLMP 688 | 1.293 | −3.199 | 1.46E−09 ± 2.60E−10 | 45 ± 8 | 4000 ± 1500 | 0.51 ± 0.05 | 9.38 ± 1.95 | 1.49 | 113±6 |
17291−2402 | GLMP 575 | 2.518 | 5.120 | 4.22E−09 ± 5.85E−10 | 131 ± 18 | 4000 ± 1500 | 1.16 ± 0.23 | 5.52 ± 1.10 | 8.99 | 130–853 |
17349−2444 | GLMP 593 | 2.652 | 3.637 | 2.17E−09 ± 3.81E−10 | 67 ± 12 | 4000 ± 1500 | 0.51 ± 0.05 | 7.70 ± 1.59 | 9.61 | 121±6 |
18371−3159 | LSE 63 | 2.918 | −11.818 | 1.92E−09 ± 3.99E−10 | 60 ± 12 | 1700 ± 750 | 0.13 ± 0.01 | 5.34 ± 1.30 | 0.77 | 134±7 |
17576−2653 | – | 3.472 | −1.853 | 2.57E−09 ± 3.66E−10 | 80 ± 11 | 4000 ± 1500 | 0.51 ± 0.05 | 7.07 ± 1.42 | 2.32 | 187±9 |
17516−2525 | GLMP 662 | 4.038 | 0.056 | 4.91E−08 ± 7.06E−09 | 1528 ± 219 | 6000 ± 1500 | 0.50 ± 0.05 | 1.98 ± 0.29 | 7.07 | 141–749 |
17074−1845 | LSE 3 | 4.100 | 12.263 | 4.07E−09 ± 9.59E−10 | 127 ± 30 | 6000 ± 1500 | 0.25 ± 0.02 | 6.88 ± 1.18 | 0.51 | 146±7 |
17441−2411 | Silkworm Nebula | 4.223 | 2.145 | 3.27E−08 ± 3.97E−09 | 1018 ± 123 | 12 000 ± 3000 | 1.65 ± 0.43 | 3.43 ± 0.48 | 17 | 208–962 |
17332−2215 | GLMP 588 | 4.542 | 5.295 | 2.52E−09 ± 3.92E−10 | 78 ± 12 | 4000 ± 1500 | 0.78 ± 0.16 | 7.14 ± 1.45 | 8.97 | 137–833 |
17360−2142 | GLMP 600 | 5.364 | 5.038 | 2.20E−09 ± 3.77E−10 | 69 ± 12 | 4000 ± 1500 | 0.72 ± 0.14 | 7.64 ± 1.57 | 4.12 | 141±7 |
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
– | LS 4825 | 1.671 | −6.628 | 3.09E−09 ± 6.10E−10 | 96 ± 19 | 4000 ± 1500 | 0.24 ± 0.05 | 6.45 ± 1.36 | – | – |
17550−2800 | GLMP 676 | 2.205 | −1.900 | 2.24E−09 ± 4.59E−10 | 70 ± 14 | 4000 ± 1500 | 0.51 ± 0.05 | 7.58 ± 1.62 | – | 125–1030 |
– | CD-30 15602 | 2.798 | −7.675 | 1.26E−09 ± 1.79E−10 | 39 ± 6 | 3500 ± 1500 | 0.22 ± 0.04 | 9.44 ± 2.13 | – | – |
17376−2040 | – | 6.437 | 5.275 | 4.48E−09 ± 1.70E−09 | 139 ± 53 | 4000 ± 1500 | 0.65 ± 0.13 | 5.36 ± 1.43 | 5.71 | 161–1232 |
17416−2112 | GLMP 625 | 6.469 | 4.198 | 2.14E−09 ± 2.34E−10 | 67 ± 7 | 4000 ± 1500 | 0.85 ± 0.17 | 7.75 ± 1.51 | 77 | 130–499 |
16476−1122 | – | 7.524 | 20.418 | 8.32E−09 ± 8.44E−10 | 259 ± 26 | 3500 ± 1500 | 0.70 ± 0.07 | 3.68 ± 0.81 | 0.11 | 186±9 |
F16277−0724 | LS IV -07 1 | 7.956 | 26.706 | 1.36E−07 ± 8.80E−09 | 4234 ± 274 | 3500 ± 1500 | 0.24 ± 0.02 | 0.91 ± 0.20 | – | – |
17433−1750 | GLMP 637 | 9.562 | 5.612 | 5.73E−09 ± 3.22E−09 | 178 ± 100 | 4000 ± 1500 | 0.53 ± 0.11 | 4.74 ± 1.60 | 5.95 | 156–431 |
17364−1238 | – | 13.183 | 9.720 | 8.80E−10 ± 1.05E−10 | 27 ± 3 | 1700 ± 750 | 0.46 ± 0.09 | 7.88 ± 1.80 | 0.28 | 134±7 |
18313−1738 | – | 15.322 | −4.268 | 7.58E−09 ± 2.36E−09 | 236 ± 73 | 6000 ± 1500 | 0.64 ± 0.13 | 5.04 ± 1.01 | 6.39 | 412–1142 |
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
– | LS 4825 | 1.671 | −6.628 | 3.09E−09 ± 6.10E−10 | 96 ± 19 | 4000 ± 1500 | 0.24 ± 0.05 | 6.45 ± 1.36 | – | – |
17550−2800 | GLMP 676 | 2.205 | −1.900 | 2.24E−09 ± 4.59E−10 | 70 ± 14 | 4000 ± 1500 | 0.51 ± 0.05 | 7.58 ± 1.62 | – | 125–1030 |
– | CD-30 15602 | 2.798 | −7.675 | 1.26E−09 ± 1.79E−10 | 39 ± 6 | 3500 ± 1500 | 0.22 ± 0.04 | 9.44 ± 2.13 | – | – |
17376−2040 | – | 6.437 | 5.275 | 4.48E−09 ± 1.70E−09 | 139 ± 53 | 4000 ± 1500 | 0.65 ± 0.13 | 5.36 ± 1.43 | 5.71 | 161–1232 |
17416−2112 | GLMP 625 | 6.469 | 4.198 | 2.14E−09 ± 2.34E−10 | 67 ± 7 | 4000 ± 1500 | 0.85 ± 0.17 | 7.75 ± 1.51 | 77 | 130–499 |
16476−1122 | – | 7.524 | 20.418 | 8.32E−09 ± 8.44E−10 | 259 ± 26 | 3500 ± 1500 | 0.70 ± 0.07 | 3.68 ± 0.81 | 0.11 | 186±9 |
F16277−0724 | LS IV -07 1 | 7.956 | 26.706 | 1.36E−07 ± 8.80E−09 | 4234 ± 274 | 3500 ± 1500 | 0.24 ± 0.02 | 0.91 ± 0.20 | – | – |
17433−1750 | GLMP 637 | 9.562 | 5.612 | 5.73E−09 ± 3.22E−09 | 178 ± 100 | 4000 ± 1500 | 0.53 ± 0.11 | 4.74 ± 1.60 | 5.95 | 156–431 |
17364−1238 | – | 13.183 | 9.720 | 8.80E−10 ± 1.05E−10 | 27 ± 3 | 1700 ± 750 | 0.46 ± 0.09 | 7.88 ± 1.80 | 0.28 | 134±7 |
18313−1738 | – | 15.322 | −4.268 | 7.58E−09 ± 2.36E−09 | 236 ± 73 | 6000 ± 1500 | 0.64 ± 0.13 | 5.04 ± 1.01 | 6.39 | 412–1142 |
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
– | LS 4825 | 1.671 | −6.628 | 3.09E−09 ± 6.10E−10 | 96 ± 19 | 4000 ± 1500 | 0.24 ± 0.05 | 6.45 ± 1.36 | – | – |
17550−2800 | GLMP 676 | 2.205 | −1.900 | 2.24E−09 ± 4.59E−10 | 70 ± 14 | 4000 ± 1500 | 0.51 ± 0.05 | 7.58 ± 1.62 | – | 125–1030 |
– | CD-30 15602 | 2.798 | −7.675 | 1.26E−09 ± 1.79E−10 | 39 ± 6 | 3500 ± 1500 | 0.22 ± 0.04 | 9.44 ± 2.13 | – | – |
17376−2040 | – | 6.437 | 5.275 | 4.48E−09 ± 1.70E−09 | 139 ± 53 | 4000 ± 1500 | 0.65 ± 0.13 | 5.36 ± 1.43 | 5.71 | 161–1232 |
17416−2112 | GLMP 625 | 6.469 | 4.198 | 2.14E−09 ± 2.34E−10 | 67 ± 7 | 4000 ± 1500 | 0.85 ± 0.17 | 7.75 ± 1.51 | 77 | 130–499 |
16476−1122 | – | 7.524 | 20.418 | 8.32E−09 ± 8.44E−10 | 259 ± 26 | 3500 ± 1500 | 0.70 ± 0.07 | 3.68 ± 0.81 | 0.11 | 186±9 |
F16277−0724 | LS IV -07 1 | 7.956 | 26.706 | 1.36E−07 ± 8.80E−09 | 4234 ± 274 | 3500 ± 1500 | 0.24 ± 0.02 | 0.91 ± 0.20 | – | – |
17433−1750 | GLMP 637 | 9.562 | 5.612 | 5.73E−09 ± 3.22E−09 | 178 ± 100 | 4000 ± 1500 | 0.53 ± 0.11 | 4.74 ± 1.60 | 5.95 | 156–431 |
17364−1238 | – | 13.183 | 9.720 | 8.80E−10 ± 1.05E−10 | 27 ± 3 | 1700 ± 750 | 0.46 ± 0.09 | 7.88 ± 1.80 | 0.28 | 134±7 |
18313−1738 | – | 15.322 | −4.268 | 7.58E−09 ± 2.36E−09 | 236 ± 73 | 6000 ± 1500 | 0.64 ± 0.13 | 5.04 ± 1.01 | 6.39 | 412–1142 |
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
– | LS 4825 | 1.671 | −6.628 | 3.09E−09 ± 6.10E−10 | 96 ± 19 | 4000 ± 1500 | 0.24 ± 0.05 | 6.45 ± 1.36 | – | – |
17550−2800 | GLMP 676 | 2.205 | −1.900 | 2.24E−09 ± 4.59E−10 | 70 ± 14 | 4000 ± 1500 | 0.51 ± 0.05 | 7.58 ± 1.62 | – | 125–1030 |
– | CD-30 15602 | 2.798 | −7.675 | 1.26E−09 ± 1.79E−10 | 39 ± 6 | 3500 ± 1500 | 0.22 ± 0.04 | 9.44 ± 2.13 | – | – |
17376−2040 | – | 6.437 | 5.275 | 4.48E−09 ± 1.70E−09 | 139 ± 53 | 4000 ± 1500 | 0.65 ± 0.13 | 5.36 ± 1.43 | 5.71 | 161–1232 |
17416−2112 | GLMP 625 | 6.469 | 4.198 | 2.14E−09 ± 2.34E−10 | 67 ± 7 | 4000 ± 1500 | 0.85 ± 0.17 | 7.75 ± 1.51 | 77 | 130–499 |
16476−1122 | – | 7.524 | 20.418 | 8.32E−09 ± 8.44E−10 | 259 ± 26 | 3500 ± 1500 | 0.70 ± 0.07 | 3.68 ± 0.81 | 0.11 | 186±9 |
F16277−0724 | LS IV -07 1 | 7.956 | 26.706 | 1.36E−07 ± 8.80E−09 | 4234 ± 274 | 3500 ± 1500 | 0.24 ± 0.02 | 0.91 ± 0.20 | – | – |
17433−1750 | GLMP 637 | 9.562 | 5.612 | 5.73E−09 ± 3.22E−09 | 178 ± 100 | 4000 ± 1500 | 0.53 ± 0.11 | 4.74 ± 1.60 | 5.95 | 156–431 |
17364−1238 | – | 13.183 | 9.720 | 8.80E−10 ± 1.05E−10 | 27 ± 3 | 1700 ± 750 | 0.46 ± 0.09 | 7.88 ± 1.80 | 0.28 | 134±7 |
18313−1738 | – | 15.322 | −4.268 | 7.58E−09 ± 2.36E−09 | 236 ± 73 | 6000 ± 1500 | 0.64 ± 0.13 | 5.04 ± 1.01 | 6.39 | 412–1142 |
4.1 Is the 22 μm luminosity a standard candle for dusty PAGB stars?
5 MIMICS
As discussed by Szczerba et al. (2007), several classes of objects can mimic PAGB stars in terms of their optical and IR colours, just as PNe are confused with a wide range of mimics (e.g. Cohen et al. 2007; Frew & Parker 2010; Frew et al. 2010; Parker et al. 2012; De Marco et al. 2013). While H ii regions can be readily differentiated from PAGB stars and PNe (Anderson et al. 2012), some circumstellar nebulae around massive stars (e.g. Wachter et al. 2010), especially the rare class of yellow hypergiants with dusty ejecta, e.g. IRC+10420 (Jones et al. 1993; Oudmaijer & de Wit 2013), Hen 3-1379 (Lagadec et al. 2011b; Hutsemékers, Cox & Vamvatira-Nakou 2013) and probably AFGL 4106 (van Loon et al. 1999), have very similar MIR colours to some PPNe. For other objects, their true nature remains uncertain, e.g. HD 179821 (AFGL 2343) (Hawkins et al. 1995; Kastner & Weintraub 1995; Reddy & Hrivnak 1999; Josselin & Lèbre 2001; Kipper 2008; Ferguson & Ueta 2010). Young stellar objects are the other main contaminant in lists of PAGB stars (e.g. Adams, Lada & Shu 1987; Urquhart et al. 2007) often having quite similar SEDs and MIR colours. We differentiate the two classes on the basis of secondary criteria such as the z-height from the Galactic plane (after first estimating the distance), and the surrounding environment, particularly checking for areas of active star formation.
On the other hand, dusty D-type symbiotic stars (Corradi 1995; Belczyński et al. 2000), and some B[e] stars (Lamers et al. 1998), might be expected to have luminosities of 103–104 L⊙, at least if the accretion-disc luminosity is not too high. Hence, the SED technique should be applicable to these objects too. Several objects with ambiguous classifications, such as the Ant nebula, Mz 3 (Cohen et al. 1978, 2011; Santander-García et al. 2004) and the Calabash Nebula, OH 231.8+4.2 (Section 5.3; Choi et al. 2012), have distances included in Table 4 as we base our classifications on the current edition of the Toruń catalogue. The yellow symbiotic stars, particularly the D′-symbiotic systems (e.g. Gutiérrez-Moreno & Moreno 1998; Jorrisen 2003; Pereira, Smith & Cunha 2005; Frankowski & Jorissen 2007; Miszalski et al. 2012), are more problematic as the optical thickness and the luminosity of the cool star are generally unknown a priori. Thus, we have not determined distances to any of these stars, pending further analysis.
5.1 Derived parameters and distances for 54 miscellaneous nebulae
Several classes of objects are amenable to the SED distance technique, such as the dust-enshrouded AGB stars (Jura & Kleinmann 1989; Olivier, Whitelock & Marang 2001), and some B[e] and D-type symbiotic systems (e.g. Parthasarathy & Pottasch 1989; Miszalski, Mikołajewska & Udalski 2013). As long as the same assumptions apply to these objects as for PAGB stars, such as a typical PAGB luminosity, approximate isotropy and optical thickness, as shown by a large thermal IR excess, then a SED distance can be determined.
Of course the youngest, dustiest PNe (Kwok, Hrivnak & Milone 1986; Zhang & Kwok 1991; Sahai, Morris & Villar 2011) can also be modelled, especially dusty PNe with very late Wolf–Rayet (or [WCL]) and similar central stars (Kwok, Hrivnak & Langill 1993; De Marco, Barlow & Storey 1997; De Marco et al. 2002; Gesicki et al. 2006; De Pew et al. 2011), as these often show a very strong MIR excess (Zijlstra 2001; Clayton et al. 2011). As discussed in Section 3.1, we assume an intrinsic luminosity of 4000 L⊙ for these stars.
Other PNe and transitional objects with signatures of youth such as detectable OH maser emission (Zijlstra et al. 1989; Uscanga et al. 2012) have also been investigated.11 This approach is only approximately correct, as ionized nebulae reprocess the energy of their central stars in other components besides the thermal dust continuum, such as in emission lines and free–free and bound–free continua, which may not be modelled correctly using our SED fitting method. Furthermore, for strongly anisotropic nebulae, this method will give an upper limit of the distance. To assess these caveats, we determined an SED distance to the strongly bipolar type I nebula NGC 6302 (Wright et al. 2011), which exceeds the known distance of 1.17 kpc (Meaburn et al. 2008) by a factor of 2. This large nebula is clearly not optically thick in all directions, illustrated in the difference between the total observed luminosity (5700 L⊙) and the assumed stellar luminosity of 14 300 L⊙ (see Wright et al. 2011).
Lastly, since several well-known objects formerly considered to be PPNe are excluded from the current edition of the Toruń catalogue, such as M2-9 (Allen & Swings 1972; Balick 1989; Hora & Latter 1994; Corradi, Balick & Santander-García 2011; Castro-Carrizo et al. 2012) and Hen 2-90 (Costa, de Freitas-Pacheco & Maciel 1993; Sahai & Nyman 2000; Sahai et al. 2002; Kraus et al. 2005), we have also determined distances to them. We also include the B[e] star MWC 922 (the Red Square; Tuthill & Lloyd 2007), as its nebular morphology and SED are remarkably similar to the better known Red Rectangle (Cohen et al. 2004), as well as the D-type symbiotic outflow Hen 2-147 (Corradi et al. 1999) as it has an independent distance estimate (Santander-García et al. 2007). We have added the final flash objects, FG Sge (included in the Toruń catalogue), as well as V605 Aql which is embedded in recently formed hot dust (Clayton & De Marco 1997; Clayton et al. 2013), and is hosted by the old faint PN, Abell 58 (Abell 1966; Bond 1976). In Table 6, we provide data for these additional sources, primarily taken from our extensive data base of Galactic PNe and related objects (Bojičić et al., in preparation). The meanings of the column headings are identical to those in Table 4. The SED plots for these 54 objects are presented only in the online supplement, following the format of Fig. 1.
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
17574−2921 | H 1-47 | 1.295 | −3.040 | 9.84E−10 ± 3.07E−10 | 31 ± 10 | 4000 ± 1500 | 1.21 ± 0.12 | 11.43 ± 2.79 | 8.50 | 113–337 |
18129−3053 | SwSt 1 | 1.591 | −6.176 | 1.41E−08±3.41E−09 | 437 ± 106 | 4000 ± 1500 | 0.25 ± 0.03 | 3.02 ± 0.68 | 15 | 200–861 |
18022−2822 | M 1-37 | 2.681 | −3.468 | 1.56E−09 ± 6.39E−10 | 48 ± 20 | 4000 ± 1500 | 0.65 ± 0.07 | 9.09 ± 2.52 | 5.72 | 118–349 |
17074−1845 | Hen 3-1347 | 4.100 | 12.263 | 3.33E−09 ± 9.18E−10 | 104 ± 29 | 6000 ± 1500 | 0.25 ± 0.03 | 7.61 ± 1.42 | 0.34 | 146±7 |
19288−3419 | Hen 2-436 | 4.871 | −22.727 | 1.88E−10 ± 5.76E−11 | 5.85 ± 1.79 | 4000 ± 1500 | 0.11 ± 0.02 | 26.16 ± 6.33 | 14 | 267–880 |
18061−2502 | MaC 1-10 | 5.974 | −2.611 | 2.93E−09 ± 4.54E−10 | 91 ± 14 | 4000 ± 1500 | 0.51 ± 0.05 | 6.62 ± 1.34 | 196 | 227–781 |
18170−2416 | H 1-65 | 7.883 | −4.407 | 1.42E−09 ± 4.06E−10 | 44 ± 13 | 4000 ± 1500 | 0.65 ± 0.13 | 9.51 ± 2.24 | 2.24 | 116–356 |
17028−1004 | M 2-9 | 10.900 | 18.055 | 3.22E−08 ± 9.23E−09 | 1002 ± 287 | 2000 ± 750 | 0.50 ± 0.05 | 1.41 ± 0.33 | 146 | 169–767 |
17514−1555 | PM 1-188 | 12.219 | 4.923 | 7.03E−09 ± 2.47E−09 | 219 ± 77 | 4000 ± 1500 | 0.68 ± 0.07 | 4.28 ± 1.10 | 141 | 183–657 |
19016−2330 | GLMP 869 | 13.134 | −13.219 | 8.94E−09 ± 2.43E−09 | 278 ± 76 | 3500 ± 1500 | 0.21 ± 0.02 | 3.55 ± 0.90 | 109 | 196–659 |
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
17574−2921 | H 1-47 | 1.295 | −3.040 | 9.84E−10 ± 3.07E−10 | 31 ± 10 | 4000 ± 1500 | 1.21 ± 0.12 | 11.43 ± 2.79 | 8.50 | 113–337 |
18129−3053 | SwSt 1 | 1.591 | −6.176 | 1.41E−08±3.41E−09 | 437 ± 106 | 4000 ± 1500 | 0.25 ± 0.03 | 3.02 ± 0.68 | 15 | 200–861 |
18022−2822 | M 1-37 | 2.681 | −3.468 | 1.56E−09 ± 6.39E−10 | 48 ± 20 | 4000 ± 1500 | 0.65 ± 0.07 | 9.09 ± 2.52 | 5.72 | 118–349 |
17074−1845 | Hen 3-1347 | 4.100 | 12.263 | 3.33E−09 ± 9.18E−10 | 104 ± 29 | 6000 ± 1500 | 0.25 ± 0.03 | 7.61 ± 1.42 | 0.34 | 146±7 |
19288−3419 | Hen 2-436 | 4.871 | −22.727 | 1.88E−10 ± 5.76E−11 | 5.85 ± 1.79 | 4000 ± 1500 | 0.11 ± 0.02 | 26.16 ± 6.33 | 14 | 267–880 |
18061−2502 | MaC 1-10 | 5.974 | −2.611 | 2.93E−09 ± 4.54E−10 | 91 ± 14 | 4000 ± 1500 | 0.51 ± 0.05 | 6.62 ± 1.34 | 196 | 227–781 |
18170−2416 | H 1-65 | 7.883 | −4.407 | 1.42E−09 ± 4.06E−10 | 44 ± 13 | 4000 ± 1500 | 0.65 ± 0.13 | 9.51 ± 2.24 | 2.24 | 116–356 |
17028−1004 | M 2-9 | 10.900 | 18.055 | 3.22E−08 ± 9.23E−09 | 1002 ± 287 | 2000 ± 750 | 0.50 ± 0.05 | 1.41 ± 0.33 | 146 | 169–767 |
17514−1555 | PM 1-188 | 12.219 | 4.923 | 7.03E−09 ± 2.47E−09 | 219 ± 77 | 4000 ± 1500 | 0.68 ± 0.07 | 4.28 ± 1.10 | 141 | 183–657 |
19016−2330 | GLMP 869 | 13.134 | −13.219 | 8.94E−09 ± 2.43E−09 | 278 ± 76 | 3500 ± 1500 | 0.21 ± 0.02 | 3.55 ± 0.90 | 109 | 196–659 |
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
17574−2921 | H 1-47 | 1.295 | −3.040 | 9.84E−10 ± 3.07E−10 | 31 ± 10 | 4000 ± 1500 | 1.21 ± 0.12 | 11.43 ± 2.79 | 8.50 | 113–337 |
18129−3053 | SwSt 1 | 1.591 | −6.176 | 1.41E−08±3.41E−09 | 437 ± 106 | 4000 ± 1500 | 0.25 ± 0.03 | 3.02 ± 0.68 | 15 | 200–861 |
18022−2822 | M 1-37 | 2.681 | −3.468 | 1.56E−09 ± 6.39E−10 | 48 ± 20 | 4000 ± 1500 | 0.65 ± 0.07 | 9.09 ± 2.52 | 5.72 | 118–349 |
17074−1845 | Hen 3-1347 | 4.100 | 12.263 | 3.33E−09 ± 9.18E−10 | 104 ± 29 | 6000 ± 1500 | 0.25 ± 0.03 | 7.61 ± 1.42 | 0.34 | 146±7 |
19288−3419 | Hen 2-436 | 4.871 | −22.727 | 1.88E−10 ± 5.76E−11 | 5.85 ± 1.79 | 4000 ± 1500 | 0.11 ± 0.02 | 26.16 ± 6.33 | 14 | 267–880 |
18061−2502 | MaC 1-10 | 5.974 | −2.611 | 2.93E−09 ± 4.54E−10 | 91 ± 14 | 4000 ± 1500 | 0.51 ± 0.05 | 6.62 ± 1.34 | 196 | 227–781 |
18170−2416 | H 1-65 | 7.883 | −4.407 | 1.42E−09 ± 4.06E−10 | 44 ± 13 | 4000 ± 1500 | 0.65 ± 0.13 | 9.51 ± 2.24 | 2.24 | 116–356 |
17028−1004 | M 2-9 | 10.900 | 18.055 | 3.22E−08 ± 9.23E−09 | 1002 ± 287 | 2000 ± 750 | 0.50 ± 0.05 | 1.41 ± 0.33 | 146 | 169–767 |
17514−1555 | PM 1-188 | 12.219 | 4.923 | 7.03E−09 ± 2.47E−09 | 219 ± 77 | 4000 ± 1500 | 0.68 ± 0.07 | 4.28 ± 1.10 | 141 | 183–657 |
19016−2330 | GLMP 869 | 13.134 | −13.219 | 8.94E−09 ± 2.43E−09 | 278 ± 76 | 3500 ± 1500 | 0.21 ± 0.02 | 3.55 ± 0.90 | 109 | 196–659 |
IRAS no. . | Other name . | l . | b . | Flux . | Flux . | Luminosity . | E(B−V) . | Distance . | FIR/F⋆ . | TD . |
---|---|---|---|---|---|---|---|---|---|---|
. | . | (°) . | (°) . | (erg s−1 cm−2) . | (L⊙ kpc−2) . | (L⊙) . | (mag) . | (kpc) . | . | (K) . |
17574−2921 | H 1-47 | 1.295 | −3.040 | 9.84E−10 ± 3.07E−10 | 31 ± 10 | 4000 ± 1500 | 1.21 ± 0.12 | 11.43 ± 2.79 | 8.50 | 113–337 |
18129−3053 | SwSt 1 | 1.591 | −6.176 | 1.41E−08±3.41E−09 | 437 ± 106 | 4000 ± 1500 | 0.25 ± 0.03 | 3.02 ± 0.68 | 15 | 200–861 |
18022−2822 | M 1-37 | 2.681 | −3.468 | 1.56E−09 ± 6.39E−10 | 48 ± 20 | 4000 ± 1500 | 0.65 ± 0.07 | 9.09 ± 2.52 | 5.72 | 118–349 |
17074−1845 | Hen 3-1347 | 4.100 | 12.263 | 3.33E−09 ± 9.18E−10 | 104 ± 29 | 6000 ± 1500 | 0.25 ± 0.03 | 7.61 ± 1.42 | 0.34 | 146±7 |
19288−3419 | Hen 2-436 | 4.871 | −22.727 | 1.88E−10 ± 5.76E−11 | 5.85 ± 1.79 | 4000 ± 1500 | 0.11 ± 0.02 | 26.16 ± 6.33 | 14 | 267–880 |
18061−2502 | MaC 1-10 | 5.974 | −2.611 | 2.93E−09 ± 4.54E−10 | 91 ± 14 | 4000 ± 1500 | 0.51 ± 0.05 | 6.62 ± 1.34 | 196 | 227–781 |
18170−2416 | H 1-65 | 7.883 | −4.407 | 1.42E−09 ± 4.06E−10 | 44 ± 13 | 4000 ± 1500 | 0.65 ± 0.13 | 9.51 ± 2.24 | 2.24 | 116–356 |
17028−1004 | M 2-9 | 10.900 | 18.055 | 3.22E−08 ± 9.23E−09 | 1002 ± 287 | 2000 ± 750 | 0.50 ± 0.05 | 1.41 ± 0.33 | 146 | 169–767 |
17514−1555 | PM 1-188 | 12.219 | 4.923 | 7.03E−09 ± 2.47E−09 | 219 ± 77 | 4000 ± 1500 | 0.68 ± 0.07 | 4.28 ± 1.10 | 141 | 183–657 |
19016−2330 | GLMP 869 | 13.134 | −13.219 | 8.94E−09 ± 2.43E−09 | 278 ± 76 | 3500 ± 1500 | 0.21 ± 0.02 | 3.55 ± 0.90 | 109 | 196–659 |
5.2 Comparison with independent distances
To assess the SED technique as a viable distance tool, we undertook a comparison of our estimated distances with the highest quality, independent literature distances that were available. These distances are derived from several primary techniques, e.g. trigonometric and expansion parallaxes. We have compiled a sample of PAGB stars, supplemented with several (nascent) PPNe12 and young compact PNe, which have independently determined distances that we have deemed reliable.
We have omitted from the comparison the trigonometric distance calculated by Imai et al. (2011) for IRAS 19312+1950, as it is not clear if this is a bona fide PAGB star (Nakashima et al. 2011). We have ignored for now any kinematic distances derived from radial velocities and the assumption of circular motion around the Galaxy (e.g. Sahai, Sánchez Contreras & Morris 2005). We justify this by noting that the peculiar velocities of many PAGB stars can be large, which will lead to inaccurate distances from this method. Only for PPNe with high-mass progenitors will this approach work (see also Frew, Parker & Bojičić 2014b).
The literature objects that we have used for the distance comparisons are given in Table 7 along with their literature distances and the corresponding SED derived distance from this work, taken from Tables 4–6. Extended notes on some of these objects are given in Section 5.3. The sample of literature distances includes one object IRAS 07399−1435 (more commonly known as OH 231.8+4.2, Calabash, or Rotten Egg nebula) found in the possible section of the Toruń catalogue as well as three likely PAGB objects. The others are taken from Table 6. The distances for the globular clusters M13 and ω Cen have been taken from the 2010 edition of the Harris (1996) catalogue. Our distance for the PAGB star in NGC 6712 is imprecise, owing to a lack of good photometry, so we have excluded it from the distance comparison. For the dust-enshrouded Miras and symbiotic outflows, we have determined the luminosity for the SED fit from the P–L relations of Feast et al. (1989) and Groenewegen & Whitelock (1996) for O-rich and C-rich stars, respectively.
IRAS no. . | Common name . | L⋆ (L⊙) . | DSED (kpc) . | DLit (kpc) . | Method . | Reference . |
---|---|---|---|---|---|---|
(nascent)-PPNe / PAGB objects . | ||||||
09452+1330 | CW Leo (IRC+10216) | 9800 ± 3000 | 0.13 ± 0.03 | 0.123 ± 0.014 | G | Groenewegen et al. (2012) |
03507+1115 | IK Tau (NML Tau) | 9100 ± 3000 | 0.24 ± 0.06 | 0.265 ± 0.02 | E | Hale et al. (1997) |
21003+3630 | CRL 2688 | 6000 ± 1500 | 0.58 ± 0.08 | 0.42 ± 0.06 | E | Ueta, Murakawa & Meixner (2006) |
10158−2844 | HR 4049 | 3500 ± 1500 | 0.64 ± 0.15 | 0.64 ± 0.19 | D | Acke et al. (2013) |
18286−0959 | OH 21.80-0.13 | 6000 ± 1500 | 3.10 ± 0.45 | |$3.61^{+0.63}_{-0.47}$| | T | Imai et al. (2013a) |
18460−0151 | OH 31.0-0.2 | 6000 ± 1500 | 3.14 ± 0.80 | 2.1 ± 0.6 | S | Imai et al. (2013b) |
– | NGC 5139 WOR 1957 | 1700 ± 750 | 6.13 ± 1.36 | 5.2 ± 0.5 | M | Harris (1996) |
17423−1755 | Hen 3-1475 | 20 000 ± 5000 | 7.17 ± 1.02 | 8.30( ± 1.0) | E | Borkowski & Harrington (2001) |
– | NGC 6205 BARN 29 | 1700 ± 750 | 7.97 ± 1.90 | 7.1 ± 0.7 | M | Harris (1996) |
19134+2131 | [HSD93b] 16 | 6000 ± 1500 | 8.37 ± 1.21 | |$8.00^{+0.90}_{-0.70}$| | T | Imai, Sahai & Morris (2007) |
Young PNe | ||||||
– | NGC 7027 | 6000 ± 1500 | 0.95 ± 0.15 | 0.87 ± 0.10b | E | Zijlstra, van Hoof & Perley (2008, and references therein) |
19327+3024 | BD+30 3639 | 4000 ± 1500 | 1.25 ± 0.25 | 1.52 ± 0.21 | V | Akras & Steffen (2012) |
18240−0244 | M 2-43 | 4000 ± 1500 | 3.68 ± 0.75 | 6.9 ± 1.5 | E | Guzmán, Gómez & Rodríguez (2006) |
19219+0947 | Vy 2-2 | 6000 ± 1500 | 3.74 ± 0.60 | 3.6 ± 0.4 | E | Christianto & Seaquist (1998) |
19255+2123 | K 3-35 | 6000 ± 1500 | 5.77 ± 0.83 | |$3.90^{+0.70}_{-0.50}$| | T | Tafoya et al. (2011) |
19288−3419 | Hen 2-436 | 4000 ± 1500 | 26.2 ± 6.3 | 26.0 ± 2.0c | M | Zijlstra et al. (2006) |
Symbiotic stars/Other objects | ||||||
17028−1004 | M 2-9 | 2000 ± 750 | 1.41 ± 0.33 | 1.3 ±0.2 | G | Corradi et al. (2011) |
20097+2010 | Hen 1-5 (FG Sge) | 6000 ± 1500 | 2.08 ± 0.52 | 2.5 (±0.5) | P | Mayor & Acker (1980) |
07399−1435 | OH 231.8+4.2a | 15 000 ± 3000 | 2.53 ± 0.43 | |$1.54^{+0.02}_{-0.01}$| | T | Choi et al. (2012) |
16099−5651 | Hen 2-147 | 7000 ± 1500 | 2.96 ± 0.38 | 3.0 ±0.4 | PL | Santander-García et al. (2007) |
17317−3331 | V1018 Sco | 35 000 ± 5000 | 3.76 ± 0.66 | 3.20 ± 0.64 | L | Cohen, Parker & Chapman (2005) |
19158+0141 | Abell 58 (V605 Aql) | 4000 ± 1500 | 5.25 ± 1.21 | 4.60 ± 0.60 | E | Clayton et al. (2013) |
IRAS no. . | Common name . | L⋆ (L⊙) . | DSED (kpc) . | DLit (kpc) . | Method . | Reference . |
---|---|---|---|---|---|---|
(nascent)-PPNe / PAGB objects . | ||||||
09452+1330 | CW Leo (IRC+10216) | 9800 ± 3000 | 0.13 ± 0.03 | 0.123 ± 0.014 | G | Groenewegen et al. (2012) |
03507+1115 | IK Tau (NML Tau) | 9100 ± 3000 | 0.24 ± 0.06 | 0.265 ± 0.02 | E | Hale et al. (1997) |
21003+3630 | CRL 2688 | 6000 ± 1500 | 0.58 ± 0.08 | 0.42 ± 0.06 | E | Ueta, Murakawa & Meixner (2006) |
10158−2844 | HR 4049 | 3500 ± 1500 | 0.64 ± 0.15 | 0.64 ± 0.19 | D | Acke et al. (2013) |
18286−0959 | OH 21.80-0.13 | 6000 ± 1500 | 3.10 ± 0.45 | |$3.61^{+0.63}_{-0.47}$| | T | Imai et al. (2013a) |
18460−0151 | OH 31.0-0.2 | 6000 ± 1500 | 3.14 ± 0.80 | 2.1 ± 0.6 | S | Imai et al. (2013b) |
– | NGC 5139 WOR 1957 | 1700 ± 750 | 6.13 ± 1.36 | 5.2 ± 0.5 | M | Harris (1996) |
17423−1755 | Hen 3-1475 | 20 000 ± 5000 | 7.17 ± 1.02 | 8.30( ± 1.0) | E | Borkowski & Harrington (2001) |
– | NGC 6205 BARN 29 | 1700 ± 750 | 7.97 ± 1.90 | 7.1 ± 0.7 | M | Harris (1996) |
19134+2131 | [HSD93b] 16 | 6000 ± 1500 | 8.37 ± 1.21 | |$8.00^{+0.90}_{-0.70}$| | T | Imai, Sahai & Morris (2007) |
Young PNe | ||||||
– | NGC 7027 | 6000 ± 1500 | 0.95 ± 0.15 | 0.87 ± 0.10b | E | Zijlstra, van Hoof & Perley (2008, and references therein) |
19327+3024 | BD+30 3639 | 4000 ± 1500 | 1.25 ± 0.25 | 1.52 ± 0.21 | V | Akras & Steffen (2012) |
18240−0244 | M 2-43 | 4000 ± 1500 | 3.68 ± 0.75 | 6.9 ± 1.5 | E | Guzmán, Gómez & Rodríguez (2006) |
19219+0947 | Vy 2-2 | 6000 ± 1500 | 3.74 ± 0.60 | 3.6 ± 0.4 | E | Christianto & Seaquist (1998) |
19255+2123 | K 3-35 | 6000 ± 1500 | 5.77 ± 0.83 | |$3.90^{+0.70}_{-0.50}$| | T | Tafoya et al. (2011) |
19288−3419 | Hen 2-436 | 4000 ± 1500 | 26.2 ± 6.3 | 26.0 ± 2.0c | M | Zijlstra et al. (2006) |
Symbiotic stars/Other objects | ||||||
17028−1004 | M 2-9 | 2000 ± 750 | 1.41 ± 0.33 | 1.3 ±0.2 | G | Corradi et al. (2011) |
20097+2010 | Hen 1-5 (FG Sge) | 6000 ± 1500 | 2.08 ± 0.52 | 2.5 (±0.5) | P | Mayor & Acker (1980) |
07399−1435 | OH 231.8+4.2a | 15 000 ± 3000 | 2.53 ± 0.43 | |$1.54^{+0.02}_{-0.01}$| | T | Choi et al. (2012) |
16099−5651 | Hen 2-147 | 7000 ± 1500 | 2.96 ± 0.38 | 3.0 ±0.4 | PL | Santander-García et al. (2007) |
17317−3331 | V1018 Sco | 35 000 ± 5000 | 3.76 ± 0.66 | 3.20 ± 0.64 | L | Cohen, Parker & Chapman (2005) |
19158+0141 | Abell 58 (V605 Aql) | 4000 ± 1500 | 5.25 ± 1.21 | 4.60 ± 0.60 | E | Clayton et al. (2013) |
Notes. D: dynamical parallax; E: expansion parallax; G: geometric model; L: phase-lag method; M: membership of system at known distance; P: pulsation theory; PL: Mira P–L relationship; S: statistical parallax; T: trigonometric parallax; V: velocity field mapping. aPhysical member of the open cluster M 46; bthis is the average of the distances given in Zijlstra et al. (2008); crefer to the text.
IRAS no. . | Common name . | L⋆ (L⊙) . | DSED (kpc) . | DLit (kpc) . | Method . | Reference . |
---|---|---|---|---|---|---|
(nascent)-PPNe / PAGB objects . | ||||||
09452+1330 | CW Leo (IRC+10216) | 9800 ± 3000 | 0.13 ± 0.03 | 0.123 ± 0.014 | G | Groenewegen et al. (2012) |
03507+1115 | IK Tau (NML Tau) | 9100 ± 3000 | 0.24 ± 0.06 | 0.265 ± 0.02 | E | Hale et al. (1997) |
21003+3630 | CRL 2688 | 6000 ± 1500 | 0.58 ± 0.08 | 0.42 ± 0.06 | E | Ueta, Murakawa & Meixner (2006) |
10158−2844 | HR 4049 | 3500 ± 1500 | 0.64 ± 0.15 | 0.64 ± 0.19 | D | Acke et al. (2013) |
18286−0959 | OH 21.80-0.13 | 6000 ± 1500 | 3.10 ± 0.45 | |$3.61^{+0.63}_{-0.47}$| | T | Imai et al. (2013a) |
18460−0151 | OH 31.0-0.2 | 6000 ± 1500 | 3.14 ± 0.80 | 2.1 ± 0.6 | S | Imai et al. (2013b) |
– | NGC 5139 WOR 1957 | 1700 ± 750 | 6.13 ± 1.36 | 5.2 ± 0.5 | M | Harris (1996) |
17423−1755 | Hen 3-1475 | 20 000 ± 5000 | 7.17 ± 1.02 | 8.30( ± 1.0) | E | Borkowski & Harrington (2001) |
– | NGC 6205 BARN 29 | 1700 ± 750 | 7.97 ± 1.90 | 7.1 ± 0.7 | M | Harris (1996) |
19134+2131 | [HSD93b] 16 | 6000 ± 1500 | 8.37 ± 1.21 | |$8.00^{+0.90}_{-0.70}$| | T | Imai, Sahai & Morris (2007) |
Young PNe | ||||||
– | NGC 7027 | 6000 ± 1500 | 0.95 ± 0.15 | 0.87 ± 0.10b | E | Zijlstra, van Hoof & Perley (2008, and references therein) |
19327+3024 | BD+30 3639 | 4000 ± 1500 | 1.25 ± 0.25 | 1.52 ± 0.21 | V | Akras & Steffen (2012) |
18240−0244 | M 2-43 | 4000 ± 1500 | 3.68 ± 0.75 | 6.9 ± 1.5 | E | Guzmán, Gómez & Rodríguez (2006) |
19219+0947 | Vy 2-2 | 6000 ± 1500 | 3.74 ± 0.60 | 3.6 ± 0.4 | E | Christianto & Seaquist (1998) |
19255+2123 | K 3-35 | 6000 ± 1500 | 5.77 ± 0.83 | |$3.90^{+0.70}_{-0.50}$| | T | Tafoya et al. (2011) |
19288−3419 | Hen 2-436 | 4000 ± 1500 | 26.2 ± 6.3 | 26.0 ± 2.0c | M | Zijlstra et al. (2006) |
Symbiotic stars/Other objects | ||||||
17028−1004 | M 2-9 | 2000 ± 750 | 1.41 ± 0.33 | 1.3 ±0.2 | G | Corradi et al. (2011) |
20097+2010 | Hen 1-5 (FG Sge) | 6000 ± 1500 | 2.08 ± 0.52 | 2.5 (±0.5) | P | Mayor & Acker (1980) |
07399−1435 | OH 231.8+4.2a | 15 000 ± 3000 | 2.53 ± 0.43 | |$1.54^{+0.02}_{-0.01}$| | T | Choi et al. (2012) |
16099−5651 | Hen 2-147 | 7000 ± 1500 | 2.96 ± 0.38 | 3.0 ±0.4 | PL | Santander-García et al. (2007) |
17317−3331 | V1018 Sco | 35 000 ± 5000 | 3.76 ± 0.66 | 3.20 ± 0.64 | L | Cohen, Parker & Chapman (2005) |
19158+0141 | Abell 58 (V605 Aql) | 4000 ± 1500 | 5.25 ± 1.21 | 4.60 ± 0.60 | E | Clayton et al. (2013) |
IRAS no. . | Common name . | L⋆ (L⊙) . | DSED (kpc) . | DLit (kpc) . | Method . | Reference . |
---|---|---|---|---|---|---|
(nascent)-PPNe / PAGB objects . | ||||||
09452+1330 | CW Leo (IRC+10216) | 9800 ± 3000 | 0.13 ± 0.03 | 0.123 ± 0.014 | G | Groenewegen et al. (2012) |
03507+1115 | IK Tau (NML Tau) | 9100 ± 3000 | 0.24 ± 0.06 | 0.265 ± 0.02 | E | Hale et al. (1997) |
21003+3630 | CRL 2688 | 6000 ± 1500 | 0.58 ± 0.08 | 0.42 ± 0.06 | E | Ueta, Murakawa & Meixner (2006) |
10158−2844 | HR 4049 | 3500 ± 1500 | 0.64 ± 0.15 | 0.64 ± 0.19 | D | Acke et al. (2013) |
18286−0959 | OH 21.80-0.13 | 6000 ± 1500 | 3.10 ± 0.45 | |$3.61^{+0.63}_{-0.47}$| | T | Imai et al. (2013a) |
18460−0151 | OH 31.0-0.2 | 6000 ± 1500 | 3.14 ± 0.80 | 2.1 ± 0.6 | S | Imai et al. (2013b) |
– | NGC 5139 WOR 1957 | 1700 ± 750 | 6.13 ± 1.36 | 5.2 ± 0.5 | M | Harris (1996) |
17423−1755 | Hen 3-1475 | 20 000 ± 5000 | 7.17 ± 1.02 | 8.30( ± 1.0) | E | Borkowski & Harrington (2001) |
– | NGC 6205 BARN 29 | 1700 ± 750 | 7.97 ± 1.90 | 7.1 ± 0.7 | M | Harris (1996) |
19134+2131 | [HSD93b] 16 | 6000 ± 1500 | 8.37 ± 1.21 | |$8.00^{+0.90}_{-0.70}$| | T | Imai, Sahai & Morris (2007) |
Young PNe | ||||||
– | NGC 7027 | 6000 ± 1500 | 0.95 ± 0.15 | 0.87 ± 0.10b | E | Zijlstra, van Hoof & Perley (2008, and references therein) |
19327+3024 | BD+30 3639 | 4000 ± 1500 | 1.25 ± 0.25 | 1.52 ± 0.21 | V | Akras & Steffen (2012) |
18240−0244 | M 2-43 | 4000 ± 1500 | 3.68 ± 0.75 | 6.9 ± 1.5 | E | Guzmán, Gómez & Rodríguez (2006) |
19219+0947 | Vy 2-2 | 6000 ± 1500 | 3.74 ± 0.60 | 3.6 ± 0.4 | E | Christianto & Seaquist (1998) |
19255+2123 | K 3-35 | 6000 ± 1500 | 5.77 ± 0.83 | |$3.90^{+0.70}_{-0.50}$| | T | Tafoya et al. (2011) |
19288−3419 | Hen 2-436 | 4000 ± 1500 | 26.2 ± 6.3 | 26.0 ± 2.0c | M | Zijlstra et al. (2006) |
Symbiotic stars/Other objects | ||||||
17028−1004 | M 2-9 | 2000 ± 750 | 1.41 ± 0.33 | 1.3 ±0.2 | G | Corradi et al. (2011) |
20097+2010 | Hen 1-5 (FG Sge) | 6000 ± 1500 | 2.08 ± 0.52 | 2.5 (±0.5) | P | Mayor & Acker (1980) |
07399−1435 | OH 231.8+4.2a | 15 000 ± 3000 | 2.53 ± 0.43 | |$1.54^{+0.02}_{-0.01}$| | T | Choi et al. (2012) |
16099−5651 | Hen 2-147 | 7000 ± 1500 | 2.96 ± 0.38 | 3.0 ±0.4 | PL | Santander-García et al. (2007) |
17317−3331 | V1018 Sco | 35 000 ± 5000 | 3.76 ± 0.66 | 3.20 ± 0.64 | L | Cohen, Parker & Chapman (2005) |
19158+0141 | Abell 58 (V605 Aql) | 4000 ± 1500 | 5.25 ± 1.21 | 4.60 ± 0.60 | E | Clayton et al. (2013) |
Notes. D: dynamical parallax; E: expansion parallax; G: geometric model; L: phase-lag method; M: membership of system at known distance; P: pulsation theory; PL: Mira P–L relationship; S: statistical parallax; T: trigonometric parallax; V: velocity field mapping. aPhysical member of the open cluster M 46; bthis is the average of the distances given in Zijlstra et al. (2008); crefer to the text.
IRAS no. . | DSED (kpc) . | DSahai (kpc) . | |ΔD (kpc)| . | Morph . |
---|---|---|---|---|
01037+1219 | – | 0.65 | – | E |
11385−5517 | 1.05 | 1.20 | 0.15 | I |
13428−6232 | 1.91 | 1.90 | 0.01 | B |
13557−6442 | 3.14 | 3.30 | 0.16 | B |
15045−4945 | – | 4.90 | – | B |
15452−5459 | 1.77 | 1.70 | 0.07 | B |
15553−5230 | 4.49 | 4.30 | 0.19 | B |
16559−2957 | 5.48 | 5.20 | 0.28 | B |
17253−2831 | 7.39 | 7.70 | 0.31 | E |
17347−3139 | 3.68 | 3.10 | 0.58 | B |
17440−3310 | 7.19 | 7.60 | 0.41 | B |
17543−3102 | 7.99 | 7.30 | 0.69 | B |
18276−1431 | 2.94 | 3.00 | 0.06 | B |
18420−0512 | 6.87 | 6.60 | 0.27 | E |
19024+0044 | 5.99 | 5.10 | 0.89 | M |
19134+2131 | 8.35 | 8.40 | 0.05 | B |
19292+1806 | 5.48 | 5.20 | 0.28 | B |
19306+1407 | 5.99 | 4.90 | 1.09 | B |
19475+3119 | 1.54 | 4.20 | 2.66 | M |
20000+3239 | 3.50 | 3.50 | 0.00 | E |
22036+5306a | 3.83 | – | – | B |
22223+4327 | 4.29 | 4.30 | 0.01 | E |
23304+6147 | 4.26 | 4.20 | 0.06 | M |
IRAS no. . | DSED (kpc) . | DSahai (kpc) . | |ΔD (kpc)| . | Morph . |
---|---|---|---|---|
01037+1219 | – | 0.65 | – | E |
11385−5517 | 1.05 | 1.20 | 0.15 | I |
13428−6232 | 1.91 | 1.90 | 0.01 | B |
13557−6442 | 3.14 | 3.30 | 0.16 | B |
15045−4945 | – | 4.90 | – | B |
15452−5459 | 1.77 | 1.70 | 0.07 | B |
15553−5230 | 4.49 | 4.30 | 0.19 | B |
16559−2957 | 5.48 | 5.20 | 0.28 | B |
17253−2831 | 7.39 | 7.70 | 0.31 | E |
17347−3139 | 3.68 | 3.10 | 0.58 | B |
17440−3310 | 7.19 | 7.60 | 0.41 | B |
17543−3102 | 7.99 | 7.30 | 0.69 | B |
18276−1431 | 2.94 | 3.00 | 0.06 | B |
18420−0512 | 6.87 | 6.60 | 0.27 | E |
19024+0044 | 5.99 | 5.10 | 0.89 | M |
19134+2131 | 8.35 | 8.40 | 0.05 | B |
19292+1806 | 5.48 | 5.20 | 0.28 | B |
19306+1407 | 5.99 | 4.90 | 1.09 | B |
19475+3119 | 1.54 | 4.20 | 2.66 | M |
20000+3239 | 3.50 | 3.50 | 0.00 | E |
22036+5306a | 3.83 | – | – | B |
22223+4327 | 4.29 | 4.30 | 0.01 | E |
23304+6147 | 4.26 | 4.20 | 0.06 | M |
Notes.aThe Sahai et al. (2007) distance for IRAS 22036+5306 has been excluded since it was not derived via the SED technique.
IRAS no. . | DSED (kpc) . | DSahai (kpc) . | |ΔD (kpc)| . | Morph . |
---|---|---|---|---|
01037+1219 | – | 0.65 | – | E |
11385−5517 | 1.05 | 1.20 | 0.15 | I |
13428−6232 | 1.91 | 1.90 | 0.01 | B |
13557−6442 | 3.14 | 3.30 | 0.16 | B |
15045−4945 | – | 4.90 | – | B |
15452−5459 | 1.77 | 1.70 | 0.07 | B |
15553−5230 | 4.49 | 4.30 | 0.19 | B |
16559−2957 | 5.48 | 5.20 | 0.28 | B |
17253−2831 | 7.39 | 7.70 | 0.31 | E |
17347−3139 | 3.68 | 3.10 | 0.58 | B |
17440−3310 | 7.19 | 7.60 | 0.41 | B |
17543−3102 | 7.99 | 7.30 | 0.69 | B |
18276−1431 | 2.94 | 3.00 | 0.06 | B |
18420−0512 | 6.87 | 6.60 | 0.27 | E |
19024+0044 | 5.99 | 5.10 | 0.89 | M |
19134+2131 | 8.35 | 8.40 | 0.05 | B |
19292+1806 | 5.48 | 5.20 | 0.28 | B |
19306+1407 | 5.99 | 4.90 | 1.09 | B |
19475+3119 | 1.54 | 4.20 | 2.66 | M |
20000+3239 | 3.50 | 3.50 | 0.00 | E |
22036+5306a | 3.83 | – | – | B |
22223+4327 | 4.29 | 4.30 | 0.01 | E |
23304+6147 | 4.26 | 4.20 | 0.06 | M |
IRAS no. . | DSED (kpc) . | DSahai (kpc) . | |ΔD (kpc)| . | Morph . |
---|---|---|---|---|
01037+1219 | – | 0.65 | – | E |
11385−5517 | 1.05 | 1.20 | 0.15 | I |
13428−6232 | 1.91 | 1.90 | 0.01 | B |
13557−6442 | 3.14 | 3.30 | 0.16 | B |
15045−4945 | – | 4.90 | – | B |
15452−5459 | 1.77 | 1.70 | 0.07 | B |
15553−5230 | 4.49 | 4.30 | 0.19 | B |
16559−2957 | 5.48 | 5.20 | 0.28 | B |
17253−2831 | 7.39 | 7.70 | 0.31 | E |
17347−3139 | 3.68 | 3.10 | 0.58 | B |
17440−3310 | 7.19 | 7.60 | 0.41 | B |
17543−3102 | 7.99 | 7.30 | 0.69 | B |
18276−1431 | 2.94 | 3.00 | 0.06 | B |
18420−0512 | 6.87 | 6.60 | 0.27 | E |
19024+0044 | 5.99 | 5.10 | 0.89 | M |
19134+2131 | 8.35 | 8.40 | 0.05 | B |
19292+1806 | 5.48 | 5.20 | 0.28 | B |
19306+1407 | 5.99 | 4.90 | 1.09 | B |
19475+3119 | 1.54 | 4.20 | 2.66 | M |
20000+3239 | 3.50 | 3.50 | 0.00 | E |
22036+5306a | 3.83 | – | – | B |
22223+4327 | 4.29 | 4.30 | 0.01 | E |
23304+6147 | 4.26 | 4.20 | 0.06 | M |
Notes.aThe Sahai et al. (2007) distance for IRAS 22036+5306 has been excluded since it was not derived via the SED technique.
In Fig. 4 , we plot the SED derived distances against the independent literature distances given in Table 7. A least-squares fit of the data with a slope of 0.99 ± 0.06 indicates that the two data sets are in very good agreement. Fig. 4 demonstrates clearly that the SED technique is a viable method for determining distances to dusty PAGB objects. In addition to the independent distance comparison, we have compared our SED distances (with a default luminosity, L⋆ = 6000 L⊙) with those derived by Sahai et al. (2007), previously the largest sample of SED derived distances for PAGB objects. In the bottom panel of Fig. 4, we show that our independent SED distances are in excellent agreement with those of Sahai et al. (2007). Note that we have excluded the distance of IRAS 22036+5306 from this comparison on the grounds that the Sahai et al. (2007) distance was not derived using the SED technique. The only outlying object is IRAS 19475+3119 where the large reddening [E(B − V) ∼ 0.8 mag] coupled with the strong flux contribution from the central star significantly alters the total isotropic flux thereby giving a distance vastly different to the reddened value of Sahai et al. (2007).
5.3 Notes on individual objects with independent distances
5.3.1 IRAS 03507+1115 (NML Tau)
This nearby dust-enshrouded Mira (Decin et al. 2010), or nascent PPN, is historically known as NML Tau (Neugebauer, Martz & Leighton 1965). The 469 d pulsation period leads to a bolometric luminosity of 9100 ± 3000 L⊙ from the P–L relation of Feast et al. (1989), and a distance of 240 ± 60 pc, in good agreement with the expansion distance from Hale et al. (1997) of 265 ± 20 pc.
5.3.2 IRAS 06176−1036 (HD 44179)
This is the Red Rectangle, quite possibly the most famous PPN for its beautiful reflection nebulosity (Cohen et al. 2004; Van Winckel 2014). We required three blackbody curves to represent the dust component, ranging from T ∼ 250 to 1400 K, to represent the Keplerian disc in this system (Men'shchikov et al. 2002). The integrated flux of the Red Rectangle gives a distance of DSED = 0.85 ± 0.19 kpc compared to the currently accepted distance of 0.73 kpc derived by Men'shchikov et al. (2002). These authors also derived a core mass of 0.58 M⊙ which corresponds to a luminosity of ≈4000 L⊙ using our prescription, whereas those authors used a higher luminosity of 6050 L⊙. Using a luminosity of 4000 L⊙, we derive a distance to the Red Rectangle of DSED = 0.69 ± 0.18 kpc, in excellent agreement with their distance.
5.3.3 IRAS 09452+1330 (IRC+10216)
CW Leo is the nearest dust-enshrouded AGB star to the Sun (Becklin et al. 1969; Morris 1975; Menten et al. 2012). The surrounding bipolar nebula (Christou et al. 1991; Kastner & Weintraub 1994; Osterbart et al. 2000) indicates that the star is very close to leaving the AGB and becoming a bona fide PPN. We use the distance from Groenewegen et al. (2012) as a comparison with our own distance. By modelling the time lag between the star and surrounding bow-shock, these authors determined a distance of 123 ± 14 pc. We refine our SED distance by estimating the bolometric luminosity directly from the 630 d pulsation period and adopting the revised P–L relation for carbon-rich Miras from Groenewegen & Whitelock (1996). We determine L = 9800 L⊙, and a distance DSED = 130 ± 30 pc, in excellent agreement.
5.3.4 IRAS 10158−2844 (HR 4049)
AG Ant is a long-period spectroscopic binary with a core mass of 0.58 M⊙ and a secondary mass of 0.34 M⊙ (Hinkle, Brittain & Lambert 2007). The large IR excess of HR 4049 has been modelled using a blackbody of T ∼ 1100 K by Dominik et al. (2003), De Ruyter et al. (2006) and Hinkle et al. (2007) in agreement with our SED modelling. Several authors have determined the stellar parameters of HR 4049 (Lambert, Hinkle & Luck 1988; Waelkens et al. 1991) finding an effective temperature Teff ≈ 7500 K higher than the photospheric temperature determined via blackbody fitting. Our SED distance of 0.64 kpc is in perfect agreement with a new dynamical parallax distance of 0.64 kpc from Acke et al. (2013).
5.3.5 IRAS 17028−1004 (M2-9)
An extensive literature (e.g. Allen & Swings 1972; Schwarz et al. 1997; Smith & Gehrz 2005, and references therein) exists on this beautiful and remarkable bipolar nebula. It is considered here to be a resolved symbiotic outflow, based on its spectral characteristics and red NIR colours (see Schmeja & Kimeswenger 2001). Corradi et al. (2011) have modelled the light changes of the Butterfly Nebula to determine a distance of 1.3 ± 0.2 kpc. The physical properties of this system are then consistent with a symbiotic-like interacting binary as the central engine, but the implied luminosity (2500 L⊙) is less than expected for a thermally-pulsing AGB (TP-AGB) star. It may be a red giant branch or early-AGB star, or alternatively a thick-disc TP-AGB star.
5.3.6 IRAS 17317−3331 (V1018 Sco)
This heavily obscured, high-mass, oxygen-rich OH/IR star is surrounded by an ionized ring nebula (Cohen et al. 2005, 2006), which is either an unusual PN or a symbiotic outflow. Our SED distance compares well with the phase-lag distance of 3.2 ± 0.6 kpc from Cohen et al. (2005). Assuming that the pulsation period of the star is given by the maser light curve, the 1486 d period (Chapman, Habing & Killeen 1995) leads to a bolometric luminosity of 35 000 L⊙ adopting the Mira P–L relation for O-rich stars of Feast et al. (1989). Our measured bolometric flux, Ftot = 2475 L⊙ kpc−2, leads directly to a distance of 3.76 ± 0.66 kpc, in agreement with the phase-lag distance within the uncertainties.
5.3.7 IRAS 17423−1755 (Hen 3-1475)
This is a particularly interesting example of an oxygen-rich PAGB star with highly collimated outflows (Riera et al. 1995; Manteiga et al. 2011). Assuming L = 20 000 L⊙, appropriate for a high-mass progenitor, we estimate a distance of 7.2 kpc, in reasonable agreement with the expansion parallax distance of Borkowski & Harrington (2001). Similar objects to Hen 3-1475 are IRAS 19343+2926 and IRAS 17393−2727,13 which shows a strong 12.8 μm [Ne ii] nebular emission line, showing that its central region is already ionized by its fast evolving central star (García-Hernández et al. 2007).
5.3.8 IRAS 19288−3419 (Hen 2-436)
Hen 2-436 is the optically brightest PN in the Sagittarius dwarf spheroidal galaxy (see Zijlstra et al. 2006, and references therein). Though not in the Toruń catalogue, this dusty object belongs to the rare group of IR-[WC] PNe (Zijlstra 2001), so is amenable to the SED technique. Recent distances to the Sgr dSph galaxy range from 24.3 kpc (McDonald et al. 2013) to 28–30 kpc (Siegel et al. 2011), suggesting that there is considerable line-of-sight depth. We adopt a distance of 26 ± 2 kpc which is the average of the values tabulated by Kunder & Chaboyer (2009) and McDonald et al. (2013). This is in very good agreement with our SED distance to Hen 2-436 of 26.2 kpc, after assuming a stellar luminosity of 4000 L⊙ for its [WC] nucleus (see Section 3.1).
5.3.9 IRAS 21003+3630 (CRL 2688)
The Cygnus Egg is the closest PPN to the Sun, with a distance of 0.42 kpc, derived from two-epoch Hubble Space Telescope (HST) proper motions of the expanding system of nebular features (Ueta et al. 2006). Our bolometric flux, Ftot = 17 850 L⊙ kpc−2, combined with an assumed luminosity of 6000 L⊙ leads to an estimated distance of 0.58 ± 0.08 kpc. It appears that the intrinsic luminosity must be less than our assumed value, despite CRL 2668 having a DUPLEX morphology.
5.3.10 OH 231.8+04.2, a bipolar outflow in an open cluster
There is an extensive literature on this object so only a brief sketch is provided here. Commonly known as the Calabash or Rotten Egg Nebula, this strongly bipolar reflection/emission nebulosity lies in the outskirts of the rich open cluster, M46 (NGC 2437), and is generally regarded as a PPN (Sánchez Contreras et al. 2000; Bujarrabal et al. 2002; Meakin et al. 2003). However, the fact that there is a binary system at the centre, comprising a very late type (M9–10 III) Mira star (Cohen 1981; Kastner et al. 1998) with a 700 d pulsation period and an A0 V companion (Sańchez Contreras, Gil de Paz & Sahai 2004), seems to indicate that the nebula has more in common with D-type symbiotic outflows (Corradi 1995), than with true PPNe.
Unlike the case of the planetary, NGC 2438 which is projected on M46 (Kiss et al. 2008), the systemic radial velocity of the Calabash nebula (Jura & Morris 1985; Morris et al. 1987; Sánchez Contreras, Bujarrabal & Alcolea 1997; Zijlstra et al. 2001; Bujarrabal et al. 2012) is consistent with the mean radial velocity of M46 (Mermilliod et al. 2007; Frinchaboy & Majewski 2008; Kiss et al. 2008). Both the phase-lag distance to OH 231.8 of 1.3–1.5 kpc (Bowers & Morris 1984; Kastner et al. 1992) and the much more accurate maser trigonometric distance of 1.54 kpc from Choi et al. (2012) are in good agreement with recent distance determinations to M46 from Sharma et al. (2006) and Davidge (2013). Lastly, the proper motions of its associated water masers (Choi et al. 2012) also agree with the cluster's proper motion (Dias, Lépine & Alessi 2002). As the membership of OH 231.8+04.2 in M46 is now secure, we can use the age of the cluster (2.2–2.5 × 108 yr; Sharma et al. 2006; Davidge 2013) to infer the progenitor mass of the Mira, which turns out to be ∼3.4 M⊙, in agreement with previous mass estimates (Kastner et al. 1998). The substantial nitrogen enrichment compared to the solar abundance (Cohen et al. 1985) seen in the shock-ionized lobes of the Calabash (Reipurth 1987; Bujarrabal et al. 2002) is consistent with mild HBB in this star. Hence, this observation may indicate that nitrogen enrichment can occur at stellar masses under ∼4.0 M⊙ (see Karakas et al. 2009).
6 SUMMARY AND FUTURE WORK
The compilation of the Toruń catalogue (Szczerba et al. 2007, 2012) gathered a wide assortment of flux data for all Galactic PAGB objects known at that time. We have used these data, adding more recent fluxes from the literature to build a homogenized set of distances by modelling their observed SEDs with one or more blackbody curves. Total fluxes were calculated for each object by numerically integrating the fitted curves. This approach was adopted both for simplicity, given the large amount of data used, and internal consistency, as spectral types, needed for more detailed fitting, were unknown in the majority of cases. In a follow-up paper, we will investigate any potential discrepancy in the integrated fluxes when using model atmospheres as opposed to blackbody fitting. We expect that for severely reddened stars and stars with a high-FIR/F⋆ model atmospheres will make little difference compared to the UV-bright stars where we expect a larger difference in integrated fluxes. The assumed luminosity (in solar units) was derived using the empirical core-mass luminosity relation for PAGB evolution from VW94, and a set of criteria separating the different populations of PAGB objects. Distances were computed by equating an assumed luminosity with the integrated flux of each source thus creating a homogenized set of distances, presented in full as an online supplement. The calculated distances were compared to several independently derived literature values measured using a variety of methods, in order to ascertain the accuracy of our approach. In Section 5.2, we showed that our derived distances are in good agreement with a range of literature values. In this way, we have effectively demonstrated that the SED technique is a valid method for calculating statistical distances to PAGB and related objects. In a follow-up paper (Vickers et al., in preparation), we will determine distances to the remaining objects in the Toruń catalogue, namely the RV Tauri stars, using instead empirical P–L relations, as well as the R CrB stars (and related hydrogen-deficient objects).
In a further paper, we will investigate the population characteristics of a relatively complete volume-limited sample of Galactic disc PAGB objects for the first time. Such a census can be used for understanding the population demographics of Galactic PAGB objects, and their relationships with their precursor AGB stars and descendent PNe (Frew & Parker 2006; Frew 2008; Frew et al. 2014b). Volume-limited samples are an underappreciated tool for studying stellar populations (Frew & Parker 2012), having the power to unlock the vital characteristics of Galactic PAGB objects, needed to understand the possible shaping mechanisms of their progenitors (Balick & Frank 2002). Specifically, we will endeavour to determine the scaleheights (and hence progenitor ages and masses) of the various subgroups of PAGB stars and relate these to the morphological and dust properties of the resolved nebulae (Ueta et al. 2000; Siódmiak et al. 2008), and those objects that possess Keplerian dust discs (see Van Winckel et al. 2006; Hinkle et al. 2007; van Aarle et al. 2011; Acke et al. 2013). Our upcoming analysis will be undertaken with much larger samples than have been utilized previously (Likkel, te Lintel Hekkert & Chapman 1993).
Finally, we expect the data avalanche from modern multiwavelength surveys to aid in the discovery of many more PAGB stars and PPNe in the Galaxy. These will be incorporated into our new relational data base of PNe and PAGB stars currently under construction at Macquarie University, in conjunction with the CDS, Strasbourg (Bojičić et al., in preparation).
We thank the referee, Ryszard Szczerba, for constructive comments that improved this paper. This research has made use of the SIMBAD data base and the VizieR service, operated at CDS, Strasbourg, France. Additional data were obtained from the Mikulski Archive for Space Telescopes (MAST). STScI is operated by the Association of Universities for Research in Astronomy, Inc. Support for MAST for non-HST data is provided by the NASA Office of Space Science. DJF thanks Macquarie University for an MQ Research Fellowship and ISB is the recipient of an Australian Research Council Super Science Fellowship (project ID FS100100019), while QAP acknowledges additional support from the Australian Astronomical Observatory.
This research made use of data products from the Midcourse Space Experiment. Processing of the data was funded by the Ballistic Missile Defense Organization with additional support from NASA Office of Space Science. This research has also made use of the NASA/IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.
The less-luminous type II Cepheids (W Vir and BL Her stars) are probably in an intermediate evolutionary phase between the blue horizontal branch and the base of the AGB, or are on a blue loop from the lower AGB (e.g. Maas, Giridhar & Lambert 2007). They are not considered further.
Vizier is hosted by the Centre de Données astronomiques de Strasbourg (CDS). See http://vizier.u-strasbg.fr/viz-bin/VizieR.
Planck (http://www.esa.int/Planck) is a project of the European Space Agency (ESA) with instruments provided by two scientific consortia funded by ESA member states, with contributions from NASA (USA) and telescope reflectors provided by a collaboration between ESA and a scientific consortium led and funded by Denmark.
This object is not in the current edition of the Toruń catalogue.
REFERENCES
SUPPORTING INFORMATION
Additional Supporting Information may be found in the online version of this article:
Table 4. Parameters and distances for 209 likely PAGB stars in the Torun catalogue, ordered by Galactic longitude.
Table 5. Parameters and distances for 87 possible PAGB stars in the Torun catalogue, ordered by Galactic longitude.
Table 6. Parameters and distances for 54 miscellaneous objects not in the Torun Catalogue, ordered by Galactic longitude.
Appendix B. Spectral energy distributions for 209 likely PAGB stars.
Appendix C. Spectral energy distributions for 87 possible PAGB stars.
Appendix D. Spectral energy distributions for 54 miscellaneous objects.
Please note: OUP is not responsible for the content or functionality of any supporting materials supplied by the authors. Any queries (other than missing material) should be directed to the corresponding author for the paper.