Reduced Maximum Mass-Loss Rate of OH/IR Stars Due to Overlooked Binary Interaction

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Reduced maximum mass-loss rate of OH/IR stars due to overlooked binary interaction

DOI: 10.1038/s41550-019-0703-5

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Citation for published version (APA): Decin, L., Homan, W., Danilovich, T., de Koter, A., Engels, D., Waters, L. B. F. M., Muller, S., Gielen, C., Garca- Hernandez, D. A., Stancliffe, R., Van de Sande, M., Molenberghs, G., Kerschbaum, F., Zijlstra, A., & El Mellah, I. (2019). Reduced maximum mass-loss rate of OH/IR stars due to overlooked binary interaction. Nature Astronomy, 3(5), 408-415. https://doi.org/10.1038/s41550-019-0703-5 Published in: Nature Astronomy

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Download date:08. Oct. 2021 Reduced maximum mass-loss rate of OH/IR stars due to overlooked binary interaction

L. Decin1,2,?, W. Homan1, T. Danilovich1, A. de Koter1,3, D. Engels4, L. B. F. M. Waters3,5, S. Muller6, C. Gielen1, D. A. Garc´ıa-Hernandez´ 7,8, R. J. Stancliffe9,10, M. Van de Sande1, G. Molenberghs11,12, F. Kerschbaum13, A. A. Zijlstra14 & I. El Mellah15

1Instituut voor Sterrenkunde, KU Leuven, Celestijnenlaan 200D, 3001 Leuven, Belgium ?e-mail: [email protected] 2University of Leeds, School of Chemistry, Leeds LS2 9JT, United Kingdom 3Astronomical Institute Anton Pannekoek, University of Amsterdam, Science Park 904, PO Box 94249, 1090 GE, Amsterdam, The Netherlands 4Hamburger Sternwarte, Gojenbergsweg 112, D-21029 Hamburg, Germany 5SRON Netherlands Institute for Space Research, Sorbonnelaan 2, 3584 CA Utrecht, The Netherlands 6Chalmers University of Technology, Department of Space, Earth and Environmnet, Onsala Space Obser- vatory, S-439 92 Onsala, Sweden 7Instituto de Astrof´ısica de Canarias, E-38205 La Laguna, Tenerife, Spain 8Departamento de Astrof´ısica, Universidad de La Laguna (ULL), E-38206 La Laguna, Tenerife, Spain 9E. A. Milne Centre for Astrophysics, Department of Physics & Mathematics, University of Hull, HU6 7RX, UK 10School of Physics and Astronomy, University of Birmingham, Birmingham B15 2TT, UK 11I-BioStat, Universiteit Hasselt, Martelarenlaan 42, 3500, Hasselt, Belgium 12I-BioStat, KU Leuven, Kapucijnenvoer 35, 3000, Leuven, Belgium 13University of Vienna, Department of Astrophysics, Turkenschanzstrasse¨ 17, 1180 Wien, Austria 14Jodrell Bank Centre for Astrophysics, School of Physics & Astronomy, University of Manchester, Manch- ester, UK 15Centre for mathematical Plasma-Astrophysics, KU Leuven, Celestijnenlaan 200B, 3001 Leuven, Belgium

In 1981, the idea of a superwind that ends the life of cool giant stars was proposed1. Extreme OH/IR- 4 2–12 stars develop superwinds with the highest mass-loss rates known so far, up to a few 10 M /yr , hence epitomizing our understanding of the maximum mass-loss rate achieved during the Asymptotic Giant Branch (AGB) phase. A condundrum arises whereby the observationally determined duration of the superwind phase is too short for these stars to become white dwarfs 2–4, 6, 8–10. Here, we report on the detection of spiral structures around two cornerstone extreme OH/IR-stars, OH 26.5+0.6 and OH 30.1 0.7, identifying them as wide binary systems. Hydrodynamical simulations show that the companion’s gravitational attraction creates an equatorial density enhancement mimicking a short extreme superwind phase, thereby solving the decades-old conundrum. This discovery restricts the

1 maximum mass-loss rate of AGB stars around the single-scattering radiation-pressure limit of a few 5 10 M /yr. This brings about crucial implications for nucleosynthetic yields, planet survival, and the wind-driving mechanism.

All stars with initial mass between 0.8-8 solar masses (M ) lose 40-80% of their mass via a dust- ⇠ driven wind during the Asymptotic Giant Branch (AGB) phase, their ﬁnal destiny being white dwarfs. This wind is caused by radiation pressure on grains formed in the extended cool atmosphere due to pulsations13. 7 For mass-loss rates above 10 M /yr, the associated timescale for stars to shed their mantle is much ⇠ shorter than the nuclear burning timescale, such that mass loss determines the further evolution13. An empirical relation between mass-loss rate and luminosity is given by Reimers’ original formula14 widely used in evolutionary calculations: M˙ = 4 10 13⌘L/gR, with M˙ the mass-loss rate in units of solar masses ⇥ per year (M /yr), ⌘ a unit-less parameter, and the stellar luminosity L, gravity g, and radius R in solar units. For ⌘ 1, the AGB lifetime is of the order of one million years, and the maximum AGB mass- ⇠ 6 1 loss rate is only a few 10 M /yr . Hence, a Reimers-like wind cannot explain the distribution of white dwarf masses and characteristics of the planetary-nebulae population1. Therefore, Renzini proposed in 1981 that a superwind (see Supplementary Information for terminology) with mass-loss rate of at least a 5 few 10 M /yr can develop at the high luminosity tip of the AGB phase leading to the ejection of the residual envelope on a timescale that is very short compared with the overall AGB lifetime. The debate over a sudden onset of the superwind phase or a gradual increase in mass-loss rate has not yet been settled, neither has the maximum mass-loss rate been established (see Supplementary Information).

Since 1981, many AGB stars were found to have mass-loss rates consistent with the suggested superwind values15–17. In this context, extreme OH/IR-stars are key objects. They represent a class of intermediate-mass (2–8 M ) oxygen-rich stars at the tip of the AGB evolution with the highest mass-loss 5 4 rates discovered (between 10 and a few 10 M /yr) (see Supplementary Information). As such, these stars underpin a crucial part of our understanding of the maximum AGB mass-loss rate achieved, and hence of the process by which stars return mass to the interstellar medium. However, for more than a dozen extreme OH/IR-stars it has been shown that the superwind phase lasts only 200–1 000 years and was preceded ⇠ 6 2–4,9, 10 by a phase with mass-loss rate one or two orders of magnitude lower (around a few 10 M /yr) . Given an expected superwind lifetime of some ten thousand years to allow the stars to reduce their mass below the Chandrasekhar limit, the probability of observing a dozen stars with almost the same superwind age (of 200–1 000 years) is unrealistically small (see Methods). Statistical reasoning implies that the 4 ⇠ 10 M /yr superwind duration for these extreme OH/IR-stars cannot be much longer than 1 000 years ⇠ ⇠ (see Methods). This duration is, however, too short for these stars to become white dwarfs resulting in a conﬂict with initial-ﬁnal mass relations deduced from young stellar clusters18.

Two prototype stars of this class of extreme OH/IR-stars have been observed with ALMA. Obser- vations of low-excitation rotational lines of 12CO were obtained for OH 26.5+0.6 and OH 30.1 0.7 (see

2 Fig. 1, Fig. 2, see Methods, see Supplementary Information). The CO data show a centralized bright emission region with full width at half maximum (FWHM) of 0. 4–0. 6. This value is similar to the FWHM ⇠ 00 00 of the dusty envelope which was interpreted as a measure for the superwind envelope with duration of only 130–200 ( 440–650) years for OH 26.5+0.6 (OH 30.1 0.7)4, 9, 10. A distinct feature in the ALMA data ⇠ ⇠ of both targets is an incomplete ring-like pattern (see Fig. 1, Fig. 2). The position-velocity diagrams (see Fig. 3) reveal that this incomplete ring-like pattern is characteristic for a shell-like spiral created by the orbital motion of the mass-losing AGB star around the center of mass of a binary system19,20 (see Meth- ods, see Supplementary Information). We deduce an orbital period of 430 ( 590) years and mean binary ⇠ ⇠ separation of 24 ( 38) stellar radii (R ) for OH 26.5+0.6 (OH 30.1 0.7) (see Methods). ⇠ ⇠ ?

The detection of a binary-induced shell-like spiral in both OH/IR-stars puts the evolutionary and statistical problem of the short superwind phase in extreme OH/IR-stars in a new perspective. Hydrodynamical simulations for a binary system resembling OH 26.5+0.6 (Fig. 15 in Mastrodemos et al.21) show a wind pattern similar to the ALMA data with the appearance of an incomplete ring-like pattern arising from a spiral shock with spiral-arm spacing (of 400 R ) comparable to the one derived for OH 26.5+0.6 ( 1 or ⇠ ? ⇠ 00 350 R ). Of prime importance for this study is the emergence of a compact high-density region within ⇠ ? 150 R from the primary star caused by the gravitational focusing by the binary companion (Fig. 15 in ⇠ ? Mastrodemos et al.21). This extent of 150 R corresponds to a radial crossing time of 100 years, much ⇠ ? ⇠ in resemblance with the duration of the high mass-loss (superwind) phase of OH 26.5+0.6 ( 130 years). ⇠ This compact equatorial density enhancement (CEDE) has a density contrast of a factor 10 compared to ⇠ the background intrinsic wind density19 (see Methods). This CEDE renders superﬂuous the concept of a very short extreme superwind phase in OH 26.5+0.6, i.e. the compact higher density region in the vicinity of OH 26.5+0.6 —- and likewise in OH 30.1 0.7 –– is caused by the direct gravitational interaction with a binary companion and does not signal an abrupt change in mass-loss rate. The ALMA position-velocity diagrams and asymmetries reported in other observational diagnostics allow us to deduce the orientation of the lower density orbital axis and higher density orbital plane (see Fig. 4, see Methods).

The ALMA data of OH 26.5+0.6 and OH 30.1-0.7, and statistical reasoning imply a general argument 4 that a large fraction, if not all, extreme OH/IR-stars reported to have a short 10 M /yr superwind phase ⇠ are part of a binary system. With an observed binary frequency for intermediate mass stars of at least 30– 45%, increasing to 60–100% for extreme OH/IR stars with initial mass around 5–8 M 22,23, the chance of observing the imprints of binary interaction is high. As such, the intense superwind phase in extreme OH/IR- stars is neither short nor problematic. A CEDE resolves the statistically unlikely situation of detecting a 4 dozen stars having the same short 10 M /yr superwind age (of 200–1000 years) if the superwind ⇠ ⇠ lifetime would be much longer, as required from evolutionary considerations. Moreover, one cannot use the measured geometrical extent of the compact high-density region to estimate a ‘duration’ since the velocity vector ﬁeld does not follow a radially-streaming pattern in the vicinity of the binary stars21.

3 The idea of a binary-induced CEDE mimicking a short intense superwind phase also explains (i) the continuous distribution of infrared colors of observed OH/IR samples, (ii) the absence of evidence of episodic superwind bursts, (iii) the bipolar or multipolar morphology frequently detected in (pre-)planetary nebulae (PNe), and (iv) the detection of micron-sized grains in some (pre-)PNe (see Supplementary Infor- mation). An important consideration is that all extreme OH/IR-stars with suggested short intense superwind phase display crystalline silicate features. Since we attribute the density enhancement to binary interaction, and following argumentation on crystal formation24,25, we suggest processing of grains into crystalline lattice structures in the CEDE (see Methods). Conversely the presence of crystalline silicates in the envelope surrounding OH/IR-stars can be seen as an indicator for binary interaction creating an environment facilitating the crystallization of grains.

A crucial implication of this paradigm shift, linking the behavior of extreme OH/IR-stars to binary interaction, deals with stellar evolution, in particular with the maximum mass-loss rate attained during the AGB phase. This rate rules the fate of the star and in a broader context dictates the chemical enrichment of the interstellar medium by evolved stars. Observational studies26 suggest that the maximum mass-loss rate during the superwind phase can exceed the single-scattering radiation pressure limit (M˙ = L/vc, with 5 v the wind velocity, c the speed of light, and typical values around a few 10 M /yr) by a factor of 10. ⇠ Dynamical models put forward a more conservative factor around 1–227,28. The realization that all reported high mass-loss rate values concern extreme OH/IR-stars (in binary systems) and are based on simpliﬁed 1D analyses of warm (crystalline) dust grains residing close to the star, hence in the CEDE, leads to the conclusion that the reported mass-loss rates have been overestimated with a factor of a few up to 100 ⇠ (see Methods). Low-excitation CO lines are the prime accurate probes of the mass-loss rate yielding values 5 around (0.1–3) 10 M /yr (see Methods). This establishes a threshold for the maximum mass-loss rate ⇠ ⇥ around the single-scattering radiation pressure limit. This conclusion has important consequences for stellar evolution models implementing mass-loss rate prescriptions (see Supplementary Information) and for our understanding of the wind-driven mechanism in AGB stars. Lowering the maximum mass-loss rate in stellar evolution models implies an increase in the AGB lifetime and impacts the nucleosynthetic yields returned to the interstellar medium both in terms of light and heavy (s-process) elements, due to the occurrence of extra thermal pulses and more third dredge-up episodes, with expected differences up to an order of magnitude29. A lower mass-loss rate over a longer time lapse also allows for a more progressive mass transfer and possibly a higher total AGB material captured by the companion in binary systems undergoing wind-mass transfer. This can explain why a fraction of carbon enhanced metal poor stars have accreted as much as 0.1–0.2 solar masses of material from their AGB companion30. A lower mass-loss rate also facilitates the survival of close- by gas planets by avoiding thermal evaporation31. State-of-the-art hydrodynamical models investigating the wind driving in oxygen-rich AGB stars based on stellar pulsations and radiation pressure on dust need 5 27 substantial ﬁnetuning to reach mass-loss rates of 10 M /yr and very speculative mechanisms have ⇠ been invoked32 to explain mass-loss rates substantially beyond this limit. For a maximum mass-loss rate not exceeding the single-scattering radiation pressure limit, we infer that the dust-driven wind scenario can

4 explain the mass loss in all AGB stars with pulsation periods above 300 days. ⇠

References

1. Renzini, A. Red giants as precursors of planetary nebulae. In Iben, I., Jr. & Renzini, A. (eds.) Physical Processes in Red Giants, vol. 88 of Astrophysics and Space Science Library, 431–446 (1981).

2. Heske, A., Forveille, T., Omont, A., van der Veen, W. E. C. J. & Habing, H. J. Deﬁciency of CO emission from massive envelopes around cool OH/IR stars. A&A 239, 173–185 (1990).

3. Delfosse, X., Kahane, C. & Forveille, T. Superwind in evolved OH/IR stars. A&A 320, 249–256 (1997).

4. Chesneau, O. et al. The mid-IR spatially resolved environment of OH 26.5+0.6 at maximum luminosity. A&A 435, 563–574 (2005).

5. van Loon, J. T., Cioni, M.-R. L., Zijlstra, A. A. & Loup, C. An empirical formula for the mass-loss rates of dust-enshrouded red supergiants and oxygen-rich Asymptotic Giant Branch stars. A&A 438, 273–289 (2005).

6. Justtanont, K., Olofsson, G., Dijkstra, C. & Meyer, A. W. Near-infrared observations of water-ice in OH/IR stars. A&A 450, 1051–1059 (2006).

7. Groenewegen, M. A. T., Sloan, G. C., Soszynski,´ I. & Petersen, E. A. Luminosities and mass-loss rates of SMC and LMC AGB stars and red supergiants. A&A 506, 1277–1296 (2009).

8. Groenewegen, M. A. T. An extension of the DUSTY radiative transfer code and an application to OH 26.5 and TT Cygni. A&A 543, A36 (2012).

9. Justtanont, K. et al. OH/IR stars and their superwinds as observed by the Herschel Space Observatory. A&A 556, A101 (2013).

10. de Vries, B. L. et al. The problematically short superwind of OH/IR stars. Probing the outﬂow with the 69 µm spectral band of forsterite. A&A 561, A75 (2014).

11. Goldman, S. R. et al. The wind speeds, dust content, and mass-loss rates of evolved AGB and RSG stars at varying metallicity. MNRAS 465, 403–433 (2017).

12. McDonald, I., De Beck, E., Zijlstra, A. A. & Lagadec, E. Pulsation-triggered dust production by asymptotic giant branch stars. MNRAS 481, 4984–4999 (2018).

13. Hofner,¨ S. & Olofsson, H. Mass loss of stars on the asymptotic giant branch. Mechanisms, models and measurements. A&AR 26, 1 (2018).

5 14. Reimers, D. Circumstellar absorption lines and mass loss from red giants. Memoires of the Societe Royale des Sciences de Liege 8, 369–382 (1975).

15. Knapp, G. R. & Morris, M. Mass loss from evolved stars. III - Mass loss rates for ﬁfty stars from CO J = 1-0 observations. ApJ 292, 640–669 (1985).

16. Bedijn, P. J. Dust shells around Miras and OH/IR stars - Interpretation of IRAS and other infrared measurements. A&A 186, 136–152 (1987).

17. Wood, P. R. et al. OH/IR stars in the Magellanic Clouds. ApJ 397, 552–569 (1992).

18. Catalan,´ S., Isern, J., Garc´ıa-Berro, E. & Ribas, I. The initial-ﬁnal mass relationship of white dwarfs revisited: effect on the luminosity function and mass distribution. MNRAS 387, 1693–1706 (2008).

19. Kim, H. & Taam, R. E. Wide Binary Effects on Asymmetries in Asymptotic Giant Branch Circumstellar Envelopes. ApJ 759, 59 (2012).

20. Homan, W. et al. Simpliﬁed models of stellar wind anatomy for interpreting high-resolution data. Analytical approach to embedded spiral geometries. A&A 579, A118 (2015).

21. Mastrodemos, N. & Morris, M. Bipolar Pre-Planetary Nebulae: Hydrodynamics of Dusty Winds in Binary Systems. II. Morphology of the Circumstellar Envelopes. ApJ 523, 357–380 (1999).

22. Raghavan, D. et al. A Survey of Stellar Families: Multiplicity of Solar-type Stars. ApJS 190, 1–42 (2010).

23. Duchene,ˆ G. & Kraus, A. Stellar Multiplicity. ARA&A 51, 269–310 (2013).

24. Molster, F. J. et al. Low-temperature crystallization of silicate dust in circumstellar disks. Nat 401, 563–565 (1999).

25. Edgar, R. G., Nordhaus, J., Blackman, E. G. & Frank, A. The Formation of Crystalline Dust in AGB Winds from Binary-induced Spiral Shocks. ApJ 675, L101 (2008).

26. van Loon, J. T. et al. Mass-loss rates and luminosity functions of dust-enshrouded AGB stars and red supergiants in the LMC. A&A 351, 559–572 (1999).

27. Bladh, S., Hofner,¨ S., Aringer, B. & Eriksson, K. Exploring wind-driving dust species in cool luminous giants. III. Wind models for M-type AGB stars: dynamic and photometric properties. A&A 575, A105 (2015).

28. Schroder,¨ K.-P., Winters, J. M. & Sedlmayr, E. Tip-AGB stellar evolution in the presence of a pulsating, dust-induced “superwind“. A&A 349, 898–906 (1999).

6 29. Karakas, A. I. et al. Heavy-element yields and abundances of asymptotic giant branch models with a Small Magellanic Cloud metallicity. MNRAS 477, 421–437 (2018).

30. Abate, C., Stancliffe, R. J. & Liu, Z.-W. How plausible are the proposed formation scenarios of CEMP- r/s stars? A&A 587, A50 (2016).

31. Villaver, E. & Livio, M. Can Planets Survive Stellar Evolution? ApJ 661, 1192–1201 (2007).

32. Szyszka, C., Zijlstra, A. A. & Walsh, J. R. The expansion proper motions of the planetary nebula NGC 6302 from Hubble Space Telescope imaging. MNRAS 416, 715–726 (2011).

33. Decin, L., Richards, A. M. S., Danilovich, T., Homan, W. & Nuth, J. A. ALMA spectral line and imaging survey of a low and a high mass-loss rate AGB star between 335 and 362 GHz. A&A 615, A28 (2018).

34. Kim, H., Hsieh, I.-T., Liu, S.-Y. & Taam, R. E. Evidence of a Binary-induced Spiral from an Incomplete Ring Pattern of CIT 6. ApJ 776, 86 (2013).

Acknowledgements This paper uses the ALMA data ADS/JAO.ALMA2015.1.00054.S, 2016.1.00005.S, and 2016.2.00088.S. ALMA is a partnership of ESO (representing its member states), NSF (USA) and NINS (Japan), together with NRC (Canada) and NSC and ASIAA (Taiwan), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO and NAOJ. This paper makes use of the CASA data reduction package: http://casa.nra.edu – Credit: International consortium of sci- entists based at the National Radio Astronomical Observatory (NRAO), the European Southern Observatory (ESO), the National Astronomical Observatory of Japan (NAOJ), the CSIRO Australia Telescope National Facility (CSIRO/ATNF), and the Netherlands Institute for Radio Astronomy (ASTRON) under the guidance of NRAO. L.D., T.D., W.H., and M.V.d.S. acknowledge support from the ERC consolidator grant 646758 AEROSOL, T.D. acknowledges support from the Fund of Scientiﬁc Research Flanders (FWO). D.A.G.H. acknowledges support provided by the Spanish Ministry of Economy and Competitiveness (MINECO) under grant AYA-2017-88254-P. We acknowledge the help of Carl Gottlieb (Harvard University, US) for his editorial advice on the manuscript.

Author Contributions L.D. identiﬁed the spiral structure in the ALMA data of OH 26.5+0.6 and OH 30.1 0.7, performed the full analysis, and lead the consortium, W.H., T.D., and A.d.K. contributed to the interpretation of the data, D.E., D.A.G.H. and S.M. proposed the ALMA observations (ALMA proposals 2015.1.00054.S, 2016.1.00005.S, and 2016.2.00088.S), S.M. reduced the ALMA data, D.E. did the sample analysis of the extreme OH/IR stars, G.M. give advice on the statistical matters, Y.E.M. ran the ballistic simulations, C.G. made Fig. 4, all authors contributed to the discussion.

Author Information Reprints and permissions information is available at www.nature.com/reprints.

7 Correspondence Correspondence and requests for materials should be addressed to Leen Decin (email: [email protected]).

Competing Interests The authors declare that they have no competing ﬁnancial interests.

Supplementary Information is linked to the online version of the paper at www.nature.com/nature.

Figure 1: ALMA 12CO J=3-2 channel map of OH 26.5+0.6. The map shows the incomplete ring-like pattern due to a binary-induced shell-like spiral structure. The ordinate and co-ordinate axis give the offset of the right ascension and declination, respectively. North is up, East is left. The ALMA beam size is shown as a white ellipse in the bottom left corner of each panel. The systemic velocity is at 28 km/s. An animated video slicing through the velocity frames of the ⇠ channel maps is provided in Supplementary Video 1. The contrast is best visible on screen.

Figure 2: ALMA 12CO J=3-2 channel map of OH 30.1 0.7. The map shows the incomplete ring-like pattern due to a binary-induced shell-like spiral structure. The white arrows indicate the knots caused by the overlap of gravitationally induced density wake by the companion with the spiral-shell pattern19 (see Supplementary Information). The systemic velocity is at 99 km/s. An ⇠ animated video slicing through the velocity frames of the channel maps is provided in Supplemen- tary Video 2. The contrast is best visible on screen.

8 Figure 3: Position-velocity (PV) diagrams of the 12CO J = 3–2 emission in OH 26.5+0.6 (top) and of the the 12CO J = 2–1 emission in OH 30.1-0.7 (bottom). Top panel: PV-diagram of the 12 CO J = 3–2 emission in OH 26.5+0.6 for a position angle (PA) of 5 (left) and 85 (right). The resulting PV-diagrams are characteristic for a binary-induced shell-like spiral structure19,20. The binary system is viewed at an inclination angle, i, around 60–75 (see Fig. 11 and Fig. 13 in Homan et al.20). The systemic velocity is at 28 km/s. The lack of emission at velocities between ⇠ 12 – 16 km/s is caused by the blue-wing absorption33. Bottom panel: PV-diagram of the 12CO

J = 2–1 emission in OH 30.1-0.7 for a position angle (PA) of 75 (left) and 165 (right). As in the case of OH 26.5+0.6, the resulting PV-diagrams are characteristic for a binary-induced shell-like spiral structure19,20. In the PV-diagram along the line of nodes of the knots (see Fig. 2), hence at a PA of 165, the knotty structure is present at the systemic velocity (of 99 km/s, see white ⇠ arrow), conﬁrming the existence of the companion’s wake in the orbital plane34. The binary system is viewed at an inclination angle, i, around 50–70. The lack of emission at velocities between 80 – 85 km/s is caused by the blue-wing absorption33. The contrast is best visible on screen.

Figure 4: Sketch of the OH 26.5+0.6 binary system. Sketch of the inner 1000 wind region of OH 26.5+0.6 indicating the different observations that signal a binary companion. The reﬂect motion of the primary OH/IR-star causes a one-arm spiral structure with an extent almost reaching the orbital axis. This spiral structure is seen at a high inclination angle (i = 60 – 75) and as such reveals itself in an incomplete ring-like pattern (see grey areas in the sketch). From the ALMA data, we infer an orbital axis with a position angle (PA) of 5 10. A higher density region ± with FWHM of 000.6 (in pink) is detected in molecular (CO, SiO) and dust emission. This region ⇠ was hitherto misinterpreted as a sign for a very short superwind phase. We show that this region is an equatorial density enhancement situated in the orbital plane caused by the direct gravitational focussing of the binary companion.

9 Methods

4 The statistical probability of observing stars having a 10 M /yr superwind with age of 1 000 ⇠ ⇠ years

For extreme OH/IR-stars (see Supplementary Information for the deﬁnition of this class of objects) like OH 26.5+0.6 and OH 30.1 0.7 with an initial mass between 4–8 solar masses, the estimated Asymptotic Giant Branch (AGB) lifetime during the thermally pulsating phase is between 0.01–0.35 million years35. More than a dozen extreme OH/IR stars within a distance of 4 0000 parsec have been shown to have 4 ⇠ a current 10 M /yr superwind that was preceded by a phase with mass-loss rate one to two orders ⇠ 6 2–4, 6, 8–10 of magnitude lower (mean mass-loss rate 7.4 10 M /yr) (see Supplementary Information). ⇠ ⇥ Modelling suggests that the onset of the present-day high mass-loss rate phase was only 200–1 000 years ⇠ 2–4, 6, 8–10 4 ago . For a superwind with mass-loss rate of 10 M /yr a superwind lifetime of some ten thou- ⇠ sand years is required so that these stars lose enough mass to reach white dwarf masses. The probability of observing a dozen stars with almost the same superwind age (of 200 – 1 000 years) is extremely small if the lifetime of the superwind is of the order of ten thousand years. This can easily be derived by deﬁning

N: the number of stars in the population, • n: the number of stars observed; obviously n N, •  p: the number of ‘successful’ stars, i.e. stars within the window speciﬁed, in the case having a • 4 10 M /yr superwind with age of 1 000 years; obviously 0 p n, ⇠ ⇠   ⇡: probability that a star falls within the speciﬁed window; for this problem ⇡ 1 000/10 000 = 0.1. • ⇠

Assuming observations are independent of each other, the probability of observing p ‘successes’ out of n is binomial: n p n p P (p n, ⇡)= ⇡ (1 ⇡) . | p ! One can roughly estimate the number of AGB stars which are type II OH/IR-stars (see Supplementary In- formation for terminology) and with infrared colours similar to OH 26.5+0.6 and OH 30.1 0.7, using the space densities derived by Ortiz et al.36. This yields 400 stars within a sphere of radius 4 000 parsec ⇠ ⇠ around the Sun, hence N 400. For a dozen of extreme OH/IR-stars, all necessary (spectroscopic) data of ⇠ the infrared dust and CO emission are available to estimate the age of the superwind. For all these galactic extreme OH/IR-stars, a short superwind age was retrieved2–4, 6, 8–10. Hence, p = n 12 or ⇡

12 12 0 12 12 P (p = 12 n = 12,⇡)= ⇡ (1 ⇡) = ⇡ = 10 . | 12 ! This low probability implies that the superwind lifetime in intermediate-mass stars cannot be of the order 4 of ten thousand years. This is confirmed by the non-detection of extreme OH/IR-stars with a 10 M /yr ⇠ 10 superwind with age between 2 000 and 10 000 years. Under the assumption of single-star evolution (but ⇠ see main text), the detection of all observed extreme OH/IR-stars in that extremely short evolutionary phase 4 can statistically only be understood if the duration of the 10 M /yr superwind phase is not much longer ⇠ than 1 000 years (hence ⇡ approaches 1). However, this duration is too short for these stars to become ⇠ white dwarfs resulting in a conflict with initial-final mass relations deduced from young stellar clusters18.

ALMA observations of OH 26.5+0.6 and OH 30.1 0.7 Aiming to better quantify the short superwind phase of OH 26.5+0.6 and OH 30.1 0.7, low-excitation 12CO lines of both targets have been observed with ALMA (see Supplementary Information). The data were reduced using the CASA software package37. The full potential of the ALMA 12CO data is laid out in the channel maps (see Fig. 1, Fig. 2). Both the channel maps and zero-moment maps (see Supplementary Fig. 1, Supplementary Fig. 2 show bright emission in the central region with full width half maximum (FWHM) of 0. 4–0. 6. For a spherically symmetric wind, this corresponds to a duration of 130–200 ( 440– ⇠ 00 00 ⇠ ⇠ 650) years for OH 26.5+0.6 (OH 30.1 0.7) being similar to the superwind age derived for both targets (see Supplementary Information).

Of particular interest for this study is the emergence of an incomplete ring-like pattern in the 12CO channel map of both targets (see Fig. 1, Fig. 2, see Supplementary Information). This ring-like pattern signals the presence of a binary companion19,20, 34,38, 39. Position-velocity (PV) diagrams allow one to gauge if and how density structures are correlated. The PV diagrams prove that the circumstellar envelope of both targets is shaped by a binary-induced shell-like spiral19,20 created by the mass-losing AGB star wobbling around the center-of-mass of a binary system (see Fig. 3, see Supplementary Information). In addition, the high signal-to-noise data of OH 30.1 0.7 allow us to detects knots (see white arrows in Fig. 2) which indicate the region where the shell-like spiral pattern overlaps with the density wake caused by Bondi-Hoyle accretion of the binary companion (see Supplementary Information).

Binary properties of OH 26.5+0.6 and OH 30.1 0.7 Hydrodynamical simulations show that for a mass-losing star in a binary system wobbling around the center- of-mass in a circular orbit an Archimedean spiral is formed in the circumstellar envelope19. If the orbital plane is located exactly on the plane of the sky (i =0), this Archimedean spiral can easily be recognised in an angle-radius plot as straight lines (see upper panel in Supplementary Fig. 3). Including the effect of inclination, a spiral-shell structure exhibits undulations with respect to these straight lines34. This is illustrated for the OH 26.5+0.6 binary system in the bottom panel of Supplementary Fig. 3. As for CIT 634, the inclination effect of the OH 26.5+0.6 binary system can easily be discerned in Supplementary Fig. 3. Exploiting the asymmetry in the PV-diagrams and the presence of knots in the OH 30.1 0.7 data, one can use the outcome of radiative transfer calculations20,34 to estimate the inclination angle, i, of the binary

11 system which is between 60 – 75 for OH 26.5+0.6, and between 50 – 70 for OH 30.1 0.7. ⇠ ⇠ Assuming a circular orbit, the orbital period can be estimated from the slope of the observed emission in the angle-radius plot40, which for OH 26.5+0.6 yields 430 years. Another method exploits the spiral- ⇠ 39,41 arm spacing . The spiral-arm spacing is determined by the product of the binary orbital period Tp and the pattern propagation speed in the orbital plane, which is given by41

2 2⇡rp rarm = Vw + Vp , (1) h i 3 ⇥ Vp ⇣ ⌘ with V the wind velocity, V the orbital velocity, r the orbital radius of the primary, and the orbital period h wi p p T given by the last term (2⇡r /V ). The ﬁrst term in the right-hand part of Eq. 1, V + 2 V , is the pattern p p p h wi 3 p propagation speed throughout the orbital plane. In AGB binary simulations, this pattern speed is close to the 19 wind speed . At the distance where the spiral arcs are detected (beyond 100. 5 for OH 26.5+0.6 and 100. 0 for OH 30.1 0.7, corresponding to 500 and 1 000 stellar radii, respectively), the wind has already reached its ⇠ terminal velocity. The derived arm-spacing of 1 for OH 26.5+0.6 and of 0. 58 for OH 30.1 0.7 yields ⇠ 00 ⇠ 00 an orbital period of 350–430 years for OH 26.5+0.6 and of 460–590 years for OH 30.1 0.7, assuming ⇠ ⇠ orbital speeds ranging up to 5 km/s. Orbital periods for other (carbon-rich) AGB binary stars displaying spiral structures range from 55 – 830 years34,38, 39,42, 43.

Isotopic ratios of carbon and oxygen indicate a progenitor mass for both targets between 5 to 8 M 44 (see Supplementary Information). The current ALMA data do not allow an accurate estimate of the companion’s mass. Assuming a total mass for the binary system between 5 to 9 M and using Kepler’s third law, the mean binary separation is 100 20 astronomical units for OH 26.5+0.6 and 123 22 astronomical units for ± ± OH 30.1 0.7. For a stellar radius of 900 solar radii (OH 26.5+0.6) and 700 solar radii (OH 30.1 0.7)45, ⇠ ⇠ this corresponds to 24 stellar radii and 38 stellar radii, respectively. ⇠ ⇠ Currently, we know of some dozen extreme OH/IR stars within 4 000 parsec having similar mass- ⇠ loss rate characteristics as OH 26.5+0.6 and OH 30.1 0.7. Statistical reasoning implies a general argument that a large fraction of them is in a binary system. Using the space densities of Ortiz et al.36, we expect some 400 stars within a 4 000 parsec sphere having infrared colours similar to OH 26.5+0.6 and OH 30.1 0.7. ⇠ The frequency for binary companions with a mass ratio of 0.1 to 1.0 and an orbital period ranging from 100 -– 1000 years is 13% for stars with mass between 2 and 5 M and is 17% for stars between 5 and ⇠ ⇠ ⇠ 9M 46. Only considering mass ratios between 0.1 and 0.3, the binary frequency reduces to 6% and 9%, ⇠ ⇠ respectively. As such, we expect more extreme OH/IR stars to be detected being part of a binary system.

Properties of the compact equatorial density enhancement (CEDE)

The seminal work of Mastrodemos & Morris21 is one of the few studies exploring the effects of the hydrodynamical interaction in AGB binaries for a large parameter space of wind velocities, circular orbit separations, and companion masses. Although their simulations were not ﬁne-tuned toward the speciﬁc sit-

12 uation of OH 26.5+0.6 or OH 30.1 0.7 (i.e. the estimated mass of the OH/IR-stars is larger than the adopted mass of the primary of 1.5 solar masses), their results can be used to gauge the family of wind patterns that can arise. Of particular interest for this study is their model M7 which has a large binary period of 227 years, a binary separation of 24 stellar radii, and a wind velocity of 15 km/s, very similar to the deduced binary parameters of OH 26.5+0.6. The modelled wind pattern perpendicular to the orbital plane clearly shows the appearance of an incomplete ring pattern arising from a spiral shock (see Fig. 15 in Mastrodemos & Morris21). The arm-spacing is around 400 stellar radii, which in the case of OH 26.5+0.6 would translate to 1. 2, comparable to the one derived from the ALMA data. Of prime importance for this study is the emer- ⇠ 00 gence of a compact high-density region within 150 stellar radii (hence 6 times the binary separation) ⇠ ⇠ from the primary star caused by the gravitational focussing of the binary companion21.

A lower limit for the density enhancement of the compact equatorial density enhancement (CEDE) can be obtained using ballistic trajectory calculations47. Doing so, we have calculated the density ampliﬁ- cation factor deﬁned as the invert of the fraction of solid angle covered by the wind:

1/ sin (✓c)=1/ sin [arctan (H/r)] (2) where H is half of the extent of the wind in the direction normal to the orbital plane, r is the distance at which H is measured and ✓c is the angle subtended (see top panel in Supplementary Fig. 4). A small grid of simulations with binary parameters similar to the ones of OH 26.5+0.6 and OH 30.1 0.7 led to the estimates in the bottom panel of Supplementary Fig. 4, where the binary mass ratio has been set to 5 and the ﬁlling factor of the donor star to 10%. The variable represents how fast the wind reaches its terminal speed using a -type velocity law48: for =2(resp. =3), the wind reaches 90% of its terminal wind speed after 20 stellar radii (resp. 30 stellar radii). The beaming toward the equator is strongly dependent on ⌘, being the ratio of the terminal wind speed to the orbital speed, and on . These simulations show a density ampliﬁcation of a factor of a few, which is a lower limit since cooling, which is not accounted for, will compress the wind even more49.

The hydrodynamical simulations of Kim & Taam19 include all physics to better quantify the density enhancement factor. Using their simulation results, one can determine that the CEDE has a density contrast of a factor 10 with respect to the background wind density. This can be deduced by comparing Fig. 4 ⇠ and Fig. 10 of Kim & Taam19. Fig. 4 only includes the effect of the wobbling of the mass-loss AGB star around the center-of-mass (model M6), while in Fig. 10 both the wobbling of the mass-loss AGB star and the gravitational focusing by the secondary are accounted for (model M7). The density of model M7 is a linear superposition of the amounts of density enhancement due to the two mechanisms. The difference in density in the inner regions between Fig. 4 and Fig. 10 is roughly a factor of 10. The density enhancement increases for shorter radial distances and is strongly latitude-dependent. We note, however, that values between a factor of a few up to 100 are not unusual depending on the speciﬁc parameters of the binary ⇠ system19,21, 25,50, 51.

13 The bright centralized emission seen in the 12CO ALMA data of OH 26.5+0.6 and OH 30.1 0.7 signals the CEDE (Fig. 1, Fig. 2, Supplementary Video 1, Supplementary Video 2). The creation of a CEDE in OH 26.5+0.6 is also supported by the infrared image of the dusty envelope of OH 26.5+0.6 that displays a slightly oblate shape with a PA for the major axis being at 95 6 4 . The infrared asymmetry ± is also correlated with the asymmetry reported for the OH 1612 MHz maser for which a projected axis of rotation oriented in the north-south direction was deduced52,53. As such, both the ALMA PV diagrams of OH 26.5+0.6 and the oblateness of the OH and infrared dust images show morphological evidence for the north-south direction corresponding to the lower density orbital axis and the east-west direction to the higher density orbital plane (see Fig. 4). The direction of the orbital plane of OH 30.1 0.7 is more difﬁcult to constrain due to a lack of complementary observations. At the other hand, the presence of knots in the channel map (Fig. 2) is a strong diagnostic. In the PV-diagram at a PA of 165, being along the line of nodes of the knots, the knotty structure is present at the systemic velocity (see Fig. 3), conﬁrming the existence of the companion’s wake in the orbital plane34.

This CEDE also offers the perfect conditions for amorphous grains to be annealed into crystalline minerals. As shown in the simulations of Edgar et al.25, the local deflection of the wind flow by the secondary drives a shock formation lying in the orbital plane with an associated increase in temperature. A first estimate of the post-shock temperature can be calculated using Eq. 2 of Edgar et al.25. For the OH 26.5+0.6 and OH 30.1 0.7 binary system parameters derived above, it is straightforward to calculate that post-shock temperatures between 1067 K (the annealing temperature of silicates54) and 2000 K (when grains are vapor- ized) are reached for the primary mass being between 6.3 M and 8.8 M , in very good agreement with the mass estimates from the isotope analysis44. The associated timescale for annealing is so short compared to other timescales25 that the annealing is almost instantaneous. Hence, the shaping of the AGB wind by a binary companion provides a mechanism for forming crystalline dust in the orbital plane.

Derived mass-loss rates for extreme OH/IR-stars

While mass-loss rates for OH/IR-stars are sometimes derived from the OH maser intensity55,56, the un- certainties on the derived values is large due to the inherent non-linear behaviour of masers57. Therefore, the methods most frequently used for mass-loss rate determinations are the analyses of (i) infrared colours and/or dust spectral features or (ii) rotational transitions of CO gas. While both methods often suffer from the same simplification of a 1D (spherically symmetric) geometry, the first method involves the unknown dust-to-gas ratio and the unknown dust outflow velocities. The latter are often assumed to equal the gas outflow velocities, which is a valid assumption in the case of gas and dust being well coupled. For the context of this paper, one needs to assess how the binary-induced morphology of the circumstellar wind, in particular the spiral structure and the compact equatorial enhancement (CEDE), complicates the derivation of the mass-loss rate.

14 Tracing the mass-loss rate: warm dust versus cool gas: The mass-loss rates derived from the analyses of dust spectral features (or infrared colours) and from low-excitation CO lines yield significantly different results for extreme OH/IR-stars with the former ones being one or two orders of magnitude larger than the values derived from CO gas (see Supplementary Information). For OH 26.5+0.6 and OH 30.1 0.7, the 4 6 2, 9, 10,45, 58 values are 2–3 10 M /yr from dust versus 1 – 9.7 10 M /yr from CO , with the single- ⇥ ⇠ ⇥5 scattering radiation pressure limit being around 5 10 M /yr. This discrepancy was historically inter- ⇥ preted as a sign of a short intense superwind phase in which the (warm) dust traces the recent high density material close to the star, while the (cool) gas signals a lower mass-loss rate in the past (see Supplementary Information). However, realising that OH 26.5+0.6 and OH 30.1 0.7 are part of a binary system in which both a compact equatorial density enhancement (CEDE) and spiral structure are formed puts this discrepancy in another context. I.e., the (near-)infrared spectral energy distribution (SED) mainly traces the warm (crystalline) dust residing close to the star, hence in the CEDE. The density in the CEDE is roughly a factor 10 higher than the background wind density. Therefore, the reported high mass-loss rate superwind values are mainly a reflection of the higher density in the CEDE created by binary interaction, but not of the actual mass-loss rate which will be lower. An additional complication is that SED fits for extreme OH/IR-stars have always been based on a simplified 1D geometry. This assumption might lead to an erroneous derivation of the density structure (see Supplementary Information). Moreover, with the velocity vector field being highly complex in the inner wind region due to the binary interaction 21, a simple relation between density structure and mass-loss rate is ruled out.

The real mass-loss rate is better estimated from radiative transfer modelling of various low-excitation CO rotational lines observed with telescopes having a large beam size (hence incorporating the total wind extent). The advantage of low-excitation CO lines is that they probe the bulk of the material in the cooler less dense region of a stellar wind. As such, the mass in the CEDE is negligible compared to the mass probed. One might wonder if (i.) the occurrence of a large spiral structure, (ii.) the change in mass-loss rate during the AGB evolution, or (iii.) CO line saturation effects do not complicate the mass-loss rate retrieval from low-excitation CO gas. Concerning the ﬁrst aspect, we can rely on the studies performed by Homan et al.20. They have shown (see their Fig. 16, upper left panel) that one can safely analyse low-excitation CO emission lines using a simpliﬁed 1D approach in the case of a high mass-loss rate wind in which a shell-like spiral is embedded. To constrain the second aspect, we turn our attention again to OH 26.5+0.6 and OH 30.1 0.7, since both targets have been studied with various instruments. For OH 26.5+0.6, the CO J=1–0 line was imaged with BIMA and a CO envelope diameter of 8.5 5.5 (FWHM) was derived59. For a distance of 00⇥ 00 1370 parsec and a wind velocity of 15 km/s (see Supplementary Information), this translates to 2600 years, ⇠ which is 25% of the interpulse period for a star with initial mass of 5 M 35. Analogously, the 12CO J=1–0 ⇠ data of OH 30.1 0.7 yield an extent of only 3780 years (see Supplementary Information). For stars with ⇠ initial mass above 1.5 M , the mass-loss rate at the high-luminosity tip of the AGB will only marginally ⇠ change during that time interval28, and hence the assumption of a constant mass-loss rate when analysing a set of low-excitation CO lines is valid. The third aspect occurs for stars experiencing a high mass-loss

15 rate, since the saturation of the optically thick CO lines weakens the dependence of the line intensity on the 60 5 amount of CO in the envelope . This effect starts to play a role for mass-loss rates above 3 10 M /yr ⇠ ⇥ (i.e., similar to the single-scattering pressure limit for extreme OH/IR-stars), and the application of analytical formulae relating the CO line intensity with mass-loss rate should be done with care45. For mass-loss rates above that limit, dedicated radiative transfer modelling can cope with that effect61. Hence, none of the three aforementioned aspects impedes the use of low-excitation CO lines to derive the mass-loss rate of (extreme) OH/IR-stars. Additional advantages of the use of low-excitation CO lines are that the dust-to-gas conversion ratio is not needed and that the width of the CO lines yields the overall wind expansion velocity since the spiral-pattern propagation speed (Eq. 1) is close to the wind speed19. Therefore, the equation of mass- conservation, M˙ = 4⇡r2⇢(r)v(r), can be used to convert the derived gas density structure, ⇢(r), into a value for the mass-loss rate. Hence, we conclude that various low-excitation CO lines should be used to determine 5 2, 9, 58 the mass-loss rate in (extreme) OH/IR-stars yielding typical values around (0.1–3) 10 M /yr , i.e. ⇥ not exceeding the single-scattering radiation pressure limit, and potentially even remaining lower by a factor of a few.

The uncertainty on mass-loss rate values as derived from low-excitation CO lines is typically a factor of 360, which is partly caused by the impossibility to have direct knowledge on the [CO/H ]-ratio. The ⇠ 2 latter value is taken from AGB stellar nucleosynthesis models which predict a value around 3 10 4 with ⇥ an uncertainty of a factor of 215,62, 63. It might be expected that the mass-loss rate varies during thermal pulse cycles. Following the argumentation given in the Supplementary Information based on dynamical atmosphere models28, we expect that as the star evolves and the mass-loss rate reaches the single-scattering radiation pressure limit, the impact of thermal pulses or pulsations on variations in the density structure is modest64, being a factor 2. Taking this into account and realizing that the single-scattering limit for ⇠ 5 extreme OH/IR stars is 5 10 M /yr, this study puts forward a maximum mas-loss rate during the ⇠ ⇥ AGB-phase around the single-scattering radiation pressure limit.

Data availability The ALMA data from proposals 2015.1.00054.S, 2016.1.00005.S, and 2016.2.00088.S can be retrieved from the ALMA data archive at http://almascience.eso.org/aq/. The data that support the plots within this paper and other ﬁndings of this study are available from the corresponding author upon reasonable request.

Code availability Figs. 1–3, and Supplementary Figs. 1–2 are made using CASA37. Supplementary Fig. 3 is made using a simple python script that can be distributed upon request. Supplementary Fig. 4 is based on ballistic trajectory calculations by one of the co-authors47 solving the equation of motion using a classical fourth order Runge-Kutta scheme. The code is available for collaboration with I.E.M. upon reasonable request.

16 References

35. Karakas, A. I. Helium enrichment and carbon-star production in metal-rich populations. MNRAS 445, 347–358 (2014).

36. Ortiz, R. & Maciel, W. J. AGB stars: densities and formation rates obtained from OH/IR stars. A&A 313, 180–190 (1996).

37. McMullin, J. P., Waters, B., Schiebel, D., Young, W. & Golap, K. CASA Architecture and Applications. In Shaw, R. A., Hill, F. & Bell, D. J. (eds.) Astronomical Data Analysis Software and Systems XVI, vol. 376 of Astronomical Society of the Paciﬁc Conference Series, 127 (2007).

38. Maercker, M. et al. Unexpectedly large mass loss during the thermal pulse cycle of the red giant star R Sculptoris. Nat 490, 232–234 (2012).

39. Decin, L. et al. ALMA data suggest the presence of spiral structure in the inner wind of CW Leonis. A&A 574, A5 (2015).

40. Kim, H. et al. The large-scale nebular pattern of a superwind binary in an eccentric orbit. Nature Astronomy 1, 60 (2017).

41. Kim, H. & Taam, R. E. A New Method of Determining the Characteristics of Evolved Binary Systems Revealed in the Observed Circumstellar Patterns: Application to AFGL 3068. ApJ 759, L22 (2012).

42. Mauron, N. & Huggins, P. J. Imaging the circumstellar envelopes of AGB stars. A&A 452, 257–268 (2006).

43. Dinh-V.-Trung & Lim, J. Tracing the Asymmetry in the Envelope Around the Carbon Star CIT 6. ApJ 701, 292–297 (2009).

44. Justtanont, K. et al. Herschel observations of extreme OH/IR stars. The isotopic ratios of oxygen as a sign-post for the stellar mass. A&A 578, A115 (2015).

45. De Beck, E. et al. Probing the mass-loss history of AGB and red supergiant stars from CO rotational line proﬁles. II. CO line survey of evolved stars: derivation of mass-loss rate formulae. A&A 523, A18 (2010).

46. Moe, M. & Di Stefano, R. Mind Your Ps and Qs: The Interrelation between Period (P) and Mass-ratio (Q) Distributions of Binary Stars. ApJS 230, 15 (2017). 1606.05347.

47. El Mellah, I. & Casse, F. A numerical investigation of wind accretion in persistent supergiant X-ray binaries - I. Structure of the ﬂow at the orbital scale. MNRAS 467, 2585–2593 (2017). 1609.01532.

48. Lamers, H. J. G. L. M. & Cassinelli, J. P. Introduction to Stellar Winds (1999).

17 49. El Mellah, I., Sander, A. A. C., Sundqvist, J. O. & Keppens, R. Formation of wind-captured discs in Supergiant X-ray binaries : consequences for Vela X-1 and Cygnus X-1. ArXiv e-prints (2018).

50. Liu, Z.-W., Stancliffe, R. J., Abate, C. & Matrozis, E. Three-dimensional Hydrodynamical Simulations of Mass Transfer in Binary Systems by a Free Wind. ApJ 846, 117 (2017).

51. Chen, Z., Frank, A., Blackman, E. G., Nordhaus, J. & Carroll-Nellenback, J. Mass transfer and disc formation in AGB binary systems. MNRAS 468, 4465–4477 (2017).

52. Bowers, P. F. & Johnston, K. J. Sensitive VLA observations of OH 127.8 - 0.0 and OH 26.5 + 0.6. ApJ 354, 676–686 (1990).

53. Etoka, S. & Diamond, P. J. Insight into the OH polarimetric structure of OH26.5+0.6. MNRAS 406, 2218–2234 (2010).

54. Hallenbeck, S. L., Nuth, J. A., III & Nelson, R. N. Evolving Optical Properties of Annealing Silicate Grains: From Amorphous Condensate to Crystalline Mineral. ApJ 535, 247–255 (2000).

55. Baud, B. & Habing, H. J. The maser strength of OH/IR stars, evolution of mass loss and the creation of a superwind. A&A 127, 73–83 (1983).

56. van der Veen, W. E. C. J. & Rugers, M. A comparison between CO-, OH-, and IR-mass-loss rates of evolved stars. A&A 226, 183–202 (1989).

57. Marshall, J. R. et al. Asymptotic giant branch superwind speed at low metallicity. MNRAS 355, 1348– 1360 (2004).

58. Justtanont, K., Skinner, C. J., Tielens, A. G. G. M., Meixner, M. & Baas, F. Modeling of the Dust and Gas Outﬂows from OH 26.5+0.6: The Superwind. ApJ 456, 337 (1996).

59. Fong, D., Justtanont, K., Meixner, M. & Campbell, M. T. Imaging the circumstellar envelope of OH 26.5+0.6. A&A 396, 581–587 (2002).

60. Ramstedt, S., Schoier,¨ F. L., Olofsson, H. & Lundgren, A. A. On the reliability of mass-loss-rate estimates for AGB stars. A&A 487, 645–657 (2008).

61. Decin, L. et al. Probing the mass-loss history of AGB and red supergiant stars from CO rotational line proﬁles. I. Theoretical model - Mass-loss history unravelled in VY CMa. A&A 456, 549–563 (2006).

62. Karakas, A. I. Updated stellar yields from asymptotic giant branch models. MNRAS 403, 1413–1425 (2010). 0912.2142.

63. Cristallo, S., Straniero, O., Piersanti, L. & Gobrecht, D. Evolution, Nucleosynthesis, and Yields of AGB Stars at Different Metallicities. III. Intermediate-mass Models, Revised Low-mass Models, and the ph-FRUITY Interface. ApJS 219, 40 (2015). 1507.07338.

18 64. Vassiliadis, E. & Wood, P. R. Evolution of low- and intermediate-mass stars to the end of the asymptotic giant branch with mass loss. ApJ 413, 641–657 (1993).

19 20 21 PA=-5◦ PA=85◦

40 0.18

0.16 30 0.14 [ 0.12 Jy /beam] 20 0.1

LSRK radio velocity [km/s] velocity LSRKradio 0.08

10 0.06 0.04 0.02 0 0 -4 -3 -2 -1 0 1 2 3 4 -4 -3 -2 -1 0 1 2 3 4 Offset [arcsec] Offset [arcsec]

22 cf. preferred axis of •5 ring-like arcs PA = -5° Chesneau et al. 2005 detected by ALMA and Bowers & Johnston 1990 •due to reﬂex motion of AGB primary star for inclination i = 60°- 75° •arm spacing ≃ 1”

~ 1” ≃ 350R* binary separation

~ 0.073” ≃ 24R*

gravitational ~ 0.3” ≃ 100R* focus of orbital plane matter by “short companion superwind phase” detected in CO, SiO = • (Justtanont et al. 2013, 2018) compact equatorial •detected in dust density (Chesneau et al. 2005) enhancement •FWHM radius ≃ 0.3” ≃ 100R* orbital axis

23 Reduced maximum mass-loss rate of OH/IR stars due to overlooked binary interaction

Supplementary Information

1 The superwind scenario

The evidence of mass loss during the late stages of evolution has been known for a long time1,2. In 3 13 1975, Reimers published his famous mass-loss law for red giants, M˙ = 4 10 ⌘L/gR [M /yr], ⇥ with 1/3 <⌘<3 and the stellar luminosity L, stellar gravity g, and and stellar radius R in solar units. This formula was a ﬁt of mass-loss rate versus L/gR using an assortment of stars for which mass-loss rate estimates were available in 1975. The prescription yields quite good mass-loss rates for stars on the Red Giant Branch (RGB) for which mass loss can arise from chromospheric activity or can also be pulsation-enhanced and dust-driven4. However, for the Asymptotic Giant Branch phase, it yields mass-loss rates which increase too slowly over too long a period of time and the 6 5,6 maximum mass-loss rate achieved is only 10 M /yr . With typical AGB lifetimes of about 1 ⇠ million years, the cumulative mass loss is too low to explain the distribution of white dwarf masses and a Reimers-like wind cannot produce the characteristics of the planetary-nebulae population5. To resolve these issues, Renzini pointed out in 1981 that a much stronger wind must develop quite 5 5 rapidly at the end of the AGB evolution with mass-loss rates above a few 10 M /yr . Renzini suggested the provisional term ‘superwind’ for this phenomenon since the physical cause for it was not known at that time.

We note that this historical terminology of the word ‘superwind’ is not always followed in the literature. I.e. some authors use the word ‘superwind’ to indicate the phase when the rate associated with the mass-loss phenomenon exceeds the nuclear burning rate, such that mass loss determines the further evolution7, 8. For the latter deﬁnition, the transition to a superwind would happen around 7 10 M /yr. In this paper, we follow the historical deﬁnition of the word ‘superwind’ proposed by Renzini.

Following the suggestion of Renzini, speculations about the physical nature of the superwind have been made by many authors among whom we recall Baud & Habing (1983)9, Bedijn (1988)10, van der Veen (1989)11, and Bowen & Willson (1991)6. Growing evidence arising from dynamical modelling support the idea that as the star evolves, its mass decreases and radius increases yielding an increase in the atmospheric density scale-height. Since the density scale-height appears in the argument of an exponential function describing the density, even a moderate increase in density scale-height brings forth a very rapidly increase in mass-loss rate6. These dynamical models exhibit the properties and behaviour expected for terminal AGB evolution with a superwind6.

Even so, the debate over a sudden onset of the superwind phase or a gradual increase in mass-loss rate has not yet been settled, neither has the maximum mass-loss rate been established. The detection of a detached shell, with a density contrast of some 2 orders of magnitude, in some lower mass ( 1–3 solar masses) carbon-rich AGB stars is sometimes linked to a superwind burst ⇠ triggered by a thermal pulse12 occurring every 10 000 to 100 000 years and lasting a few hundred ⇠ years13. However, it then remains unclear why no analogous shells are detected around the more massive oxygen-rich counterparts. Dynamical atmosphere models14 support the idea of a more gradual increase in mass-loss rate for stars with an initial mass above 1.3 M . As the star evolves and the mass-loss rate reaches the single-scattering radiation pressure limit, the impact of thermal pulses or pulsations on variations in the density structure is predicted to be modest15 (see also Sect 7).

2 OH/IR-stars - The case of OH 26.5+0.6 and OH 30.1 0.7

For oxygen-rich AGB stars experiencing a substantial mass loss, the OH maser phenomenon occurs both in optically thick envelopes at the frequency of 1612 MHz (type II OH masers) and in optically thin envelopes at frequencies of 1665 and 1667 MHz (type I OH masers)16. A subclass of the type II OH/IR-stars is detected as infrared point source without optical counterpart9. These extreme OH/IR-stars represent a class of intermediate-mass oxygen-rich stars at the tip of the AGB evolution with the highest AGB mass-loss rates discovered. They currently experience an ex- 5 treme superwind phase (see Sect. 1 for terminology) with mass-loss rates between 10 and a few 4 10 M /yr, i.e. often exceeding the single-scattering radiation pressure limit (M˙ =L/vc, with v the 5 wind velocity, c the speed of light, and typical values around a few 10 M /yr) by a factor of a few. Their initial masses range between 2 to 8 solar masses, they exhibit bright OH masers, are large amplitude long-period (between 500–2000 days) variables, and have luminosities around 104 solar luminosities9,17. The characteristic absorption features at 10 and 18 micrometer by amorphous silicate grains conﬁrm that these objects are heavily obscured by dust and experience high mass-loss rates18. A subclass of these extreme OH/IR-stars even exhibit crystalline silicate features18,19.

OH 26.5+0.6 and OH 30.1 0.7 are two prototype targets of this class of extreme OH/IR- stars. OH 26.5+0.6 was discovered in 197420 and is known to have the strongest observed OH 1612 MHz maser emission among galactic AGB stars21. It is located at a distance of 1370 parsec22 and has a terminal wind velocity of 15 km/s23. The estimated current superwind mass-loss rate of 4 24 OH 26.5+0.6 is around 3 10 M /yr . Modelling of the gas and dust emission suggests that ⇥ the onset of the present-day high mass-loss rate (superwind) phase was only 130–200 years 6 19,24, 25⇠ ago and was preceded by a phase with mass-loss rate around 10 M /yr . The distance to OH 30.1 0.7 has recently been determined to be 3 900 parsec26. The current superwind phase 4 24 ⇠ has a mass-loss rate of 1.8 10 M /yr . The extent of the superwind envelope is estimated ⇠ ⇥ around 2.5 3.7 1016 cm 19,24, which for a wind velocity of 18.1 km/s24 translates to a duration ⇥ of 440–650 years. The observed carbon and oxygen isotopic ratios indicate that the progenitor ⇠ mass of both OH/IR-stars is between 5 to 8 solar masses27.

More than a dozen other extreme OH/IR-stars (within a distance of 4 000 parsec) have ⇠ been shown to display a similar behaviour with a sudden jump in mass-loss rate by one to two orders of magnitude19,23, 24,28–30. This behaviour was ﬁrst noted by Heske et al.28 who analysed a sample of 13 OH/IR-stars. For this sample, the mean mass-loss rate derived from low-excitation 6 CO lines was 7.4 10 M /yr, while the analysis of IRAS infrared data yielded a mean value 4 ⇥ of 1.1 10 M /yr. This discrepancy was historically interpreted as an indication that the mass- ⇥ loss rate is lower in the cool outer parts of the envelope (where CO is emitting) than in the warm inner parts (where dust is emitting)28. The derived duration of the current high mass-loss rate (‘superwind’) phase varies between 200–1 000 years19,23, 24,30. ⇠

3 ALMA observations of OH 26.5+0.6 and OH 30.1 0.7

Aiming to better quantify the short superwind phase of OH 26.5+0.6 and OH 30.1 0.7, several molecular lines have been observed with ALMA (proposals 2015.1.00054.S and 2016.2.00088.S for OH 26.5+0.6, and proposal 2016.1.00005.S for OH 30.1 0.7). Of particular interest are the low-excitation rotational transitions of 12CO, since CO is a well-known tracer of the density in stellar winds31,32. The interferometric capabilities of ALMA allow an excellent ﬁltering of the interstellar CO emission that otherwise complicates the study of CO in the extreme OH/IR-stars residing in the galactic plane.

3.1 ALMA 12CO J=3-2 data of OH 26.5+0.6 OH 26.5+0.6 was observed with ALMA in three different configurations including a compact configuration of the main array (12 m antennas), an extended configuration of the main array, and the

ALMA compact array (7 m antennas). The main array data sample the small scale structures (000.1 up to a few arcsec), while the compact array is sensitive to structures between 4–1000. The differ- 25 ⇠ ent datasets have been combined . The spatial resolution of the data is 000.22 000.18 with a velocity ⇥ resolution of 1 km/s. The noise level of the channel map, , is 8 mJy/beam, the maximum rms ⇠ intensity value reached at systemic velocity (of 28 km/s) is 0.52 Jy/beam. ⇠ The zero-moment map is shown in Supplementary Fig. 1. The full extent of the zero-moment map is 300.7 300.2, which for a distance of 1370 parsec translates into an envelope with radius ⇠ ⇥ of 3.8 1016 cm. Interpreting this outer CO edge as a lower limit for the CO photodissociation ⇠ ⇥ 7 33,34 radius, this extent indicates a value for the mass-loss rate in the outer wind > 1 10 M /yr . ⇠ ⇥ The CO photodissociation radius is better constrained from the 12CO J=1-0 BIMA map having a 35 deconvolved full width at half maximum (FWHM) of 800.5 500.5 or an outer-wind mass-loss rate 7 33,34 ⇥ around 8 10 M /yr . ⇥

The channel map of the 12CO J = 3–2 emission is shown in Fig. ??. An animated video slicing through the velocity frames of the channel map is provided in Supplementary Video 1. Both the channel map and zero-moment map (Supplementary Fig. 1) show bright emission with a deconvolved FWHM for the zero-moment map of 000.41. Of particular interest for this study is the emergence of an incomplete ring-like pattern for velocities between 16 – 38 km/s at distances beyond 100.5 as measured from the central pixel (Fig. ??). This ring-like pattern signals the ⇠ presence of a binary companion causing OH 26.5+0.6 to wobble around the center-of-mass of the binary system (see Sect. 4). This reﬂex motion shapes the envelope surrounding OH 26.5+0.6 into a shell-like spiral pattern.

3.2 ALMA 12CO J=2-1 data of OH 30.1 0.7 The 12CO J=1–0 and J=2–1 rotational transitions imaging the envelope surrounding OH 30.1 0.7 12 were obtained with ALMA The ALMA restoring beam for the CO J=1–0 was 100.41 100.04, while 12 ⇥ a higher spatial resolution of 000.41 000.26 was achieved for the CO J=2–1 transition. The reduced ⇥ datasets are available at the ALMA archive. The noise level of the 12CO J=2–1 channel map, , is 1.7 mJy/beam, the maximum intensity value reached at systemic velocity (of 99 km/s) rms ⇠ ⇠ is 0.29 Jy/beam.

The zero-moment maps of both transitions are displayed in Supplementary Fig. 2. The full extent of the 12CO J=1–0 map is signiﬁcantly larger than for 12CO J=2–1, which is easily understood in terms of the higher excitation temperature for the 12CO J=2–1 line. The CO photodissociation 12 radius can be constrained from the CO J=1–0 map and yields a value of 300.7. At a distance ⇠ of 3 900 parsec, this translates to an envelope radius of 2 1017 cm which marks a mass-loss rate 6 33,34 ⇥ in the outer wind of 5 10 M /yr . The centralized bright emission in the high-resolution 12 ⇠ ⇥ CO J=2–1 data has a deconvolved FWHM of 000.60. ⇠

The spatial resolution of the 12CO J=1–0 data is too low to show any morphological structure in the envelope surrounding OH 30.1 0.7, but various ring-like arcs are detected in the 12CO J = 2– 1 data (see Fig. ??, Supplementary Video 2). In analogy to OH 26.5+0.6, this morphology points toward binary interaction (see Sect. 4).

4 OH 26.5+0.6 and OH 30.1 0.7 as binary systems

PV diagrams of the ALMA 12CO data: A particular potency of channel maps is gauged by using position-velocity (PV) diagrams. A PV diagram is created by taking a slit with a specific length and width at a specific position angle (PA) through the central pixel (representing the position of the star). The PA is measured from North to East. The slit width is used for averaging the intensity (from the channel map) along the slice. The resulting 2D plot hence yields the average emission along the chosen axis versus velocity. The strength of this visualization technique is to detect correlated density structures and guide their interpretation36. For a spherically symmetric wind, no differences between PVs at different position angles (PAs) are expected36. In the case of asymmetric (but correlated) density structure, the key is to define two orthogonal PVs so as to maximally exploit the asymmetry of the data37. For the ALMA 12CO data of OH 26.5+0.6, this is achieved by taking a PA of 5 and 85, while for OH 30.1 0.7 the PA values of 75 and 165 are optimal (see Fig. ??). The resulting PV diagrams for both targets are characteristic of a binary-induced shell-like spiral38,39 (see below). An analogous morphology was described for the carbon-rich AGB star CIT 640.

Binary-induced spiral structures: Observations at increasingly higher spatial resolution have revealed the existence of asymmetries in the circumstellar envelopes of a fraction of nearby AGB stars, including spirals, arcs, and disks36,41–43. A binary companion star is often put forward as being responsible for the wind shaping. Pioneering hydrodynamical models simulating the influence of a companion on a wide orbit around an AGB star demonstrate that spiral patterns readily emerge44. If the companion is relatively massive, a spiral-shell-shaped pattern is created by the orbital motion of the AGB star around the centre of mass of the system38. The spiral windings are characterized by circular arcs in the meridional plane 38,44 (see Fig. ??). In addition, the companion’s gravity focusses a fraction of the wind material toward the equatorial plane, and a (more flattened) spiral structure attached to the companion might emerge caused by Bondi-Hoyle accretion38,45, 46. The former type of shell-like spirals are easily confused with periodical (spherically symmetric) increases in density due to periods of enhanced mass loss and only PV diagrams at various PAs can differentiate between the two causes. The differentiation between the shell-like versus flattened spiral structure is easily done in the PV-plane39. For the case of OH 26.5+0.6 and OH 30.1 0.7, the PV diagrams (Fig. ??) prove that the spiral structure detected has a shell-like pattern and hence reflects the indirect influence of the binary companion via the reflex motion of the AGB star. In addition, knots can be discerned in the channel map of OH 30.1 0.7 at a PA of 10 and 165 (see white arrows in Fig. ??). Hydrodynamical simulations show that these knots arise in the region of the gravitationally induced density wake with a flatter spiral shape overlapping with the spiral-shell pattern due to the orbital motion of the AGB star38,40. As in the case of CIT 640, the PV-diagram of

OH 30.1 0.7 is most symmetric along the line of nodes, hence a PA 165. ⇠

Hitherto, most spiral structures are detected around carbon-rich AGB stars36,40, 41,47, 48. Only recently, the ﬁrst spiral structure was reported around a low mass-loss rate oxygen-rich AGB star43. This paper reports the ﬁrst detection of a binary-induced spiral structure around intermediate- mass oxygen-rich AGB stars undergoing high mass loss. The observed fraction of main-sequence intermediate-mass stars being part of a binary system is at least 30 – 45%. For extreme OH/IR stars within initial mass between 5–8 M , the binary fraction is of the order of 60–100% at all separations, with the fraction for 5 M stars being closer to 60%49,50. Hence, this detection of binary- induced spiral patterns does not come as a surprise, but the capabilities of ALMA are needed to detect the relics of binary interaction in the wind pattern.

For a binary system the terminal wind velocity is expected to be lower since the global Roche potential to escape from is somewhat lower that the one produced by an isolated AGB (at ﬁxed AGB mass) thanks to the orbital rotation which brings some centrifugal support to the wind launching, provided the AGB star is synchronized.

5 1D versus 2D analysis of dust spectral features:

The occurrence of a compact equatorial density enhancement (CEDE), or in general a higher density along the orbital plane and a lower density along the orbital axis (see Fig. ??), might undermine the use of a 1D geometry when ﬁtting the infrared SED. As discussed in the main text and in the Supplementary Information an analogous morphology, but in a more evolved state, is seen in the AGB descendants, the (pre-)planetary nebulae. As such, we can use examples from that commu- nity to assess the error made when analysing the infrared SED of OH 26.5+0.6 and OH 30.1 0.7 (and with it other extreme OH/IR-stars in a binary system) using a spherically symmetric assumption. The post-AGB OH/IR-star HD 161796 has been analysed using both a 1D approach51 and a 2D parameterized description of a bipolar structure with an equatorial density enhancement52. Both studies used the same gas-to-dust ratio of 270. As a cautionary remark, we wish to add here that both studies have assumed the same constant velocity of 15 km/s to translate the derived density structure to a mass-loss rate. However, in the case of non-radial structures that assumption might lead to erroneous conclusions on the absolute values of the mass-loss rate (and its potential evolution through time) although a relative comparison between the derived density structures and total mass lost is still meaningful. While Hoogzaad et al.51 have derived a gas mass-loss rate of 4 51 5.1 10 M /yr which started 900 years ago and ended 430 years ago, Min et al. retrieved a ⇥ 5 mass-loss rate of 6.5 10 M /yr during the same period with a sharp increase in mass-loss rate ⇠ ⇥ in the next 100 years. The accumulated mass derived in the 1D study51 was 0.24 M and only ⇠ ⇠ 0.05 M in the 2D study51, a factor of 5 lower. ⇠ ⇠

In other cases, such as the water-fountain OH/IR post-AGB source IRAS 16342-3814, the derived dust mass from a 1D53 and 2D54 analysis is quite similar (in this case around 0.01 M ) but the 3 gas mass-loss rate differs largely and even reaches an exceptionally large value of 1 10 M /yr in ⇥ the 1D ﬁtting. This mass-loss rate is much higher than the value derived from the analysis of low- 5 55 excitation CO lines, being around 2.9 10 M /yr , conﬁrming the conclusion drawn in previous ⇥ section.

6 Application to mass-loss rate prescriptions in stellar evolution models

Stellar evolution models use parametric relations between the mass-loss rate and some fundamental stellar parameter(s) to calculate the change of the mass throughout stellar evolution. Various relations for AGB evolution exist, such as the ones proposed in refs. 15,56–62. Mass-loss rate prescriptions during the high mass-loss rate superwind phase (see Sect. 1 for terminology) differ by a factor of ten depending on, e.g., whether they adopt the single-scattering radiation pressure limit15 or the extreme OH/IR-stars56,57 to establish the maximum AGB mass-loss rate. Some of these formula are based on fitting models56,58, while others fit observationally derived mass-loss rates15,57, 59–62. These latter relations inherently depend on the method of mass-loss rate retrieval for the sample. The mass-loss rate for OH/IR-stars included in these analyses is always deduced from infrared colours or a fit to the dusty spectral energy distribution, with the exception being De Beck et al.61 who used low-rotational CO lines. Building upon the discussion in Methods, we conclude that the reported mass-loss rates for extreme OH/IR-stars based on dust spectral features are too high by a factor of a few up to 100. This conclusion impacts some mass-loss rate prescriptions56,57, 59,62 used in stellar ⇠ evolution models. We exemplify this statement using Fig. 1 of Vassiliadis & Wood15: for pulsation periods between 300–500 days the mass-loss rate scales with the period and for longer periods the mass-loss rate reaches a constant maximal value. All stars in the constant mass-loss rate part of the diagram of Vassiliadis & Wood15 are OH/IR-stars with deduced mass-loss rates well above the single-scattering radiation pressure limit (represented by the dotted line in Fig. 1 of Vassiliadis & Wood15). The dozen OH/IR-stars in their sample lie in the galactic plane or in the Large Magellanic Cloud (LMC). Remarkably, all five galactic OH/IR-stars belong to the group of extreme OH/IR- stars with crystalline silicate features (OH26.5+0.6, OH30.1-0.7, OH32.8-0.3, OH104.9+2.4, and OH127.8-0.0; see results of de Vries et al.19). OH maser studies show similar structural properties between these sources63. Building upon the results of OH 26.5+0.6 and OH 30.1 0.7, we argue that the mass-loss rate of these galactic OH/IR-stars is actually lower, i.e. the values displayed in Fig. 1 of Vassiliadis & Wood15 reflect the higher density in the orbital plane of a binary system in the vicinity of the OH/IR-star and not the real value of the mass-loss rate. The higher density in the orbital plane facilitates the crystallization of grains (see Methods). The real mass-loss rate is better estimated from low-excitation CO rotational lines (see previous section). Exploiting this method, 5 24,28, 61,64 the mass-loss rate is of the order of (0.1 3) 10 M /yr for these extreme OH/IR-stars , ⇥ around or slightly below the single-scattering radiation pressure limit. The same scenario seems to hold for the OH/IR-stars in the LMC. All OH/IR-stars in the LMC sample of Vassiliadis & Wood15 display the characteristic crystalline silicate features in their infrared spectra (see results of Jones et al.65 and Van Loon et al.66). An in-depth study by Van Loon57, using the dusty spectral energy distribution as a diagnostic tool, suggests that OH/IR-stars in the LMC might exceed the single- scattering radiation pressure limit by up to a factor 10. However, here as well do all OH/IR-stars ⇠ in their sample with mass-loss rate values significantly above the single-scattering radiation pressure limit and for which infrared spectra have been obtained show the characteristic crystalline silicate features (see results of Jones et al.65). As for the galactic counterparts, we suggest the real mass-loss rates of these targets to be lower, potentially by a factor ten or more. This proposition is supported by the fact that, unexpectedly, ALMA failed to detect the low-excitation rotational CO gas emission in the two LMC OH/IR-stars with highest estimated mass-loss rates by van Loon et 4 al.: a fit of the dusty spectral energy distribution yielded estimated values around 2 10 M /yr for ⇥ IRAS04545-7000 and IRAS05298-695757, while the non-detection of CO suggests mass-loss rates 6 67 lower than 4 10 M /yr for these two stars . Since both targets belong to the LMC, and hence ⇥ have a well-determined distance, the luminosity of these sources is well constrained, being 24 900 and 37 700 L , respectively. These high luminosity values indicate that both targets have reached the high-luminosity tip of the AGB13 while still experiencing a wind with mass-loss rate lower than the single-scattering radiation pressure limit. Based on this, we conclude that the maximal mass-loss rate of OH/IR-stars at the tip of the AGB evolution will be close to the single-scattering radiation pressure limit.

7 Implication of binary scenario on reported characteristics of wind dynamics, chemistry, and morphology

The idea of a binary-induced compact high-density region mimicking a short intense superwind phase solves several issues related to wind dynamics, chemistry, and morphology. First, for extreme OH/IR-stars a sudden jump in mass-loss rate (of 1 to 2 orders-of-magnitude) does not need to be invoked which is consistent with observed OH/IR samples having a continuous distribution of infrared colours68. Second, no imprints of thermal pulses or pulsations triggering strong episodic superwind bursts have been detected in these OH/IR-stars69. Higher density shells with periodic distances of a few hundred years can be explained in the context of a binary scenario. This casts aside the suggestion of several short but strong superwind phases proposed for these stars19 and lends support to the idea of a more gradual increase in mass-loss rate as the star evolves to the tip of the AGB. Third, a binary scenario provides an evolutionary link to the morphological appearance of the AGB descendants, the (pre-)planetary nebulae (PNe), which frequently display bipolar and multi-polar structures70,71. OH 26.5+0.6 and OH 30.1 0.7 could be the progenitors of a fraction of such (pre-)PNe (see below). In particular, the lower density along the orbital axis might develop into bipolar (or multi-polar) lobes, while the compact high-density region in the equatorial plane might evolve into a torus (or ring-like) structure with a central gap due to the photodissociation by the hot central PN star. This gap and torus will further collimate the outﬂow from the central star. This suggests that the equatorial density enhancement precedes the shaping of the bipolar nebulae and is not coeval, thus being the direct cause of the nebula shaping. Fourth, the equatorial density enhancement allows favourable conditions for dust grains to grow explaining the detection of micron-sized grains in some (pre-)PNe72,73. If densities are high enough, second generation exoplanets may even form, in analogy to the results for planet formation around young stellar objects74.

8 Morphological appearance of (pre-)PNe and the potential evolutionary link to OH 26.5+0.6 and OH 30.1 0.7

Both the reﬂex motion of the primary AGB star and the gravitational focusing by the companion result in higher densities along the equatorial plane and lower density towards the poles (see sketch in Fig. ??). These density properties can be seen as the right conditions for bipolar (or multipolar) and torus (or ring-like) structures to develop during the (pre-)PNe phase.

Indeed, about 25–50% of all OH/IR-stars do not show long-period pulsational variability. These stars are thought to be in the early pre-PNe phase75,76. Among several of them, bipolar outflows were found with, in part, spectacularly large outflow velocities (so-called ‘water fountains’) 71, 77, 78. Such outflows point to the presence of a binary system and OH 26.5+0.6 and OH 30.1 0.7 could be the progenitor of such pre-PNe. Within the next few thousand years the primary star will have lost most of its stellar envelope and will have stopped its pulsation. The central binary system may then trigger the launch of a bipolar outflow, while the star is evolving toward a PN.

In PNe, bipolar and multi-polar structures are frequently detected70. In 1995, it was already suggested that bipolar PN nebula are produced by more massive progenitors than is the case for other morphological PN types79. The case-study of OH 26.5+0.6 and OH 30.1 0.7 and a direct comparison with one of the best studied young PN, NGC 6302, support this hypothesis and moreover point toward the impact of binary action on the nebula shaping. OH 26.5+0.6, OH 30.1 0.7, and NGC 6302 have similar initial masses of about 5.5 solar masses and the infrared spectra of both targets display crystalline silicate features27,80. NGC 6302 is a large butterﬂy-like young PN, in which an expanding molecular torus (or ring) with mass of about 0.1 solar mass and a system of high-velocity bipolar outﬂows were detected81. Detailed analysis shows the equatorial torus to be ejected some 5 000 years ago during an event lasting for about 2 000 years81. As such, the mass-loss 5 rate linked to this event is of the order of 5 10 M /yr, which corresponds to the single scatter- ⇥ ing radiation-pressure limit for a star similar to OH 26.5+0.6 and OH 30.1 0.7. OH 26.5+0.6 and OH 30.1 0.7 are caught in action while creating a compact higher density region in the equatorial plane thanks to the direct graviational focussing by the companion. Once the star enters the PN phase, the hot central star will photodissociate its surroundings. Consequently, the morphology of the compact equatorial region will change into a torus (or disk-like) structure with central gap. Both the compact equatorial region during the AGB phase and the torus-like structure during the PN phase provide the appropriate conditions for grains to crystallize and the high densities in these regions facilitate the detections of these features in their infrared spectra82,83.

References

1. Adams, W. S. & MacCormack, E. Systematic Displacements of Lines in the Spectra of Certain Bright Stars. Astrophys. J. 81, 119 (1935).

2. Deutsch, A. J. The Circumstellar Envelope of Alpha Herculis. Astrophys. J. 123, 210 (1956).

3. Reimers, D. Circumstellar absorption lines and mass loss from red giants. Memoires of the Societe Royale des Sciences de Liege 8, 369–382 (1975).

4. Groenewegen, M. A. T. Infrared excess around nearby red giant branch stars and Reimers law. Astron. Astrophys. 540, A32 (2012).

5. Renzini, A. Red giants as precursors of planetary nebulae. In Iben, I., Jr. & Renzini, A. (eds.) Physical Processes in Red Giants, vol. 88 of Astrophysics and Space Science Library, 431–446 (1981).

6. Bowen, G. H. & Willson, L. A. From wind to superwind - The evolution of mass-loss rates for Mira models. Astrophys. J. Lett. 375, L53–L56 (1991).

7. Lagadec, E. & Zijlstra, A. A. The trigger of the asymptotic giant branch superwind: the importance of carbon. Mon. Not. Roy. Astron. Soc. 390, L59–L63 (2008). 0807.3730. 8. Zijlstra, A. A., Lagadec, E., Sloan, G. & Matsuura, M. The AGB Superwind in Nearby Galax- ies. In Luttermoser, D. G., Smith, B. J. & Stencel, R. E. (eds.) The Biggest, Baddest, Coolest Stars, vol. 412 of Astronomical Society of the Paciﬁc Conference Series, 65 (2009).

9. Baud, B. & Habing, H. J. The maser strength of OH/IR stars, evolution of mass loss and the creation of a superwind. Astron. Astrophys. 127, 73–83 (1983).

10. Bedijn, P. J. Pulsation, mass loss, and evolution of upper asymptotic giant branch stars. Astron. Astrophys. 205, 105–124 (1988).

11. van der Veen, W. E. C. J. The mass-loss evolution of oxygen-rich AGB stars and its consequences for stellar evolution. Astron. Astrophys. 210, 127–146 (1989).

12. Olofsson, H. et al. A high-resolution study of episodic mass loss from the carbon star TT Cygni. Astron. Astrophys. 353, 583–597 (2000).

13. Karakas, A. I. Helium enrichment and carbon-star production in metal-rich populations. Mon. Not. Roy. Astron. Soc. 445, 347–358 (2014).

14. Schroder,¨ K.-P., Winters, J. M. & Sedlmayr, E. Tip-AGB stellar evolution in the presence of a pulsating, dust-induced “superwind“. Astron. Astrophys. 349, 898–906 (1999).

15. Vassiliadis, E. & Wood, P. R. Evolution of low- and intermediate-mass stars to the end of the asymptotic giant branch with mass loss. Astrophys. J. 413, 641–657 (1993).

16. Ortiz, R. & Maciel, W. J. AGB stars: densities and formation rates obtained from OH/IR stars. Astron. Astrophys. 313, 180–190 (1996).

17. Engels, D., Kreysa, E., Schultz, G. V. & Sherwood, W. A. The nature of OH/IR stars. I - Infrared Mira variables. Astron. Astrophys. 124, 123–138 (1983).

18. Sylvester, R. J. et al. 2.4-197 mu m spectroscopy of OH/IR stars: the IR characteristics of circumstellar dust in O-rich environments. Astron. Astrophys. 352, 587–599 (1999).

19. de Vries, B. L. et al. The problematically short superwind of OH/IR stars. Probing the outﬂow with the 69 µm spectral band of forsterite. Astron. Astrophys. 561, A75 (2014).

20. Andersson, C., Johansson, L. E. B., Goss, W. M., Winnberg, A. & Nguyen-Quang-Rieu. OH 26.5+0.6 - a strong OH source at 1612 MHz. Astron. Astrophys. 30, 475–477 (1974).

21. Engels, D. & Bunzel, F. A database of circumstellar OH masers. Astron. Astrophys. 582, A68 (2015). 22. van Langevelde, H. J., van der Heiden, R. & van Schooneveld, C. Phase lags from multiple ﬂux curves of OH/IR stars. Astron. Astrophys. 239, 193–204 (1990).

23. Justtanont, K., Olofsson, G., Dijkstra, C. & Meyer, A. W. Near-infrared observations of water- ice in OH/IR stars. Astron. Astrophys. 450, 1051–1059 (2006).

24. Justtanont, K. et al. OH/IR stars and their superwinds as observed by the Herschel Space Observatory. Astron. Astrophys. 556, A101 (2013).

25. Justtanont, K. et al. ALMA observations of the extreme OH/IR star OH 26.5+0.6. I. The resolved molecular envelope of the superwind. Astron. Astrophys. (subm.) (2018).

26. Engels, D., Etoka, S., Gerard,´ E. & Richards, A. Phase-lag Distances of OH Masing AGB Stars. In Kerschbaum, F., Wing, R. F. & Hron, J. (eds.) Why Galaxies Care about AGB Stars III: A Closer Look in Space and Time, vol. 497 of Astronomical Society of the Paciﬁc Conference Series, 473 (2015).

27. Justtanont, K. et al. Herschel observations of extreme OH/IR stars. The isotopic ratios of oxygen as a sign-post for the stellar mass. Astron. Astrophys. 578, A115 (2015).

28. Heske, A., Forveille, T., Omont, A., van der Veen, W. E. C. J. & Habing, H. J. Deﬁciency of CO emission from massive envelopes around cool OH/IR stars. Astron. Astrophys. 239, 173–185 (1990).

29. Delfosse, X., Kahane, C. & Forveille, T. Superwind in evolved OH/IR stars. Astron. Astrophys. 320, 249–256 (1997).

30. Chesneau, O. et al. The mid-IR spatially resolved environment of OH 26.5+0.6 at maximum luminosity. Astron. Astrophys. 435, 563–574 (2005).

31. Morris, M., Lucas, R. & Omont, A. Molecular emission from expanding circumstellar envelopes Polarization and proﬁle asymmetries. Astron. Astrophys. 142, 107–116 (1985).

32. Schonberg,¨ K. Molecular line spectra from circumstellar envelopes. I - Excitation of 12CO, 13CO, and CS, and the formation of line proﬁles. Astron. Astrophys. 195, 198–207 (1988).

33. Groenewegen, M. A. T. The photodissociation of CO in circumstellar envelopes. Astron. Astrophys. 606, A67 (2017).

34. Mamon, G. A., Glassgold, A. E. & Huggins, P. J. The photodissociation of CO in circumstellar envelopes. Astrophys. J. 328, 797–808 (1988).

35. Fong, D., Justtanont, K., Meixner, M. & Campbell, M. T. Imaging the circumstellar envelope of OH 26.5+0.6. Astron. Astrophys. 396, 581–587 (2002). 36. Decin, L. et al. ALMA data suggest the presence of spiral structure in the inner wind of CW Leonis. Astron. Astrophys. 574, A5 (2015).

37. Homan, W., Boulangier, J., Decin, L. & de Koter, A. Simpliﬁed models of circumstellar mor- phologies for interpreting high-resolution data. Analytical approach to the equatorial density enhancement. Astron. Astrophys. 596, A91 (2016).

38. Kim, H. & Taam, R. E. Wide Binary Effects on Asymmetries in Asymptotic Giant Branch Circumstellar Envelopes. Astrophys. J. 759, 59 (2012).

39. Homan, W. et al. Simpliﬁed models of stellar wind anatomy for interpreting high-resolution data. Analytical approach to embedded spiral geometries. Astron. Astrophys. 579, A118 (2015).

40. Kim, H., Hsieh, I.-T., Liu, S.-Y. & Taam, R. E. Evidence of a Binary-induced Spiral from an Incomplete Ring Pattern of CIT 6. Astrophys. J. 776, 86 (2013).

41. Maercker, M. et al. Unexpectedly large mass loss during the thermal pulse cycle of the red giant star R Sculptoris. Nature 490, 232–234 (2012).

42. Kervella, P. et al. ALMA observations of the nearby AGB star L2 Puppis. I. Mass of the central star and detection of a candidate planet. Astron. Astrophys. 596, A92 (2016).

43. Homan, W., Richards, A., Decin, L., de Koter, A. & Kervella, P. An unusual face-on spiral in the wind of the M-type AGB star EP Aquarii. Astron. Astrophys. 614, A113 (2018).

44. Mastrodemos, N. & Morris, M. Bipolar Pre-Planetary Nebulae: Hydrodynamics of Dusty Winds in Binary Systems. II. Morphology of the Circumstellar Envelopes. Astrophys. J. 523, 357–380 (1999).

45. Liu, Z.-W., Stancliffe, R. J., Abate, C. & Matrozis, E. Three-dimensional Hydrodynamical Simulations of Mass Transfer in Binary Systems by a Free Wind. Astrophys. J. 846, 117 (2017).

46. Chen, Z., Frank, A., Blackman, E. G., Nordhaus, J. & Carroll-Nellenback, J. Mass transfer and disc formation in AGB binary systems. Mon. Not. Roy. Astron. Soc. 468, 4465–4477 (2017).

47. Mauron, N. & Huggins, P. J. Imaging the circumstellar envelopes of AGB stars. Astron. Astrophys. 452, 257–268 (2006).

48. Dinh-V.-Trung & Lim, J. Tracing the Asymmetry in the Envelope Around the Carbon Star CIT 6. Astrophys. J. 701, 292–297 (2009). 49. Raghavan, D. et al. A Survey of Stellar Families: Multiplicity of Solar-type Stars. Astrophys. J. Suppl. 190, 1–42 (2010). 1007.0414.

50. Duchene,ˆ G. & Kraus, A. Stellar Multiplicity. Ann. Rev. Astron. Astrophys. 51, 269–310 (2013).

51. Hoogzaad, S. N. et al. The circumstellar dust shell of the post-AGB star HD 161796. Astron. Astrophys. 389, 547–555 (2002).

52. Min, M. et al. The color dependent morphology of the post-AGB star HD 161796. Astron. Astrophys. 554, A15 (2013).

53. Dijkstra, C. et al. The mineralogy, geometry and mass-loss history of IRAS 16342-3814. Astron. Astrophys. 399, 1037–1046 (2003).

54. Murakawa, K. & Izumiura, H. Dust shell model of the water fountain source IRAS 16342- 3814. Astron. Astrophys. 544, A58 (2012).

55. Imai, H., He, J.-H., Nakashima, J.-I. U., Nobuharu, Deguchi, S. & Koning, N. CO J = 3-2 Emission from the “Water Fountain” Sources IRAS 16342-3814 and IRAS 18286-0959. Pub. Astron. Soc. Japan 61, 1365–1372 (2009).

56. Bloecker, T. Stellar evolution of low and intermediate-mass stars. I. Mass loss on the AGB and its consequences for stellar evolution. Astron. Astrophys. 297, 727 (1995).

57. van Loon, J. T. et al. Mass-loss rates and luminosity functions of dust-enshrouded AGB stars and red supergiants in the LMC. Astron. Astrophys. 351, 559–572 (1999).

58. Willson, L. A. Mass Loss From Cool Stars: Impact on the Evolution of Stars and Stellar Populations. Ann. Rev. Astron. Astrophys. 38, 573–611 (2000).

59. van Loon, J. T., Cioni, M.-R. L., Zijlstra, A. A. & Loup, C. An empirical formula for the mass-loss rates of dust-enshrouded red supergiants and oxygen-rich Asymptotic Giant Branch stars. Astron. Astrophys. 438, 273–289 (2005).

60. Groenewegen, M. A. T., Sloan, G. C., Soszynski,´ I. & Petersen, E. A. Luminosities and mass- loss rates of SMC and LMC AGB stars and red supergiants. Astron. Astrophys. 506, 1277–1296 (2009).

61. De Beck, E. et al. Probing the mass-loss history of AGB and red supergiant stars from CO rotational line proﬁles. II. CO line survey of evolved stars: derivation of mass-loss rate formulae. Astron. Astrophys. 523, A18 (2010). 62. Goldman, S. R. et al. The wind speeds, dust content, and mass-loss rates of evolved AGB and RSG stars at varying metallicity. Mon. Not. Roy. Astron. Soc. 465, 403–433 (2017).

63. Bowers, P. F., Johnston, K. J. & Spencer, J. H. Circumstellar envelope structure of late-type stars. Astrophys. J. 274, 733–754 (1983).

64. Justtanont, K., Skinner, C. J., Tielens, A. G. G. M., Meixner, M. & Baas, F. Modeling of the Dust and Gas Outﬂows from OH 26.5+0.6: The Superwind. Astrophys. J. 456, 337 (1996).

65. Jones, O. C. et al. On the metallicity dependence of crystalline silicates in oxygen-rich asymptotic giant branch stars and red supergiants. Mon. Not. Roy. Astron. Soc. 427, 3209–3229 (2012).

66. van Loon, J. T. et al. A Spitzer Space Telescope Far-Infrared Spectral Atlas of Compact Sources in the Magellanic Clouds. I. The Large Magellanic Cloud. Astron. J. 139, 68–95 (2010).

67. Groenewegen, M. A. T. et al. The ALMA detection of CO rotational line emission in AGB stars in the Large Magellanic Cloud. Astron. Astrophys. 596, A50 (2016).

68. Jimenez-Esteban,´ F. M., Agudo-Merida,´ L., Engels, D. & Garc´ıa-Lario, P. An infrared study of galactic OH/IR stars. I. An optical/near-IR atlas of the Arecibo sample. Astron. Astrophys. 431, 779–791 (2005).

69. Cox, N. L. J. et al. A far-infrared survey of bow shocks and detached shells around AGB stars and red supergiants. Astron. Astrophys. 537, A35 (2012).

70. Sahai, R. & Trauger, J. T. Multipolar Bubbles and Jets in Low-Excitation Planetary Nebulae: Toward a New Understanding of the Formation and Shaping of Planetary Nebulae. Astron. J. 116, 1357–1366 (1998).

71. Imai, H., Deguchi, S., Nakashima, J.-I., Kwok, S. & Diamond, P. J. The Spatiokinematical

Structure of H2O and OH Masers in the “Water Fountain” Source IRAS 18460-0151. Astro- phys. J. 773, 182 (2013).

72. Gielen, C., van Winckel, H., Min, M., Waters, L. B. F. M. & Lloyd Evans, T. SPITZER survey of dust grain processing in stable discs around binary post-AGB stars. Astron. Astrophys. 490, 725–735 (2008).

73. de Vries, B. L. et al. Micron-sized forsterite grains in the pre-planetary nebula of IRAS 17150- 3224. Searching for clues to the mysterious evolution of massive AGB stars. Astron. Astrophys. 576, A98 (2015). 74. Alibert, Y. et al. Theoretical models of planetary system formation: mass vs. semi-major axis. Astron. Astrophys. 558, A109 (2013).

75. Herman, J. & Habing, H. J. Time variations and shell sizes of OH masers in late-type stars. Astron. Astrophys. Suppl. 59, 523–555 (1985).

76. Engels, D. Water vapor masers in stars departing from the AGB. Astron. Astrophys. 388, 252–267 (2002).

77. Boboltz, D. A. & Marvel, K. B. Water Maser Kinematics in the Jet of OH 12.8-0.9. Astrophys. J. 665, 680–689 (2007).

78. Chong, S.-N., Imai, H. & Diamond, P. J. Water Masers in W43A: Early Morphological Changes of a Future Planetary Nebula. Astrophys. J. 805, 53 (2015).

79. Corradi, R. L. M. & Schwarz, H. E. Morphological populations of planetary nebulae: which progenitors? I. Comparative properties of bipolar nebulae. Astron. Astrophys. 293, 871–888 (1995).

80. Wright, N. J., Barlow, M. J., Ercolano, B. & Rauch, T. A 3D photoionization model of the extreme planetary nebula NGC 6302. Mon. Not. Roy. Astron. Soc. 418, 370–389 (2011).

81. Santander-Garc´ıa, M. et al. ALMA high spatial resolution observations of the dense molecular region of NGC 6302. Astron. Astrophys. 597, A27 (2017).

82. Molster, F. J. et al. Low-temperature crystallization of silicate dust in circumstellar disks. Nature 401, 563–565 (1999).

83. Kemper, F., Waters, L. B. F. M., de Koter, A. & Tielens, A. G. G. M. Crystallinity versus mass-loss rate in asymptotic giant branch stars. Astron. Astrophys. 369, 132–141 (2001).

84. Kim, H. et al. The large-scale nebular pattern of a superwind binary in an eccentric orbit. Nature Astronomy 1, 60 (2017). Supplementary Figures

12CO J=3-2 12CO J=2-1 2 12 3 2 10 1

[ 1 [ Jy Jy

8 km/s] /beam 0 km/s] /beam 0 -1 6 -1 -2 Offset [arcsec] Offset Offset [arcsec] Offset 4 -3 -2 2 -4

2 1 0 -1 -2 5 4 3 2 1 0 -1 -2 -3 Offset [arcsec] Offset [arcsec]

Supplementary Figure 1: Frequency-integrated zero-moment map of the 12CO J = 3–2 transition of OH 26.5+0.6. The spatial resolution is 000.22 000.18. The rms-value in the zero-moment map ⇥ is 0.13 Jy/beam km/s. All values above 1 are shown. Emission is detected in a 300.7 300.2 ⇥ rms ⇠ ⇥ region. The deconvolved full width at half maximum (FWHM) is 000.41. The contrast is best ⇠ visible on screen.

12CO J=1-0 12CO J=2-1 3 7 3 2 6 2 1 [ 1 [ 5 Jy Jy 0 km/s] /beam 0 km/s] /beam 4 -1 -1

-2 3 -2 Offset [arcsec] Offset Offset [arcsec] Offset -3 -3 2 -4 -4 1 5 4 3 2 1 0 -1 -2 -3 5 4 3 2 1 0 -1 -2 -3 Offset [arcsec] Offset [arcsec]

Supplementary Figure 2: Frequency-integrated zero-moment map of the 12CO J = 1–0 and 12CO J = 2–1 transitions of OH 30.1-0.7. The spatial resolution for the 12CO J = 1–0 data is 12 100.41 100.04 and for CO J = 2–1 data is 000.41 000.26. The -value in the zero-moment map is ⇥ ⇥ rms 0.04 Jy/beam km/s and 0.09 Jy/beam km/s for the 12CO J = 1–0 and J=2–1 data, respectively. All 12 values above 1 rms are shown. Emission is detected in a 700.4 region for the CO J = 1–0 data, ⇥12 ⇠ 12 and 600 for the CO J = 2–1 data. The deconvolved FWHM for the CO J = 2–1 data is 000.60. ⇠ ⇠ The contrast is best visible on screen. Supplementary Figure 3: 12CO J = 3–2 line emission in OH 26.5+0.6 plotted in polar coordi- nates. Radial intensity proﬁle as a function of the clockwise angle through the centre of the systemic velocity channel of the 12CO J = 3–2 emission of OH 26.5+0.6. Flux units in Jy/beam plotted on a logarithmic scale. The spiral arms are visible as the diagonally expanding green/blue ﬁlaments.

The bright white/pink region within 000.6 is the high-density zone created by binary interaction ⇠ which leads to a density enhancement which mimics the so-called superwind phase. The dashed red line indicates the point beyond which the spiral structure can easily be discerned in the ALMA data. In the upper panel, the straight dotted white lines are inserted to guide the eye to the trend for an Archimedean spiral, expected for a binary system in a circular orbit with the orbital plane located exactly on the plane of the sky (inclination i =0). Deviations of the observed pattern from the straight dotted white lines (in the upper panel) are caused by the inclination of orbital plane w.r.t. 40 the plane of the sky (i 60 – 75 for OH 26.5+0.6) leading to undulations . This is illustrated by ⇠ the curved dotted white lines in the bottom panel. A potential deviation of the binary orbit from a circle would lead to bifurcations84, an effect not discernible in the data. Supplementary Figure 4: (a) Side-view of the volume covered by the ballistic trajectories (orange shaded region) for a binary system similar to OH 26.5+0.6 or OH 30.1 0.7. (b) Ampliﬁcation factor as a function of ⌘, being the ratio of the terminal wind speed v to the orbital speed vorb, for 1 2 values of the parameter. Supplementary Video 1: Video of the 12CO J = 3–2 channel map of OH 26.5+0.6. Video slic- 1 12 ing through the velocity channel maps (between 8–46 km s ) of the CO J = 3–2 emission in OH 26.5+0.6. The ordinate and co-ordinate axis give the offset of the right ascension and declination, respectively, in units of arcseconds. North is up, East is left. The ALMA beam size is shown 1 as a white ellipse in the bottom left corner. The systemic velocity, vLSR, is at 28 km s . The 1 ⇠ front and last slice show the emission at v =24km s .

Supplementary Video 2: Video of the 12CO J = 2–1 channel map of OH 30.1-0.7. Video slic- 1 12 ing through the velocity channel maps (between 75–122.5 km s ) of the CO J = 2–1 emission in OH 30.1-0.7. The ordinate and co-ordinate axis give the offset of the right ascension and declination, respectively, in units of arcseconds. North is up, East is left. The ALMA beam size is shown 1 as a white ellipse in the bottom left corner. The systemic velocity, vLSR, is at 101 km s . The 1 ⇠ front and last slice show the emission at v =100km s .