of all PNs, D11% are bipolar. The other nonspherical PNs are elliptical, and can be further classiðed into subgroups

Transcription

1 THE ASTROPHYSICAL JOURNAL, 496:833È841, 1998 April 1 ( The American Astronomical Society. All rights reserved. Printed in U.S.A. BINARY PROGENITOR MODELS FOR BIPOLAR PLANETARY NEBULAE NOAM SOKER Department of Physics, University of Haifa at Oranim, Oranim, Tivon 36006, Israel; soker=physics.technion.ac.il Received 1997 May 30; accepted 1997 November 7 ABSTRACT We propose an explanation for the positive correlation of bipolar planetary nebulae with massive progenitors in the paradigm of binary system progenitors. We list 10 critical observations, and argue that single-star models for the formation of bipolar planetary nebulae encounter difficulties complying with these observations. On the other hand, binary system progenitors can naturally explain these key observations, and in addition explain the rich varieties of structures possessed by bipolar planetary nebulae. Based on three of the critical observations, and on previous works by Corradi and Schwarz and by Morris, we postulate that the progenitors of bipolar planetary nebulae are binary stellar systems in which the secondary diverts a substantial fraction of the mass lost by the asymptotic giant branch (AGB) primary, but the systems avoid the common envelope phase for a large fraction of the interaction time. The positive correlation of bipolar planetary nebulae with massive progenitors, M Z 2 M, is attributed to the larger ratios of red giant branch (RGB) to AGB radii which low-mass stars attain, compared with massive stars. These larger radii on the RGB cause most stellar binary companions, which potentially could have formed bipolar planetary nebulae if the primary had been on the AGB, to interact with lowmass primaries already on the RGB. This scenario predicts that the central stars of most bipolar planetary nebulae are in binary systems having orbital periods in the range of a few days to few times 10 yr. Subject headings: ISM: general È planetary nebulae: general È stars: AGB and post-agb È stars: binaries: close È stars: mass loss 1. INTRODUCTION primary is a red giant or an asymptotic giant branch (AGB) star and the companion is a main-sequence star or a white Most planetary nebulae (PNs) have a large-scale axisymmetric structure rather than a spherical one. The intrinsically (i.e., not accounting for the inñuence of the interstellar medium [ISM]) spherical PNs amount to only D10% of all PNs (Soker 1997). Corradi & Schwarz (1995) estimate that of all PNs, D11% are bipolar. The other nonspherical PNs are elliptical, and can be further classiðed into subgroups (e.g., Stanghellini, Corradi, & Schwarz 1993; Stanghellini dwarf (WD). Morris (1987, 1990) suggested a binary model for the formation of bipolar PNs. In his scenario a binary companion accretes from the primary. As a result, part of the mass is lost in the equatorial plane and part is collimated to form highly collimated high-velocity Ñows within the bipolar lobes. In the present paper we summarize and present further supporting arguments for the binary progenitor nature of 1995). The bipolar PNs (also called bilobal ÏÏ and bipolar PNs, and present a more coherent scenario. In particular, butterñy ÏÏ) are deðned according to Schwarz, Corradi, & Stanghellini (1992) as axially symmetric PNs having two lobes with an equatorial ÏÏ waist between them. The bipolar PNs are distinguished from the elliptical PNs by several properties. However, for only two properties (besides the morphology that serves to deðne the groups) we try to explain why only binary systems with massive primaries form bipolar PNs. We do not discuss the relation of this scenario to other binary systems, apart from some aspects of symbiotic nebulae. A thorough study of evolution and formation of close binary central stars of PNs and related binary systems is given by Iben & Tutukov there is a clear di erence between the two groups. These are (1993) and Han, Podsiadlowski, & Eggleton (1995). Stathe maximum expansion velocities, which are much higher tistical analysis is given also by de Kool (1992) and by for bipolar PNs, and the distribution in the Galaxy. The bipolar PNs are concentrated toward the Galactic plane (Greig 1972), with an average Galactic height of o z o ^ 130 pc, versus an average of o z o ^ 325 pc for elliptical PNs (Corradi & Schwarz 1995). This results from the fact that the progenitors of the bipolar PNs are massive mainsequence Yungelson, Tutukov, & Livio (1993). We also list several processes in binary systems which can lead to a higher mass-loss rate in the equatorial plane. Recent reviews on some aspects regarding formation of axisymmetric PNs from stellar binary systems are given by Livio (1997) and Yungelson (1997), while a review that includes substellar stars (M Z 2 M ; B or A stars). The fact that (including planetary systems) as well as stellar companions bipolar PNs result from massive progenitors has been is given by Soker (1997). Soker (1997) presents nine key widely used to argue that properties of the progenitors determine the Ðnal morphology of the descendant PN, and hence no binary companion is required to account for the observations ÏÏ that should be met by any model for the formation of axisymmetric PNs. Of these, the main problem for the binary model was the observation that the progenitors bipolar morphology (e.g., Pottasch 1995). On the other of bipolar PNs are massive. Here we try to show that hand, Morris (1990), and Corradi and Schwarz in a series of the binary scenario for the formation of axisymmetric PNs papers (Corradi 1995; Corradi & Schwarz 1993, 1995; can account for this correlation as well. Section 2 reviews Schwarz & Corradi 1992, 1995, and further references in these papers) argue very convincingly that bipolar PNs are similar in many ways to symbiotic nebulae that are known to result from binary interaction. In symbiotic systems the the main properties of bipolar PNs that any model should account for. In 3 we list several processes that can inñuence the structures of the descendant PNs. This demonstrates the rich types of structures that the descendant PNs 833

2 834 SOKER Vol. 496 can acquire. In 4 we give the main assumptions of the scenario, which are in keeping with the studies of Morris (1981, 1987, 1990) and those of Schwarz and Corradi, mentioned above. We postulate that the progenitors of bipolar PNs are binary stellar systems which avoid the common envelope phase during a substantial fraction of the interaction process. We show that in order to comply with the basic assumption of the scenario the primary is very likely to be of M Z 2 M on the main sequence. Our summary is in CRITICAL OBSERVATIONS In this section we list the main observations that any model for bipolar PNs should explain and comply with. 1. Bipolar PNs versus PNs with binary central stars.è The list of PNs found to have close binary nuclei (Bond & Livio 1990; Bond 1995) and the list of 43 bipolar PNs of Corradi & Schwarz (1995; also in Schwarz & Corradi 1995) have only one object in common, NGC 2346 (see item 2 below). Corradi & Schwarz (1995) classify 13 PNs in the list of Bond & Livio (1990): one bipolar, one possible bipolar, eight elliptical PNs, and three irregular PNs. Further support for the di erent population of bipolar PNs and PNs with central close binary stars comes from the distributions in Galactic latitude b and Galactic height z of PNs with central binary stars. This distribution resembles that of elliptical rather than that of bipolar PNs. Of the 15 PNs containing binary central stars, 13 are in the catalog of Acker et al. (1992). For LoTr 1, I could not Ðnd reliable data, while the nebula around the binary central star Be UMa is at a distance of 1.1 kpc (Wood, Robinson & Zhang 1995). Of the 14 PNs with satisfactory data, seven have Galactic latitude b [ 10.0, and seven have b \ 10. For 12 of these I could Ðnd the distances z from the Galactic plane, six having z Z 400 pc and six having z [ 130 pc. Although this distribution cannot be Ðtted with the usual distribution of a Ðxed scale height, it does show that a substantial fraction of these PNs comes from low-mass progenitors, as most elliptical PNs do.we conclude, like several previous works (e.g., Corradi & Schwarz 1995; Walsh & Walton 1996), that PNs that are formed from binary systems that evolve through a common envelope phase and end with small orbital separations have a di erent morphological type from the bipolar PNs studied by Corradi & Schwarz (1995). The PNs with binary central stars, though, are likely to have a higher density contrast between equatorial plane and polar regions than most of the other elliptical PNs (Pollacco & Bell 1997). 2. NGC 2346.ÈThis is the only PN which is deðned as bipolar by Corradi & Schwarz (1995) and which is known to have a central binary star, V651 Monocerotis. A possible evolutionary root for the formation of this binary central star is discussed by Iben & Livio (1993; see also Iben 1995) based on a study by Iben & Tutukov (1993). They claim that the primary, of mass M \ 0.4 ^ 0.05 M, is a degenerate helium core, while the 1 secondary, of mass M \ 1.8 ^ 0.3 M, is a main-sequence star. The orbital period 2 is days, and the orbital semimajor axis is 34.9 R (Iben & Livio 1993). Iben & Livio suggest that the primary over- Ñowed its Roche lobe while on the late red giant branch (RGB; a late case B), and the system entered a common envelope phase. The initial orbital semimajor axis was D1 AU. During the common envelope phase the binary system spirals in to the Ðnal separation of 34.9 R. Three properties of this system are most important for the present study: (a) The mass of the central hot star is much lower than the mass expected for single-star evolution of PNs. There are more bipolar PNs with low-mass central stars (item 6 below). (b) The orbital separation is the largest of the 15 known binary central stars of PNs (Bond & Livio 1990). (c) The secondary is a relatively massive star, i.e., more massive than the lower limit of the progenitor of bipolar PNs (item 8 below). 3. Symbiotic nebulae.èas was pointed out by Morris (1990) and in several papers by Corradi and Schwarz (e.g., Corradi 1995; Corradi & Schwarz 1995; Schwarz & Corradi 1995), the morphologies of many bipolar PNs, particularly those with high expansion velocities v (e.g., He 2-111), are similar to that of many symbiotic e nebulae. Several other systems seem to make the link between the model of symbiotic systems and bipolar PNs. Kahane et al. (1996) discuss the interesting giant carbon star V Hydrae. In 5.3 of their paper they mention that this system, which seems to be on its way to becoming a bipolar PN, is similar in several respects to the symbiotic binary system R Aqr (see item 9 below). V Hya has a high rotational velocity (Barnbaum, Morris, & Kahane 1995), which was attributed to a recent common envelope phase by Barnbaum et al. Because of the similarity of V Hya to symbiotic nebulae in several properties, we propose that a substantial fraction of the circumstellar envelope was ejected while the secondary was still outside the envelope. The Red Rectangle nebula and its binary central star, HD 44179, may be on their way to becoming a bipolar PN. This system was studied in detail by Waelkens et al. (1996). They claim that the primaryïs mass is in the range 0.56È0.80 M, while the secondaryïs mass is in the range 0.77È0.91 M. The orbital period is 300 days. Waelkens et al. (1996) suggest that there is a thick extended circumstellar disk around the central star. A similar model, they suggest, applies to HR 4049, a binary system with an orbital period of 429 days. We suggest that these three systems are typical progenitors of bipolar PNs. In some AGB stars with bipolar circumstellar envelopes there are still no hints of binary companions (e.g., X Herculis: Kahane & Jura 1996), although it is tempting to consider binary companions even in these systems, as Kahane & Jura (1996) do for X Her. 4. Massive progenitors.èas noted by many surveys (e.g., Zuckerman & Gatley 1988), bipolar PNs are concentrated toward the Galactic plane. This and composition di erences from elliptical PNs strongly suggest that progenitors of bipolar PNs are more massive than those of the other PNs (e.g., Greig 1972; Acker 1980; Kaler 1983; Schwarz & Corradi 1995). In the scenario discussed here the main differences between massive (M Z 2 M on the main sequence) and low-mass progenitors are that the radii of the massive progenitors become much larger on the AGB than on the RGB ( 4), and their envelope mass is much larger than that of low-mass progenitors. 5. Expansion velocities.ècorradi & Schwarz (1995) compare properties of bipolar PNs and elliptical PNs. The only physical property for which the di erence between the groups is larger than the dispersion within each group is the expansion velocity, which is much faster for bipolars (see also for the newly discovered bipolar structure of the PN KjPn 8: Lopez et al. 1997). Although the statistic is very poor, a careful examination of Figure 3 of Corradi &

3 No. 2, 1998 BINARY PROGENITORS FOR BIPOLAR PNs 835 Schwarz (1995) reveals that expansion velocities v are grouped in three regions. The Ðrst is the expansion velocities of elliptical PNs, which are in the range v ^ 5È50 km e s~1, as expected for mass loss from AGB stars. Out e of the 21 bipolar PNs with expansion velocities in Table 1 of Corradi & Schwarz (1995), eight are in this range (three have v \ 50 km s~1). The second group is bipolar PNs with 50 e km s~1 \ v [ 80 km s~1, with Ðve members. It is possible that this group e can be grouped together with the Ðrst one. The third group, of eight members, is bipolar PNs for which v [ 150 km s~1. The three symbiotic nebulae studied by Corradi e & Schwarz have expansion velocities in this range. This group is clearly distinguished from the Ðrst two groups by having velocities that are much higher than the escape velocity from AGB stars. In the binary progenitor scenario the high velocities result from the wind blown by the secondary star during the proto-pn stage, and not from the fast wind in the PN stage. The fast wind will further increase this velocity during the interacting wind phase. 6. Masses of central stars.èwhile the masses of central stars of elliptical PNs are concentrated around M ^ 0.6 c M, the distribution of bipolar PNs is Ñat, in the range M D 0.5È0.8 M (Stanghellini et al. 1993; Stanghellini c 1995). However, as Stanghellini et al. (1993) point out, there is a great uncertainty about this conclusion, which should be treated only as a trend. In the binary model, the strong interaction of the secondary with the primaryïs envelope results in an enhanced mass loss relative to that of a single star. An enhanced mass loss can result from tidal interaction and spinning up the envelope. A more intensive mass loss can result from both Roche lobe overñow and a common envelope at the end of the process (e.g., Iben & Livio 1993 and references therein). This enhanced mass loss can leave a bare core with a much lower mass than would have resulted from a single-star evolution (e.g., NGC 2346). Considering that the progenitors of bipolar PNs are massive main-sequence stars, we get the observed wide range of central star masses in bipolar PNs.Mellema (1997) suggests that bipolar progenitors are massive mainsequence stars because they end with massive cores (the central star of the descendant PN). Massive bare cores evolve faster than low-mass cores do from the AGB to the PN phase. Mellema shows that around the slowly evolving central stars the ionization front has time to smooth out the density contrast between the equatorial plane and the polar directions. The more massive central stars evolve so fast that there is no time for the density proðle to be changed, and hence bipolar PNs are formed. We think that the Ñat distribution in the masses of central stars of bipolar PNs contradicts the scenario proposed by Mellema (1997). 7. Concentration of H in a toroidal equatorial waist.èin a recent paper Kastner 2 et al. (1996) searched for H emission in a sample of more than a hundred PNs. Of 2 these, D40% have been detected. Positively detected PNs tend to be concentrated toward the Galactic plane, as are PNs detected for CO (Huggins et al. 1996). In general, the H emission is concentrated in a ring (or a torus) located in the 2 equatorial plane of bipolar PNs. The presence of molecular gas in bipolar PNs is compatible with fast evolution from the AGB to the PN stage, as proposed for bipolar PNs by Mellema (1997). In the binary model it simply results from a very high mass-loss rate, and hence high density, in the equatorial plane. This also explains the Ðnding of Kastner et al. (1996) that the H emission is much stronger in the 2 equatorial plane than in the lobes. 8. T he fraction of bipolar progenitors.èthe analysis below su ers from two not well-determined numbers. The Ðrst is the minimum mass for the formation of PNs, M. It is somewhere in the range M \ 0.8È1 M. In principle PN a 0.8 M star can form a PN if it PN does not lose too much mass on the RGB. The question is whether these PNs can be observed, since the post-agb evolution time will be long and the nebular mass very low. We take here M \ 0.9 M. The second is the minimum mass for the formation PN of bipolar PNs. Corradi & Schwarz (1995) estimate this minimum mass to be M ^ 1.5 M. Similar minimum mass was found for H -emitting BP PNs (Kastner et al. 1996). Corradi & Schwarz (1995) 2 mention, though, that because the PNs are formed from stars at the end of their evolution, the minimum mass can be higher. We here take this mass to be M \ 2 M. For the maximum main-sequence mass of PN progenitors BP we take 8 M. For M \ 0.9 M, the fraction of PNs that are expected to be PN formed from progenitors more massive than M \ 2 M is For M \ 1.7 M (A stars) and M BP \ 2.3 M (see 4) the BP BP expected fractions are 0.36 and 0.21, respectively. Increasing the minimum PN formation mass to M \ 1 M, and PN with M \ 2 M, the expected fraction of bipolar PNs is BP In calculating these numbers we used a power-law initial mass function of f (M) P M~2.5. Recalling that D11% of all PNs are estimated to be bipolars (Schwarz & Corradi 1995), we conclude that only D30%È40% of massive stars form bipolar PNs. Despite the several uncertainties in this analysis, the fraction of massive progenitors that form bipolar PNs is very similar to the fraction of PN progenitors that interact with close stellar companions, D30%. This number is somewhat uncertain, since di erent estimates put the fraction of PN progenitors that interact with close stellar companions, mainly via a common envelope phase, at 20%È35% (Yungelson et al. 1993; Han et al. 1995; Soker 1997). Therefore, the proposed binary model for bipolar PNs requires that stellar secondaries of any mass (larger than D0.1 M ) can form bipolar PNs, and that the main di erence between massive progenitors (M Z 2 M ) and low-mass progenitors is in the progenitors themselves and not in the nature of the stellar binary companion. If this is indeed the case, it contradicts the explanation of Soker & Livio (1994) for massive progenitors of bipolar PNs. Soker & Livio (1994) assume that bipolar PNs are formed from systems that avoid the common envelope phase, and show that only massive secondaries, and hence massive primaries, comply with this assumption. Since only systems with massive secondaries can form bipolar PNs according to the assumption of Soker & Livio (1994), the fraction of massive stars that form bipolars will be very small, much below 30%. This is in contradiction to the Ðnding above, that D30% of all massive progenitors form bipolars. In the next section we relax the assumption of Soker & Livio (1994) and require that the system avoid the common envelope phase for a large fraction of the interaction time, but not necessarily for the entire evolution. 9. Precessing jets.èseveral bipolar PNs and proto-pns show signatures of precessing jets, also called bipolar rotating episodic jets ÏÏ (BRETs), e.g., KjPn 8 (Lopez et al. 1997), NGC 6210 (Phillips & Cuesta 1996), M1-92 (Trammell & Goodrich 1996). In stellar binary models, where the jets emerge from accretion disks, the companions

4 836 SOKER Vol. 496 can cause the disks, and hence the jets, to precess (e.g., Morris & Reipurth 1991; Larwood et al. 1996). Companions with orbital periods of several hundred years can lead to the formation of point-symmetric PNs (Corradi & Schwarz 1993; Cli e et al. 1995), while Manchado, Stanghellini, & Guerrero (1996) conjecture that companions with an orbital period of several years or more will cause precession which may form PNs with two pairs of lobes (they term these PNs quadrupolar PNs ÏÏ). Livio & Pringle (1996) propose that precessing jets can be formed from the radiation-induced self-warping instability of the accretion disk; a binary companion is not required for the precession itself, but it is required for the formation of the accretion disk. No good model for precession jets exists in single progenitor models. Jets can also be formed from collisions of winds from the two stars of a binary system, as in symbiotic systems (Nussbaumer & Walder 1993). R Aquarii is a symbiotic system with an orbital separation of a ^ 20 AU and variable jets (Paresce & Hack 1994). The binary companion seems to be essential for the formation and variability (or possibly precession: Hollis & Michalitsianos 1993) of the jets in this symbiotic system (Paresce & Hack 1994). 10. Deviation from axisymmetry.èthe deviation from axisymmetry can result from the inñuence of a wide binary companion, i.e., with orbital period comparable to the formation times of the nebula, D500È50,000 yr (Soker 1994), or a close stellar companion in an eccentric orbit (Soker, Harpaz, & Rappaport 1998). In progenitors of bipolar PNs, according to the theme of the present paper, the close stellar companion in an eccentric orbit will also cause the bipolar structure itself. The presence of a tertiary wide star is also possible in progenitors of bipolar PNs. CH Cygni is a known symbiotic triple system (Hinkle et al. 1993), with a 756 day orbital period of the close binary, and a 14.5 yr orbital period of the tertiary star. The displacement caused by a wide binary companion need not be in the equatorial plane of the nebula (for elliptical or bipolar PNs), since the wide binary orbital plane can be inclined to the equatorial plane of the PN. In the eccentric orbit case, however, the displacement of the central star relative to the nebula will be in the equatorial plane. The discussion of this section is summarized in Table 1, which compares the explanations of single and binary progenitor models for the di erent observations. One property we are not referring to is the kinematical age. This is because the kinematical age depends on the most distant point from the central star that is observed. This distance depends greatly on the brightness of the nebular outskirts, which itself depends on many parameters (distance to the PN; mass-loss rate; ionization front; inhomogeneity in the nebula; interaction with the ISM), and therefore it is not easy to analyze. 3. RELEVANT PROCESSES The main goal of the paper, which is an attempt to explain the correlation of bipolar PNs with massive progenitors in the binary progenitor paradigm, is postponed to TABLE 1 SINGLE AND BINARY PROGENITOR MODELS Observation Single Star Binary System Reference Only NGC 2346 is a bipolar Binary companions play no One of the motivations for PN with central binary star major role in shaping PNs the present proposed scenario NGC 2346 The central binary star is a Massive secondaries ] low-mass central 1 (binary system) coincidence stars in long periods are expected in some systems Similarity to some Coincidencea One of the motivations for the present symbiotic nebulae proposed scenariob Massive progenitors Main argument to support single Explanation proposed in the present star progenitors paper ( 4): low-mass stars attain large radii on RGB High expansion velocities From fast-slow winds interaction From mass accreted and blown by 2 (binary system) the companionc Flat mass distribution of In contradiction The secondary expels the envelope 1 (binary system) central stars and may terminate the AGB phase early Concentration of H in Mass concentration and fast The secondary expels a substantial 3 (single star) 2 equatorial waist central star evolution from fraction of the wind in the the AGB equatorial plane D30% of massive progenitors No particular explanation This is about the expected fraction form bipolar PNs of stellar companions with the right properties Precessing jets (variable jets) No good explanation Mass transfer forms a disk which 4, 5, 6 (binary system) might precess (collision of winds from two stars) O -center central star Wide companion Eccentric binary system, or, a 7 (single star), wide companion 8 (binary system) a A more extreme view is that the single progenitor model requires that the secondaries play no major role in shaping symbiotic nebulae. b The similarity to symbiotic was extensively discussed by Corradi and Schwarz in several papers (see references). c Further acceleration is expected from fast-slow wind interaction. REFERENCES.È(1) Iben & Tutukov 1993; (2) Morris 1987; (3) Mellema 1997; (4) e.g., Larwood et al. 1996; (5) Livio & Pringle 1996; (6) Paresce & Hack 1994; (7) Soker 1994; (8) Soker et al

5 No. 2, 1998 BINARY PROGENITORS FOR BIPOLAR PNs 837 the next sections. We now mention some aspects and processes involved when the secondary is outside the primaryïs envelope, and elaborate on a few previously mentioned processes. The discussion here will put the postulate of the next section on more solid ground. The evolution can be divided into three phases: preècommon envelope, common envelope, and postècommon envelope. The secondary will be outside the primaryïs envelope in the Ðrst and last phases. For some aspects of common envelope evolution, which we will not consider here, see Iben & Livio (1993) and Rasio & Livio (1996). We should note that some progenitors of bipolar PNs, those with massive secondaries (M Z 1 M ) and large initial orbital separation, will not enter 2 the common envelope phase at all. We start by mentioning two basic types of accretion, or two ways to divert the mass Ñow: the Roche lobe overñow and accretion from the wind. 1. Roche lobe overñow.èfor binary systems with M ^ M, where M is the secondaryïs mass, Roche lobe overñow 1 occurs 2 when the 2 primary radius is R Z 0.5a, where a is the orbital separation. For typical AGB stellar radii of massive progenitors, R ^ 300È500 R, this condition requires orbital separations of a [ 4 AU. As we show below, these systems will have circular orbits and they will corotate. When the secondary mass at the onset of the Roche lobe overñow obeys M [ 1.7M, the system avoids the 2 1 common envelope phase (Pastetter & Ritter 1989). For a common envelope phase to occur, Pastetter & Ritter (1989) Ðnd that the initial mass ratio should be 1.2 [ M /M [ For progenitors of PNs, this means M Z 1.4 M 1 (Soker & Livio 1994). As mentioned earlier, Soker & Livio (1994) postulate that these systems are the progenitors of bipolar PNs. In the next section we relax their postulate somewhat. 2. Accretion from the wind.èas the secondary moves inside the primaryïs wind, it accretes via the well-known Bondi and Hoyle rate. Since the wind temperature is [1000 K, we neglect the windïs sound speed. The accretion radius is R ^ 2GM /v2, where v is the velocity of the secondary a 2 r r relative to the wind: v ^ (v2]v2)1@2, where v is the wind r w o w velocity and v \ [G(M ] M )/a]1@2 is the orbital velocity o 1 2 of the secondary relative to the primary. In the proposed scenario the relevant values are v ^ 5È20 km s~1, and w v o \ 9.4 AM 1 ] M 2 2 M B1@2A a 20 AUB~1@2 km s~1. (1) Substituting v \ 15 km s~1, v \ 10 km s~1, and the wind density o \ o M0 r in w the expression for the accre- 1 o /(4na2v w ) tion rate M0 we Ðnd for the fraction of the wind 2 \ nr a 2 ov r, that is interrupted by the secondary (for a Z 10 AU and assuming spherically symmetric wind), M0 0.06A B2A 2 o M0 1 o ^ M a 2. (2) M 20 AUB~2 An expression which includes explicitly the dependence on orbital velocity is given by Morris (1990; see also Han et al. 1995). 3. Extended envelope overñow.ïïèanother possibility is that the secondary accretes from the primaryïs extended envelope (Harpaz, Rappaport, & Soker 1997). Since grains are formed at a temperature of D1200 K, while AGB surface temperatures are D3000 K, it has been suggested that there is an extended envelope around evolved AGB stars (e.g., Bowen 1988). In this extended envelope the mass Ñows at a speed lower than the escape velocity, and cools. The extended envelope radius is several times the stellar radius. When the secondary is inside this extended envelope, its inñuence on the mass Ñow can be larger than the expression given by equation (2) for accretion from the wind (Harpaz et al. 1997). This process is hard to quantify, since there is no complete model to describe this extended envelope. 4. Jet formation.èthe presence of precessing jets in several bipolar PNs is one of the critical observations of the list in the previous section. In the binary progenitor paradigm, mass transfer forms an accretion disk. The disk around the secondary, or in some cases around the bare core of the primary, may blows jets out its two sides. This has been suggested, though not demonstrated, for PNs by Morris (1987) and Soker & Livio (1994). 5. T idal spin-up.èin the common tidal model in use, the equilibrium tide mechanism (Zahn 1977, 1989), the circularization time is given by (Verbunt & Phinney 1995) A q \ 7 ] L B~1@3A R B2@3A M B~1 env circ f 2000 L 200 R 0.5M 1 A ] M B1@3AM env 2 B~1A M 1 ] 2 B~1A a yr, (3) 0.5 M M M 5RB8 1 1 where L, R, and M are the luminosity, radius, and total 1 mass of the primary, M is the primaryïs envelope mass, env and f ^ 1 is a dimensionless parameter. The synchronization time is related to the circularization time by the expression q ^ (1 ] M /M )(M /M )~1(I/M R2)(R/a)2q, syn circ where I is the primaryïs moment of inertia. Approximating the envelope density proðle of stars on the upper RGB and AGB by o P r~2, where r is the radial distance from the starïs center, we Ðnd I \ (2/9)M R2. Substituting this in env the expression for the synchronization time, we Ðnd A q \ 3 ] L B~1@3A R B2@3 syn f 2000 L 200 R A ] M B1@3AM env 2 B~2A a yr. (4) 0.5 M M 5RB6 1 Because of the intense mass loss rate from AGB stars, the envelope constantly loses angular momentum. To achieve substantial spinning, we require that the synchronization time be of the order of the mass-loss time M /M0 or env 1, shorter. For the early AGB evolution, which lasts longer and in which the mass-loss rate is relatively small, this condition gives a [ 10R. At late AGB stages, when the super- wind commences, this condition becomes a [ 5R. When synchronization is achieved, the ratio of the primary envelopeïs angular velocity to the surface Keplerian (breakup) velocity in an equal-mass binary system is )/) \ 1.41(R/a)3@2. This amounts to and for a \ Kep 10R and a \ 5R, respectively. Therefore, for the relevant orbital separation the primaryïs envelope will rotate at several percent of the breakup velocity, even if full synchronization is not achieved. This will probably result in a small, but nonnegligible, e ect on the mass-loss geometry, which will add to the other e ects. It will probably cause a higher

6 838 SOKER Vol. 496 mass-loss rate in the equatorial plane. Until there is a better mass-loss theory for AGB stars, we cannot be more quantitative. 6. T idal enhanced surface activity.èby studying the X-ray coronal activity of a large number of rotating single and binary evolved stars, de Medeiros & Mayor (1995 and references therein) conclude that tidal e ects of the secondary enhance the coronal activity. They compare single and binary stars having the same rotational velocities, and Ðnd that the X-ray luminosity of binary systems that achieved circularization is much higher. They conclude that circularization is necessary for enhanced X-ray coronal activity in binary systems, but is not a sufficient condition. High rotation is not a sufficient condition for this activity either. It is not yet clear what the mechanism is through which tidal e ects enhance the activity, though increasing dynamo activity and heating by mechanical waves are the most popular. We expect that, to some degree, similar e ects may exist in AGB stars in binary systems if circularization is achieved. By equation (3) the last condition is a [ 5R D 10 AU, or an orbital period of P [ 20 yr.tout & Eggleton (1988) proposed such a scenario for enhanced mass-loss rate in circular and synchronized orbits. They postulated that by means of tidal torque the companion enhances the massloss rate by a factor that can be formulated as M0 \ M0 0 [1 ] 104(R/R )6], (5) L where R is the primaryïs Roche lobe radius, and R is the L primaryïs radius. The mass-loss rate is doubled for R ^ 0.2R, which for an equal-mass binary system is equal to L R ^ 0.1a. It is likely that the companion will cause a higher mass-loss rate in the equatorial plane. Since most PN progenitors have R \ 3 AU on the AGB, this e ect is signiðcant for orbital separation of a [ 30 AU (for equal-mass binaries). 7. Eccentricity.ÈWhen the secondary has an eccentric orbit with a semimajor axis of a Z 5R D 10 AU, so that by equation (3) circularization is not achieved, it will period- ically come closer to the primary. It can a ect the mass-loss rate and geometry, and may lead to the formation of shells (Harpaz et al. 1997). It may also cause a displacement of some structural features relative to others in the descendant PN (Soker et al. 1998), in particular the displacement of the PN-central star ÏÏ from the geometrical center of the nebula (e.g., MyCn 18, the Hourglass Nebula [Sahai, Trauger, & Evans 1995]; there are more examples in Soker et al. 1998). The displacement is in the equatorial plane of the otherwise axisymmetric PN. For smaller separation we expect circular orbits. However, Waelkens et al. (1996) suggest that the eccentric orbit of the binary system HD (with orbital period of 300 days), in the center of the Red Rectangle nebula can be explained by a circumstellar disk around the binary system. 8. Equatorial concentration by orbital motion.èthe orbital motion of the primary around the center of mass changes its mass-loss geometry. An isotropic wind, for example, is deformed to have a higher velocity and a higher mass-loss rate in the equatorial plane (Soker 1994). We denote by h the angle from the symmetry axis of the descendant PN, and assume an isotropic wind from the primary with a mass-loss rate per unit solid angle of m5 We denote 0. by v \ v M /(M ] M ) the primary velocity around the o1 o center of mass. Averaging over many orbital periods, Soker (1994) Ðnds that the average mass-loss rate per unit solid angle in the descendant PN is given by m5 d (h) \ m5 0C 1 ] Av o1 v w B2A3 4 sin2 h [ 1 2BD, (6) while the average expansion velocity of the nebula is v dw (h) \ v w C 1 ] Av o1 v w B2A5 4 sin2 h [ 1 2BD. (7) Although the qualitative e ect exists for any orbital separation, these speciðc equations are derived for v /v > 1. In addition, the assumption of isotropic wind requires o1 w that the other e ects discussed above (e.g., tidal enhanced surface activity) be small. Overall, we require that for the above e ect to be nonnegligible compared to the other e ects, the orbital separation be a [ 10R D 20 AU. For a Z 20 AU and an equal-mass binary system we Ðnd v /v ^ 0.2È0.5. From equations (6) and (7) we see that this o1 e ect w changes the mass-loss rate and velocity di erently in the polar directions and equatorial plane, with up to D30% di erence. In addition to these processes, there are several other factors that determine the structure of the descendant PN. Among these are the mass of the secondary and its nature, e.g., a white dwarf (WD) or a main-sequence star. If it is a main-sequence star, or even more if it has evolved o the main sequence, it has a wind, and wind collisions from the two stars may shape the nebula as in symbiotic systems (Nussbaumer & Walder 1993). If the secondary is a WD, its radiation ionizes the nebula around the two stars, as in symbiotic systems. This may inñuence grain formation, and hence the mass-loss process. But most important, the nature of the secondary determines the structure of the wind which the secondary blows. The material in this wind is the accreted mass from the primary. The primaryïs Ðnal core mass can also a ect the PN structure (Mellema 1997). Mellema (1997) shows that since a less massive core evolves more slowly, the high density in the equatorial plane is smoothed after ionization starts. This reduces the density contrast between the equatorial plane and the polar direction in a PN having a low-mass central star. All these processes show that there is a rich spectrum of shapes that bipolar PNs can acquire. Common to all the structures, by deðnition, is the presence of two lobes that are formed by a large density contrast between the equatorial plane and the polar direction. In the scenario discussed in this paper, which was suggested in several previous studies (e.g., Morris 1987; Corradi & Schwarz 1993, 1995), a close binary companion, which in most cases avoids the common envelope phase, is the direct cause of this large density contrast. 4. MASSIVE PROGENITORS IN BINARY SYSTEMS 4.1. T he Postulate The scenario discussed below follows the proposed scenario of Morris (1981, 1987, 1990) and the popular model for symbiotic nebulae (e.g., Mikolajewska & Kenyon 1996, where more references and a discussion of some aspects of accretion disks can be found), and it is based on observations 1È3 listed in 2. Soker & Livio (1994) further explore some aspects of this scenario. They consider Roche lobe overñow by the primary, and speculate on the formation of an accretion disk around the secondary, and possible formation of jets by this disk. The basic scenario involves an

7 No. 2, 1998 BINARY PROGENITORS FOR BIPOLAR PNs 839 AGB star that loses mass at a high rate o M0 1 o [ 10~6 M, and a secondary star which accretes a substantial portion of this wind, and then blows a substantial portion of the accreted material. The basic postulate on which the scenario is based, therefore, is as follows. T he progenitors of most bipolar PNs are binary stellar systems in which the secondary diverts a substantial fraction of the mass lost by the AGB primary, but the systems avoid the common envelope phase for a large portion of the evolution. By diverts ÏÏ we refer to the process by which the mass becomes dominated by the secondaryïs gravity (e.g., an accretion disk around the secondary), but the secondary blows a substantial portion of the mass into the nebula rather than accreting it to its surface. The secondary can be outside the primary envelope for the entire evolution, it can be outside before it enters a common envelope phase, and, in some cases, it can blow a substantial mass after it leaves the common envelope phase. In NGC 2346, for example, the orbital separation is 34.9 R (see Iben & Livio 1993), so that when the primaryïs envelope contracted on its way to becoming the central star of this PN, the secondary found itself outside the envelope while the primary was still large (a radius of D15 R ), and con- tained an envelope of few 10~4 M. The secondary was outside the primaryïs envelope for a long time before it entered the common envelope as well. Han et al. (1995) assume that any close interaction in a binary system leads to the formation of a bipolar PN. However, their assumption does not comply with observation 1 of 2 (they get bipolar PNs from all common envelope cases), and with observation 8 (they get too many bipolar PNs). One of the problems is that they do not deðne clearly what they mean by bipolar PNs. ÏÏ They study three main channels of evolution: binary gravitational focusing, common envelope ejection, and binary merger. The binary gravitational focusing is more or less parallel to what we term here strong interaction which avoids the common envelope. Therefore, out of these three channels, it is the binary gravitational focusing which, according to our postulate, leads to the formation of bipolar PNs. They Ðnd that D13% of all PNs are formed from binary gravitational focusing, which is close to the 11% bipolar PNs among all PNs claimed by Corradi & Schwarz (1995). Therefore, the statistical analysis by Han et al. (1995) supports our postulate presented above. As noted in point 8 of 2 the di erence between massive and low-mass primaries is in the primary itself (since most stellar secondaries with M Z 0.1 M can in principle lead to the formation of bipolar 2 PNs). We propose here that this di erence is the ratio of the maximum radius the primary can attain on the AGB (if it evolves undisturbed) to its maximum radius on the RGB. For massive secondaries this ratio is large, while it is close to unity for low-mass stars. Therefore, the binary companions of low-mass stars will interact strongly already on the primaryïs RGB phase. Let us examine this idea more quantitatively. Taking the maximum radii on the AGB, R, and on the RGB, R, from Iben & Tutukov (1985), the following A approximation R (up to D20%) can be Ðtted for the radius ratio as function of mass: log (R A /R R ) \ 3.7 log2 (M/M ) [ 0.37 log (M/M ) ] 0.16, M [ 2.25 M, (8) for stars which develop degenerate helium cores, and log (R /R ) \ 2.2 [ 1.8 log (M/M ), MZ2.35 M, A R (9) where M is the primaryïs mass on the zero-age main sequence. In the narrow range 2.25 M [ M [ 2.35 M there is a steep match between the two Ðts. Several properties of this ratio are relevant to us: (a) This ratio increases with M for M \ 2.3 M [log (M/M ) \ 0.36], reaching values of 1.4, 1.6, 2.3, and 3.1 for M \ 1, 1.5, 2, and 2.25 M, respectively. (b) There is a sharp jump in the radius ratio at M ^ 2.3 M, which is caused mainly by the much smaller radii which massive stars attain on the RGB. (c) The maximum value of this ratio is D30, for zero-age mainsequence mass just above 2.3 M. (d) For M [ 2.3 M this ratio decreases, reaching a value of D4 for M \ 8 M. Approximately, for A and earlier stars the radius ratio is above 2, while for F and G stars this ratio is below T he Evolution to Becoming Bipolar PNs This subsection describes the evolutionary scenario which prevents low-mass stars from forming bipolar PNs. First, we recall that only stellar companions of mass M Z M can potentially form bipolar PNs. Less massive secondary stars, brown dwarfs, and planets will spin up the primaryïs envelope and form elliptical PNs, but cannot signiðcantly divert the mass-loss geometry to form bipolar PNs. The reasons are as follows. Low-mass secondaries will have to be close in order to accrete a substantial fraction of the mass lost by the primary. But since they cannot bring the more massive envelope into corotation, they will end in a common envelope phase (see below). As postulated in the present paper, based on point 1 of 2, the common envelope evolution with low-mass companions can lead to the formation of rings and high density contrast, but not bipolar PNs. Second, when the primary evolves along the upper AGB, where it blows the wind that later forms the nebula, the secondary, according to our postulate, should be outside the envelope, but quite close, in order to strongly inñuence the mass loss. The Ðrst condition reads a Z 1 AU. The sec- ondary can accrete from the wind, for which it should be at an orbital separation of a [ 10(M /M ) AU, by equation 2 (2). Therefore, this process is relevant only for a secondary of mass M Z 0.1 M. Lighter secondaries can cause highly nonspherical 2 mass loss by spinning up the envelope, or by causing the primary to overñow its Roche lobe. The ratio of the synchronized surface velocity to the surface Keplerian velocity, for low-mass secondaries, is v /v ^ (R/a)3@2. For 10% breakup velocity we require spin a [ Kep 5R ^ 10 AU, although we expect that higher rotation velocities are necessary for bipolar PNs. Since we expect circularization as well, by equation (3), the secondary will enhance and inñu- ence the mass-loss process, as we discussed in the previous section. To conclude, the orbital separation when the primary enters the AGB should be in the range of 1 AU [ a [ 5R ^ 10 AU in order to form bipolar PNs. Of course, the orbital separation will change as the primary evolves along the AGB due to mass loss, mass transfer, and tidal interaction. In particular, low-mass secondaries will spiral in somewhat as they spin up the primaryïs envelope. We are now in a position to understand the reason why low-mass primaries are unlikely to form bipolar PNs: for low-mass stars, M [ 1.5 M, the conditions for strong 1 interaction on the AGB imply also strong interaction on the

8 840 SOKER Vol. 496 RGB. The radius on the AGB is larger, but the time the star spends on the AGB is shorter than that on the RGB, and because of mass loss on the early AGB the orbital separation will increase further. Verbunt & Phinney (1995) follow the evolution of binary systems as the primary star evolves from the main sequence to the AGB, and show that for primaries less massive than D1.5 M most of the tidal interaction takes place on the RGB. This will cause circularization and enhanced mass loss. Han et al. (1995) Ðnd that for primaries with initial mass of M \ 2 M (in their presentation all primaries with masses 1 below 2 M are in one group), in most cases common envelope interaction takes place already on the RGB (they term it FGB [Ðrst giant branch]). These results show that the primary, which is already of low mass, will end the RGB with less mass than it would if it evolved as a single star. Low-mass stars lose D0.2 M while on the RGB. Increasing the mass-loss rate by a factor of D3.5, as is predicted by equation (5) for orbital separation a \ 5RÈ8R (depending on the secondary mass), will leave stars of initial mass [ 1.5 M with a very low envelope mass. These stars will not reach a high point on the AGB, as they would have done without the enhanced mass-loss rate. Therefore, their radii will be smaller on the AGB and they will not interact strongly with their companions. We see that, in addition to the radius ratio, the envelope mass plays a role in the proposed scenario as well. In many cases the secondary will enter the primaryïs envelope and spiral inward on the RGB. The secondaries can either collide with the cores already during the primariesï RGB phase, or reach small radii and then enter the primariesï envelopes again on the early primariesï AGB phase. Therefore, during the entire AGB phase the secondary will be inside the envelope, unless it collides with the primaryïs core. In these cases we expect to form elliptical, rather than bipolar, PNs. To conclude, the probability of primary stars with M [ 1.5 M forming bipolar PNs is extremely small. The prob- ability is hard to estimate, since the formation of bipolar PNs from low-mass stars seems to require Ðne tuning. The uncertainties in the strength of the tidal interaction and the inñuence of rotation and tidal forces on the mass loss from RGB stars are too large to allow simple estimates. A detailed evolutionary study of rotating and interacting RGB stars, followed by a study of the systems as they evolve to the AGB is required. Because of an increase of the radius ratio, equations (8) and (9), and an increase of the envelope mass, the chance to form bipolar PNs, for secondaries of M Z 0.1 M and orbital separation 1 AU [ a [ 10 AU, increases 2 for higher mass primaries. Primary stars of initial mass M [ 2.3 M will form bipolar PNs in almost all cases when the companion is in the right radial range; this occurs for D30% of these primary stars, as explained in item 8 of 2. When the orbital separation is very small, even massive progenitors will interact with their companions on the RGB. The inñuence of the interaction during the RGB, with both low-mass and massive primaries, on the later AGB phase deserve a more detailed study, which is beyond the scope of this paper. One rare route for the formation of bipolar PNs leaves a helium central star (Iben & Tutukov 1993). These are systems in which the secondary enters a common envelope phase on the RGB, and ejects the entire envelope. If the secondary is of low mass and/or the initial orbital separation is small (Han et al. 1995), the secondary will merge with the core before ejecting the envelope. These systems, now a single more massive star, will evolve to the horizontal branch and then to the AGB. Due to the large envelope angular momentum, they will form elliptical PNs. If the secondary survives and ejects the entire envelope, we are left with a helium WD with a nebula around it. Because of the helium coreïs low luminosity and its slow evolution from the RGB to the point when it starts ionizing the nebula (Iben & Tutukov 1993), these systems will be faint elliptical PNs, or will not be detected as PNs at all. Mellema (1997) shows that slow evolution results in a loss of density contrast, and less bipolar structure. In rare cases, the envelope will be ejected when the helium core is massive, M [ 0.4 M. These systems may become bipolar PNs, as core Iben & Tutukov (1993) suggest for NGC However, it is not clear that the primary in NGC 2346 is a helium WD (and not a CO WD), and in any case these systems, in which the envelope is ejected only at the tip of the RGB, are very rare, since they need Ðne tuning of the initial orbital separation and secondary mass (e.g., in NGC 2346 the secondary is massive [point 2 of 2]). 5. SUMMARY The goal of this paper is to propose an explanation for the positive correlation of bipolar planetary nebulae with massive progenitors, in the paradigm of binary system progenitors. We suggest that the main di erence between massive and low-mass progenitors is the larger ratios of RGB to AGB radii which low-mass stars attain. These larger radii on the RGB cause most stellar binary companions, which potentially could have formed bipolar PNs if the primary had been on the AGB, to interact with lowmass primaries already on the RGB. These systems may enter a common envelope phase on the RGB. In some cases the strong interaction will cause the primary to lose most, or even all, its envelope on the RGB; such systems will form only faint PNs, or no PN at all. In other cases of common envelope evolution the secondary will spiral in, but the primary will retain its envelope, and eventually evolve to the AGB phase. Such systems will enter a common envelope phase early on the AGB, and form elliptical, rather than bipolar, PNs. Massive primaries, on the other hand, reach large radii only on the AGB, and therefore their companions interact mainly on the AGB, the stage prior to the PN phase. In addition, more massive primaries retain much more massive envelopes, which result in high density concentration in the equatorial plane of the descendant PNs. To further support binary models for formation of bipolar PNs, we list 10 critical observations, and argue that single-star models for the formation of bipolar planetary nebulae have difficulties complying with these observations, while binary systems models do much better. This is summarized in Table 1. Even if the particular explanation proposed here for the correlation of bipolar PNs with massive progenitors is eventually superseded, the other nine critical observations listed in 2 still support binary models over single-star models. The presently proposed scenario should be further tested both observationally and theoretically. Observationally, the scenario predicts that the central stars of most bipolar planetary nebulae are in binary systems having orbital periods in the range of a few days to few times 10 yr; hence, bipolar PNs should be carefully searched for such companions. Theoretically, it will be very constructive if the mass trans-