‘Cataclysmic variables’ are binary star systems in which one star of the pair is a white dwarf, and which often generate bright and energetic stellar outbursts. Classical novae are one type of outburst: when the white dwarf accretes enough matter from its companion, the resulting hydrogen-rich atmospheric envelope can host a runaway thermonuclear reaction that generates a rapid brightening1–4. Achieving peak luminosities of up to one million times that of the Sun5 , all classical novae are recurrent, on timescales of months6 to millennia7 . During the century before and after an eruption, the ‘novalike’ binary systems that give rise to classical novae exhibit high rates of mass transfer to their white dwarfs8 . Another type of outburst is the dwarf nova: these occur in binaries that have stellar masses and periods indistinguishable from those of novalikes9 but much lower mass-transfer rates10, when accretiondisk instabilities11 drop matter onto the white dwarfs. The coexistence at the same orbital period of novalike binaries and dwarf novae—which are identical but for their widely varying accretion rates—has been a longstanding puzzle9 . Here we report the recovery of the binary star underlying the classical nova eruption of 11 March ad 1437 (refs 12, 13), and independently confirm its age by propermotion dating. We show that, almost 500 years after a classical-nova event, the system exhibited dwarf-nova eruptions. The three other oldest recovered classical novae14–16 display nova shells, but lack firm post-eruption ages17,18, and are also dwarf novae at present. We conclude that many old novae become dwarf novae for part of the millennia between successive nova eruptions19,

1. 5 5 8 | N A T U R E | V O L 5 4 8 | 3 1 au g u s t 2 0 1 7
Letter doi:10.1038/nature23644
Proper-motion age dating of the progeny of Nova
Scorpii ad 1437
M. M. Shara1,2
, K. Iłkiewicz3
, J. Mikołajewska3
, A. Pagnotta1
, M. F. Bode4
, L. A. Crause5
, K. Drozd3
, J. Faherty1
, I. Fuentes-
Morales6
, J. E. Grindlay7
, A. F. J. Moffat8
, M. L. Pretorius5,9
, L. Schmidtobreick10
, F. R. Stephenson11
, C. Tappert6
& D. Zurek1
‘Cataclysmic variables’ are binary star systems in which one
star of the pair is a white dwarf, and which often generate bright
and energetic stellar outbursts. Classical novae are one type of
outburst: when the white dwarf accretes enough matter from its
companion, the resulting hydrogen-rich atmospheric envelope
can host a runaway thermonuclear reaction that generates a rapid
brightening1–4
. Achieving peak luminosities of up to one million
times that of the Sun5
, all classical novae are recurrent, on timescales
of months6
to millennia7
. During the century before and after an
eruption, the ‘novalike’ binary systems that give rise to classical
novae exhibit high rates of mass transfer to their white dwarfs8
.
Another type of outburst is the dwarf nova: these occur in binaries
that have stellar masses and periods indistinguishable from those
of novalikes9
but much lower mass-transfer rates10
, when accretion-
disk instabilities11
drop matter onto the white dwarfs. The co-
existence at the same orbital period of novalike binaries and dwarf
novae—which are identical but for their widely varying accretion
rates—has been a longstanding puzzle9
. Here we report the recovery
of the binary star underlying the classical nova eruption of 11 March
ad 1437 (refs 12, 13), and independently confirm its age by proper-
motion dating. We show that, almost 500 years after a classical-nova
event, the system exhibited dwarf-nova eruptions. The three other
oldest recovered classical novae14–16
display nova shells, but lack
firm post-eruption ages17,18
, and are also dwarf novae at present.
We conclude that many old novae become dwarf novae for part of
the millennia between successive nova eruptions19,20
.
One of the best located novae of antiquity, recorded by Korean royal
astronomers, erupted on 11 March ad 1437 (ref. 21). It lay within the
asterism Wei (the tail of the constellation Scorpius), within half a chi
(roughly one degree) of one of the stars ζ​Sco or η​Sco (see Methods).
It was seen for 14 days before vanishing, consistent with it being a
fast-declining classical nova, and ruling out the possibility that it was
a supernova.
We began by searching a pair of U- and B-band photographic plates
from 1985, obtained with the Anglo–Australian 1.2-metre Schmidt
telescope, that were centred on Scorpius; our search yielded dozens
of ultraviolet-bright objects, but only one that was surrounded by a
clear, shell-like structure. The shell is marginally visible on other sky-­
survey plates, but is seen clearly in a narrowband Hα​image (Fig. 1).
The ­central star of the Scorpius nebula is not a cataclysmic variable.
However, a star located about 15″​from the shell centre is both the
strong-emission-line cataclysmic variable 2MASS J17012815-4306123
and the variable X-ray source IGR J17014-4306 (ref. 22).
The measured proper motion of this magnitude-17 cataclysmic
­variable, which is listed in multiple catalogues, displays large scatter.
This is because all of the Southern Hemisphere catalogues to date are
limited to short, 20-year baselines, and because the star is blended with
the image of a nearby neighbour. However, we have located the star on
the digitized DASCH23
(Digital Access to a Sky Century @ Harvard)
photographic plate A12425, taken in 1923 (see Fig. 2 for details). We
measured the cataclysmic variable’s position relative to that of nearby
stars both on that plate and on a sub-arcsecond charge-coupled-­
device (CCD) image taken in 2016 (see Methods). The resulting
93-year baseline has enabled us to measure the cataclysmic variable’s
proper motion with far higher precision than was previously possible,
­showing that μα (the angular change in the star’s right ascension, RA)
is −​12.74 ±​ 1.79 milliarcseconds per year, and μδ (the angular change
in the star’s declination, dec.) is −​27.72 ±​ 1.21 milliarcseconds per year.
1
Department of Astrophysics, American Museum of Natural History, Central Park West and 79th Street, New York, New York 10024, USA. 2
Institute of Astronomy, The Observatories, Madingley
Road, Cambridge CB3 0HA, UK. 3
N. Copernicus Astronomical Center, Polish Academy of Sciences, Bartycka 18, PL 00–716 Warsaw, Poland. 4
Astrophysics Research Institute, Liverpool John
Moores University, IC2 Liverpool Science Park, Liverpool L3 5RF, UK. 5
South African Astronomical Observatory, PO Box 9, Observatory, 7935 Cape Town, South Africa. 6
Instituto de Física y
Astronomía, Universidad de Valparaíso, Avenida Gran Bretaña 1111, 2360102 Valparaíso, Chile. 7
Harvard–Smithsonian Center for Astrophysics, The Institute for Theory and Computation, 60
Garden Street, Cambridge, Massachusetts 02138, USA. 8
Département de Physique and Centre de Recherche en Astrophysique du Québec (CRAQ), Université de Montréal, CP 6128 Succ. C-V,
Montréal, QC H3C 3J7, Canada. 9
Department of Astronomy, University of Capetown, Private Bag X3, Rondebosch 7701, South Africa. 10
European Southern Observatory, Alonso de Cordova 3107,
7630355 Santiago, Chile. 11
Department of Physics, Durham University, South Road, Durham DH1 3LE, UK.
1 arcmin
Figure 1 | The recovered nova of ad 1437 and its ejected shell. This Hα​
image was taken with the Swope 1-metre telescope and its CCD camera
in June 2016, with a total of 6,000 seconds of exposure. Images taken by
Swope were processed and combined with standard PyRAF and IRAF
procedures. Here, north is up and east is to the left. The location of the
cataclysmic variable in 2016 is indicated with red tick marks. Its proper
motion places the ad 1437 cataclysmic variable 7.4″​east and 16.0″​ north
of its current position, at the red plus sign. The position of the centre of
the shell in 2016 and its deduced position in 1437 (see text) are indicated
with blue and green plus signs, respectively. The 1437 positions of the shell
centre and of the cataclysmic variable agree to within 1.7″​, and their 1σ
error ellipses overlap.
© 2017 Macmillan Publishers Limited, part of Springer Nature. All rights reserved.

2. 3 1 au g u s t 2 0 1 7 | V O L 5 4 8 | N A T U R E | 5 5 9
Letter RESEARCH
(We determined the 1σ errors empirically from the proper motions of
dozens of stars on both plates.) This measured proper motion places
the cataclysmic variable, in ad 1437, at RA =​ 17:01:28.53 ±​ 1.0″​ and
dec. =​ −​43:05:56.7 ±​ 0.7″​(coordinates given in the J2000 system).
The centre of the nova shell in 2016 is determined from its edges
(see Fig. 1 and Methods) to be at RA 17:01:28.55 ±​ 1.43″​and dec.
−​43:05:59.3 ±​ 0.4″​.
Unlike their underlying cataclysmic variables, nova shells ­decelerate
by sweeping up interstellar matter, halving their speeds on a mean
timescale of 75 years24
. Assuming that the Nova 1437 ejecta initially
had the same proper motion that we measure for the cataclysmic
­variable, and that this value has been halved every 75 years since, then
the shell centre in 1437 must have been located +​1.43″​(east) in right
ascension and +​3.1″​(north) in declination of its 2016 position, at RA
17:01:28.68 ±​ 1.43″​and dec. −​43:05:56.2 ±​ 0.4″​.
The difference between the centre of the nova shell in 1437 and the
proper-motion-determined position of the cataclysmic variable in
that year is 1.7″​. The 1σ error ellipses of the two positions overlap,
­confirming that the cataclysmic variable is in fact the nova of ad 1437
and the source of the shell that we observe today. Fortuitously, this cata-
clysmic variable is a deeply eclipsing binary star, enabling us to measure
its orbital period (0.5340263 ( ±​ 5 ×​ 10−7
) days), and to characterize its
stellar components (see Methods).
Figure 3 shows a DASCH-based light curve of the cataclysmic
­variable from 1919 to 1951. Three dwarf-nova eruptions (in 1934, 1935
and 1942) are clearly seen, wherein the cataclysmic variable brightened
by two to four magnitudes. In Fig. 4 we show images before, during and
after one of these dwarf-nova eruptions (from 1942). The classical nova
of 11 March 1437, seen 497 years after eruption, has become a dwarf
nova, supporting the view that novalike binaries and dwarf novae are
indeed the same systems, seen at different times after classical nova
explosions20
.
Online Content Methods, along with any additional Extended Data display items and
Source Data, are available in the online version of the paper; references unique to
these sections appear only in the online paper.
received 11 March; accepted 26 June 2017.
1. Starrfield, S., Truran, J. W., Sparks, W. M. & Kutter, G. S. CNO abundances and
hydrodynamic models of the nova outburst. Astrophys. J. 176, 169–176
(1972).
2. Prialnik, D., Shara, M. M. & Shaviv, G. The evolution of a slow nova model with a
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3. Kraft, R. P. Binary stars among cataclysmic variables. III. Ten old novae.
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5. Warner, B. Cataclysmic Variable Stars (Cambridge Astrophys. Series 28,
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6. Darnley, M. J. et al. A remarkable recurrent nova in M31: the optical
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8. Collazzi, A. C. et al. The behavior of novae light curves before eruption. Astron. J.
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9. Knigge, C., Baraffe, I. & Patterson, J. The evolution of cataclysmic variables as
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1915 1920 1925 1930 1935 1940 1945 1950 1955
Bmagnitude
Year
Upper limit Detection
Figure 3 | The light curve of Nova Scorpii ad 1437 from 1919
through to 1951. The grey symbols show upper limits, while the
black dots are measured DASCH detections of the star, with typical 1σ
errors of ±​0.1–0.2 mag. The star was first detected in quiescence near
magnitude 17 in 1923, and near magnitude 16 in 1925. Dwarf-nova
outbursts in 1934 (reaching nearly magnitude 12), 1935 and 1942
(see Fig. 4) are evident.
1 arcmin
29 April 1942 18 May 1942
6 June 1942 9 June 1942
Figure 4 | A series of images of the old nova spanning six weeks in
1942. The star is seen to undergo a dwarf-nova eruption, brightening
substantially between the first and third images, then returning to
quiescence a few days later. The dwarf nova is indicated with an arrow
in each epoch in this series of Harvard DASCH MF-series photographic
plates (see Methods). North is up and east is to the left.
1 arcmin
Figure 2 | A B-band photographic image of the old nova, seen on
10 June 1923. The Harvard photographic plate A12425, part of which
is shown here, is a 300-minute exposure obtained using the 24″​ Bruce
Doublet telescope at the Harvard Observatory station in Arequipa, Peru.
North is up; east is to the left; and the cataclysmic variable is indicated with
red tick marks.
© 2017 Macmillan Publishers Limited, part of Springer Nature. All rights reserved.

3. 5 6 0 | N A T U R E | V O L 5 4 8 | 3 1 au g u s t 2 0 1 7
LetterRESEARCH
10. Krzemi’nski, W. & Smak, J. Eruptive binaries. III. The recurrent nova WZ
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of Classical Novae: Lecture Notes in Classical Physics Vol. 369 (eds Cassatella, A.
& Viotti, R.) 57–58 (Springer, 1990).
13. Black, D. T. R., Bode, M. F., Stephenson, F. R., Abbott, T. & Page, K. L. Search for
the Remnant of the Nova of 1437. In RS Ophiuchi (2006) and the Recurrent
Nova Phenomenon: Astron. Soc. Pacific Conference Proceedings Vol. 401
(eds Evans, A., Bode, M. F., O’Brien, T. J. & Darnley, M. J.) 351–354 (Astron. Soc.
Pacific, 2008).
14. Shara, M. M. et al. An ancient nova shell around the dwarf nova Z
Camelopardalis. Nature 446, 159–162 (2007).
15. Shara, M. M. et al. AT Cnc: a second dwarf nova with a classical nova shell.
Astrophys. J. 758, 121–125 (2012).
16. Miszalski, B. et al. Discovery of an eclipsing dwarf nova in the ancient nova
shell Te 11. Mon. Not. R. Astron. Soc. 456, 633–640 (2016).
17. Shara, M. M. et al. The inter-eruption timescale of classical novae from
expansion of the Z Camelopardalis shell. Astrophys. J. 756, 107–112
(2012).
18. Shara, M. M., Drissen, L., Martin, T., Alarie, A. & Stephenson, F. R. When does an
old nova become a dwarf nova? Kinematics and age of the nova shell of the
dwarf nova AT Cancri. Mon. Not. R. Astron. Soc. 465, 739–745 (2017).
19. Vogt, N. The structure and outburst mechanisms of dwarf novae and their
evolutionary status among cataclysmic variables. Astron. Gesell. 57, 79–118
(1982).
20. Shara, M. M., Livio, M., Moffat, A. F. J. & Orio, M. Do novae hibernate during
most of the millenia between eruptions? Links between dwarf and classical
novae, and implications for the space densities and evolution of cataclysmic
binaries. Astrophys. J. 311, 163–171 (1986).
21. Clark, D. H. & Stephenson, F. R. The Galactic Supernovae 1st edn (Pergamon,
1977).
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spectroscopy. X. A new multi-year, multi-observatory campaign. Astron.
Astrophys. 556, A120–A141 (2013).
23. Grindlay, J. et al. Preserving astronomy’s photographic legacy: current state
and the future of North American astronomical plates. Astron. Soc. Pac. Conf.
Ser. 410, 101–110 (2009).
24. Duerbeck, H. W. The interaction of nova shells with the interstellar medium.
Astrophys. Space Sci. 131, 461–466 (1987).
Acknowledgements J.M., K.I., K.D. and M.M.S. acknowledge support by the
Polish Narodowe Centrum Nauki (grant DEC-2013/10/M/ST9/00086). M.M.S.
acknowledges the late P. Newman and the Newman’s Own Foundation, whose
support made the participation of the American Museum of Natural History
(AMNH) in the South African Large Telescope (SALT) possible. A.F.J.M. thanks
the National Sciences and Engineering Research Council of Canada and the
Fonds de Recherche Nature et Technologies (Quebec) for financial support.
A.P. acknowledges support from the AMNH’s Kathryn W. Davis Postdoctoral
Scholar program, which is supported in part by the New York State Education
Department and by the National Science Foundation (NSF) under grant
numbers DRL-1119444 and DUE-1340006. M.M.S. acknowledges the
hospitality of the Institute of Astronomy at the University of Cambridge. Some of
the observations reported here were obtained with the SALT under programme
2016-1-SCI-044, and with the South African Astronomical Observatory’s
1.9-metre and 1.0-metre telescopes. Polish participation in SALT is funded by
grant no. MNiSW DIR/WK/2016/07. We thank the Harvard–Smithsonian Center
for Astrophysics team for making DASCH data available to the astronomical
community. The DASCH project at Harvard is partially supported from NSF
grants AST-0407380, AST-0909073 and AST-1313370. The Image Reduction
and Analysis Facility (IRAF) is distributed by the National Optical Astronomy
Observatories, which is operated by the Association of Universities for Research
in Astronomy under a cooperative agreement with the NSF. This research
has made use of data obtained from the Chandra Data Archive and software
provided by the Chandra X-ray Center in application packages CIAO, ChIPS
and Sherpa. All authors thank the referees for thoughtful and constructive
suggestions.
Author Contributions All authors shared in the ideas and the writing of this
paper. M.M.S., A.F.J.M. and M.F.B. carried out optical surveys for the nova on the
basis of F.R. Stephenson’s predictions. M.M.S. and A.F.J.M. located the nebula
associated with the old nova. Broadband CCD imaging and data reduction for
the candidate were carried out by L.A.C., M.L.P., I.F.-M. and K.D. Narrowband
imaging of the shell and reduction of those images were carried out by K.I., who
also produced the cataclysmic binary’s X-ray light curve and deduced its period.
J.E.G. retrieved the 1923 digitized image of the nova. A.P. and J.E.G. produced
the old nova’s historical light curve. A.P. and K.I. measured the old nova’s
proper motion. M.M.S, K.I. and J.M. determined the orbital period, while I.F.-M.,
K.I. and J.M. determined the white dwarf’s spin period. J.M. determined the
companion’s spectral type, the system’s mass function, and its distance, while
J.M. and M.M.S. found the limit on the shell’s mass.
Author Information Reprints and permissions information is available at
www.nature.com/reprints. The authors declare no competing financial
interests. Readers are welcome to comment on the online version of the paper.
Publisher’s note: Springer Nature remains neutral with regard to jurisdictional
claims in published maps and institutional affiliations. Correspondence and
requests for materials should be addressed to M.M.S. (mshara@amnh.org).
Reviewer Information Nature thanks C. Knigge and S. Shore for their
contribution to the peer review of this work.
© 2017 Macmillan Publishers Limited, part of Springer Nature. All rights reserved.

4. Letter RESEARCH
Methods
The historical nova Scorpii of 11 March ad 1437. The Sejong Sillok (‘Veritable
Records of [the reign of] King Sejong’) is a detailed chronicle of the reign of King
Sejong (who ruled Korea from 1418 to 1464), written in classical Chinese. Chapter
76 of the Sillok notes: “19th year of King Sejong, 2nd lunar month, day yichou [the
2nd day of the 60-day cycle], A meteor (liuxing) appeared... A solar halo... A guest
star (kexing) began to be (shi) seen between the second and third stars of Wei.
It was nearer to the third star, about half a chi (“half a foot”) away. It lasted (jiu)
for 14 days.” It was quite usual for East Asian astronomers to use “chi”—a linear
unit—as an angular unit. Descriptions of the position of the ecliptic relative to the
28 Chinese constellations (or ‘mansions’) through which the moon passes suggest
that 1 chi, as used in China, was roughly 1.50 degrees (ref. 25). Early Chinese
records of planetary conjunctions in which separations were quoted in chi suggest
that the chi/degree ratio is very roughly one. However, values in the range 0.44 to
2.8 have been determined25
, so that half a chi in China was roughly in the range
0.22–1.4 degrees. There is no similar determination for the Korean chi.
The Sejong Sillok refers to the guest star as being between the second and third
stars of Wei, while the date corresponds to 11 March ad 1437. Wei, the sixth lunar
lodge, contained nine stars in Scorpius26
. Listed in clockwise order they are ε​, μ​, ζ​,
η​, θ​, ι​, κ​, λ​ and ν​Sco. The numbering of the stars in Wei requires deduction, as no
star map from ancient Korea that lists star numbers in constellations exists. Korean
astronomers adopted asterisms similar to those of their Chinese colleagues, so it is
reasonable to look to Chinese star maps and lists for guidance. Two Chinese Yuan
dynastylistsdifferintheirordering:onehasε​Sco(1),μ​Sco(2),ζ​Sco(3)andη​Sco(4),
and the other has μ​Sco (1), ε​Sco (2), ζ​Sco (3) and η​Sco (4). The reasonable
deduction is that ζ​Sco is the third star of Wei.
Given that μ​Sco is adjacent to ζ​Sco, it is also reasonable to identify μ​as the
second star of Wei, as mentioned in the guest-star text. The guest star would then
be in the range 0.22–1.4 degrees north of ζ​Sco. Searches based on this predicted
position have proven fruitless12,13
.
An alternative numbering of the stars of Wei is based on μ​Sco, the ­determinant
star of Wei, which fixed the boundary of the lunar-lodge RA zone. Starting from
μ​Sco, and proceeding strictly clockwise around Wei, the second star in the text
would be ζ​Sco, while the third would be η​Sco. If this numbering is correct, then
one should find the old nova between ζ​Sco and η​Sco—that is, to the east of
ζ​Sco. This is, in fact, where we find the cataclysmic variable and nova shell that are
the subject of this study. In light of the uncertainties noted above concerning the
numbering of the stars in Wei in Korea, the size of the Chinese chi, and the even
greater uncertainty of the size of that unit in Korea, we conclude that the observed
angular distances of 1.95 degrees and 1.55 degrees from the cataclysmic variable
to η​Sco and to ζ​Sco, respectively, are in reasonable accord with the historical text.
The nova shell and its centre. The nova shell was imaged with the 1.0-metre
Swope telescope at Las Campanas Observatory on 15 and 17 June 2016. The
­observations were carried out with an E2V CCD231-84 CCD camera with a
pixel size of 0.435 arcsec. The total exposure time was 6,000 seconds through a
­narrowband Hα​filter. The images were reduced with standard IRAF procedures.
The ­cataclysmic variable is not at the centre of the nebula (Fig. 1). The nova shell
shows ­complex morphology, with a number of brighter knots, and strongest
­emission at the ­southeast edge of the nebula. Moreover, there is a faint outer lobe
of nebulosity to the northwest of the nova shell. At the northeast corner of the
nebula there is a faint tail-like structure extending up to a few arcmin away from
the nebula. Radial velocities will be required to distinguish shell material from
diffuse interstellar gas emission, which is strong in the direction of Scorpius.
We used the outer edges of the Hα​shell to determine the centre of the shell27
.
An initial cut through the shell in the north/south direction was ­perpendicularly
bisected, and the centre of the bisector was retained as the starting ­centre
­position. A new cut through the starting centre was made approximately
10 degrees ­clockwise from the starting cut, and then it was perpendicularly
bisected, again retaining the centre of the bisector as the next centre ­position.
This procedure was repeated twelve times until the centre measurements
­converged to within a pixel of each other. The last five iterations—where the con-
vergence was strongly evident—were averaged together to obtain the ­measured
centre, with the standard deviation of those five positions used as the 1σ uncer-
tainty on each measurement.
The nova orbital period and ephemeris. Photometric monitoring of the
­system was carried out with four telescopes. On 27, 29, 30 and 31 July 2016, we
observed the old nova with the 2.5-metre Du Pont telescope at the Las Campanas
Observatory, using the SIT e2K-1 camera and a V filter. Exposure times were
30 seconds on 31 July and 40 seconds on the other nights. On 15, 16, 20, 23, 24, 25
and 26 August 2016, the object was observed with a Sutherland high-speed optical
camera28
on the 1.0-metre South African Astronomical Observatory telescope,
using the g′​filter and with 4 ×​ 4 binning. The exposure times were 15 seconds
on 20, 25 and 26 August, 25 seconds on 23 August, and 20 seconds on the other
nights. On 12, 13, 20 and 21 September 2016, the nova was observed with the
Swope telescope in the B band. The exposure times were 180 seconds. On 21 July
2014 the nova was observed in white light with a Sutherland high-speed optical
camera28
on the 1.9-metre South African Astronomical Observatory telescope with
5-second exposures. The data were processed with standard IRAF procedures. The
magnitudes were transformed to a standard system, using all of the stars in the
field of view of each of the telescopes as standard stars. The reference magnitudes
were taken from the American Association of Variable Star Observers’ All-Sky
Photometric Survey29
.
The cataclysmic-variable light curve shows deep and short eclipses, and
­ellipsoidal variability (Extended Data Fig. 1). We applied the phase dispersion
minimization30
methodology, which is well suited to light curves with long gaps
and relatively short eclipses; it determined the orbital period to be 0.5340257 days,
or 12 hours 48 minutes 59.8 seconds (Extended Data Fig. 2). We also fit a linear
ephemeris to the measured times of eclipses, which gave a minimum heliocentric
Julian day (HJD) of:
= . ± × + . ± × ×− −
EHJD 2457626 3643 ( 3 10 ) 0 5340263 ( 5 10 ) (1)min
4 7
where E is the number of orbital periods since HJDmin =​ 2457626.3643. A detailed
view of the eclipses in three passbands is shown in Extended Data Fig. 3.
The white-dwarf spin period. In addition to the orbital-induced variability, there
is a periodic variability on a timescale of 30 minutes that is visible outside of the
eclipses. We fitted and removed a low-order polynomial from each data set to
remove orbital-induced variability. Discrete Fourier transforms performed on
these altered light curves with the Period04 program31
resulted in an ephemeris of:
= . ± × + . ± × ×− −
EHJD 2457625 059060 ( 6 10 ) 0 0215175 ( 8 10 ) (2)max
6 7
for the maxima of this variability. Over the 11 days monitored by the Sutherland
high-speed optical camera, the period remained stable at 1,859.112 seconds,
demonstrating that this variability is not flickering, but instead is due to the spin
period of the white dwarf. This, in turn, suggests that the system is an intermediate
polar (as does the presence of He ii emission; see below). The object shows two
pulses per period, with different amplitudes (Extended Data Fig. 1). The detailed
analysis of the pulse profiles is beyond the scope of this paper, though we note that
the pulse profiles varied slightly from night to night.
The nova was observed with the Chandra Advanced CCD Imaging
Spectrometer (ACIS)32
, which covers the energy range 0.2–10 keV. The obser-
vation was performed on 30 June 2015 and the total exposure time was 10.07 ks.
We extracted the light curve using the software programs CIAO version 4.8.1
(ref. 33) and CALDB version 4.7.2. The adopted bin size of the light curve
was 120 seconds. The object showed spin variability related to the spin period,
similar to that observed in the optical range (Fig. 2). We performed a discrete
Fourier transform on the Chandra light curve and calculated the errors using
a Monte Carlo simulation with the Period04 program. The resultant period of
0.0218 ±​ 0.0003 days (1,859 ±​ 26 seconds) is consistent with the period derived
from the optical data. The simultaneous presence of the same periodic variability
in X-ray and optical bands, measured a year apart, confirms that this permanent
period is indeed the spin period of the white dwarf.
The white-dwarf mass. We obtained low-resolution, long-slit spectra of
the cataclysmic variable and its shell (Extended Data Fig. 4) with the Robert Stobie
Spectrograph (RSS)34,35
and the 10-metre Southern African Large Telescope
(SALT)36,37
. We employed grating PG0900 and a slit with a projected width of 1.5″​,
which resulted in a resolving power of about 1,000. The observations were reduced
using standard IRAF procedures and the SALT RSS science pipeline38
. The ­brightest
knot in the nova shell was observed on 24 September 2016. The cataclysmic
variable was observed on 14 and 19 July 2016, and twice on 23 September 2016.
The spectra of the cataclysmic variable reveal broad (full width around
2,000–2,700 km per second) emission lines for the H i Balmer, He ii 4,686-Å and
He i 5,876-Å and 6,678-Å lines. A wealth of absorption lines indicate the presence
of an early K-type secondary star (Extended Data Fig. 4), while diffuse ­interstellar
bands at 5,780 Å, 5,797 Å and 6,281 Å are also seen. The relative intensities of
temperature-sensitive metal absorption lines, as well as the simultaneous presence
of a G-band, MgH bands and a very weak TiO band at 6,159 Å, are consistent
with a K3 star. This spectral classification is also consistent with the broadband
magnitudes and colours observed during the eclipses (B−​V is about 1.3–1.4), and
with 2MASS JHK magnitudes, and with moderate reddening E(B–V) of about
0.3–0.4. The H i Balmer line ratios in the low-density nebula/shell surrounding
the cataclysmic variable give a similar E(B–V) of about 0.3–0.6 (assuming case B
recombination) and an absorption A(V) of about 1. The 21-cm emission along this
line of sight39
gives a total hydrogen column of 1.4 ×​ 1022
per cm2
. This is consistent
with an E(B–V) of around 0.3, because the system velocity, from the radial-velocity
curve, implies a line-of-sight column that is about 10% of the total.
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5. LetterRESEARCH
Given that three of these spectra were, fortuitously, taken close to both spectro-
scopic quadratures, we can estimate the preliminary radial-velocity amplitude of
the secondary star (see Extended Data Fig. 5) to be Ksec =​ 220 ±​ 27 km per second
(assuming that the eclipses coincide with its inferior conjunction, and adopting
the photometric ephemeris). The mass function is then = . − .
+ .
f m M( ) 0 590 0 166
0 245
,
where M is the mass of the Sun. The best fit to the four radial-velocity data points
results in a secondary velocity of Ksec =​ 260 ±​ 6 km per second, with the spectro-
scopic conjunction occurring 0.035 orbital periods after the eclipse (Fig. 4). The
mass function is then = . − .
+ .
f m M( ) 0 974 0 055
0 105
While the phase shift might be the result of small-sample statistics (only four
radial-velocity points), such a shift between spectroscopic conjunction and pho-
tometric eclipse is not unusual in magnetic cataclysmic variables. The eclipse in
such cases is not that of the white dwarf but rather that of the principal accretion
spot, which can be displaced from the line connecting the two stars40
.
The mass functions obtained here demonstrate that the white dwarf must be
massive and that the secondary star must be evolved. Nova 1473 ad was rather fast,
being visible for only 14 days, so the white-dwarf mass7
is likely to be greater than
1.0 M . For a white-dwarf mass in the range 1.0–1.4 M , the mass of the
­companion is 0.3–0.76M if Ksec =​ 220 ±​ 27 km per second, and less than about
0.2 M if Ksec =​ 260 ±​ 6 km per second. Given that the secondary star is filling its
Roche lobe, we can estimate the system distance (d). For a white dwarf of mass
1.0 M , the system distance is in the range of some 650 to 980 parsecs (pc), while
for a white dwarf of mass 1.4 M , the system distance is about 540 pc.
The cataclysmic-variable emission lines have variable, complex profiles, but the
changes do not seem to show orbital modulation. They may be varying on shorter
timescales related to the 1,859-second rotation period of the white dwarf. The
spectrum of the brightest knot (southeast of the cataclysmic variable) in the shell
reveals many emission lines (see Fig. 2), including very strong (in comparison with
Hα​) [N ii] and [S ii] lines. The [O iii] line is relatively weak (fainter than the Hβ​
line), which is unusual, given the presence of He ii 4,686 emission. The [N ii]/Hα​,
[S ii]/Hα​and [O iii]/Hβ​flux line ratios are intermediate between those of
­planetary nebulae, supernova remnants and H ii regions, as seen in emission-line
diagnostic diagrams41
. The ratio of [S ii] 6,716/6,731 is 1.45, indicating a very low
electron density, ne, of less than about 100 cm−3
, whereas the [N ii] line ratio of
(6,548+​6,583)/5,755 =​ 100 implies an effective temperature of 9,400 K (ref. 42).
The shell mass. The Hα​ +​ [Ν​ ii] flux of the nova shell is 2.8 ×​ 10−15
W per m2
,
and our measurement of the Hα​/[Ν​ ii] flux ratio, from our SALT spectra, is about
0.6 (Extended Data Fig. 4). The upper limit on electron density is less than some
100 cm−3
, as noted above. This density and the distance derived above yield
ne
2
V =​ 2.35 ×​ 1056
cm−3
(d/500 pc)2
. Allowing for a reddening of E(B−​V) of about
0.3–0.4, derived above, and hence an A(V) of about 1, finally yields an upper limit
on the radiating hydrogen gas in the nova shell to be Mshell <​ 0.004 M (d/500 pc)2
.
The ejecta in nova shells decelerate to half of their initial velocities by sweeping
up interstellar matter and doubling their masses on timescales of about 75 years24
.
The Korean historical records note that Nova Scorpii disappeared after 14 days,
so it was a fast nova, and the velocity of material initially ejected must have been
more than 1,200 km per second7
. The ejecta spectrum (Fig. 4) is of low resolution,
but with it we are able to place an upper limit of 300 km per second on the current
ejecta velocity. Thus the shell has been decelerated to 25% or less of its initial speed,
which can happen only if it has at least quadrupled its own mass.
Even if the ejected shell only underwent two successive mass doublings (and it
may have undergone seven or eight such doublings), a hard upper limit on the mass
ejected in the nova eruption is Mej <​ 10−3
M . This rules out any chance of the
shell having a planetary nebula origin, because the masses of planetary nebula
shells are typically 0.1–1.0 M (ref. 43).
Data availability. All relevant data, including all figures, photometry and
­spectroscopy, are available from the corresponding author on reasonable request.
25. Kiang, T. The past orbit of Halley’s comet. Mem. R. Astron. Soc 76, 27–66
(1972).
26. Fang, H. L. & Ho, P. Y. The Astronomical Chapters of the Jinshu (Mouton & Co.,
1966).
27. Schaefer, B. E. & Pagnotta, A. An absence of ex-companion stars in the type Ia
supernova remnant SNR 0509-67.5. Nature 481, 164–166 (2012).
28. Coppejans, R. et al. Characterizing and commissioning the Sutherland
high-speed optical cameras (SHOC). Publ. Astron. Soc. Pacif. 125, 976–988
(2013).
29. Henden, A. A., Levine, S. E., Terrell, D., Smith, T. C. & Welch, D. Data release 3 of
the AAVSO all-sky photometric survey (APASS). J. Am. Assoc. Var. Star Obs. 40,
430–430 (2012).
30. Stellingwerf, R. F. Period determination using phase dispersion minimization.
Astrophys. J. 224, 953–960 (1978).
31. Lenz, P. & Breger, M. Period04 user guide. Comm. Asteroseismol. 146, 53–136
(2005).
32. Garmire, G. P., Bautz, M. W., Ford, P. G., Nousek, J. A. & Ricker, G. R. Jr Advanced
CCD imaging spectrometer (ACIS) instrument on the Chandra X-ray
Observatory. Proc. SPIE 4851, 28–44 (2003).
33. Fruscione, A., McDowell, J. C., Allen, G. E. & Houck, J. C. CIAO: Chandra’s data
analysis system. Proc. SPIE 6270, 62701V (2006).
34. Burgh, E. B. et al. Prime focus imaging spectrograph for the Southern African
Large Telescope: optical design. Proc. SPIE 4841, 1463–1471 (2003).
35. Kobulnicky, H. A. et al. Prime focus imaging spectrograph for the Southern
African large telescope: operational modes. Proc. SPIE 4841, 1634–1644
(2003).
36. Buckley, D. A. H., Swart, G. P. & Meiring, J. G. Completion and commissioning of
the Southern African Large Telescope. Proc. SPIE 6267, 62670–62677 (2006).
37. O’Donoghue, D. et al. First science with the Southern African Large Telescope:
peering at the accreting polar caps of the eclipsing polar SDSS J015543.40+​
002807.2. Mon. Not. R. Astron. Soc. 372, 151–162 (2006).
38. Crawford, S. M. et al. PySALT: the SALT science pipeline. Proc. SPIE. 7737,
25–36 (2010).
39. Kalberla, P. M. W. & Haud, U. GASS: the Parkes galactic all-sky survey. III.
Astron. Astrophys. 578, A78 (2015).
40. Beuermann, K. & Reinsch, K. High-resolution spectroscopy of the intermediate
polar EX Hydrae. I. Kinematic study and Roche tomography. Astron. Astrophys.
480, 199–212 (2008).
41. Kniazev, A. Y., Pustilnik, S. A. & Zucker, D. B. Spectroscopy of two PN candidates
in IC10. Mon. Not. R. Astron. Soc. 384, 1045–1052 (2008).
42. Osterbrock, D. E. Astrophysics of Gaseous Nebulae and Active Galactic Nuclei
(Univ. Science Books, 1989).
43. Frew, D. J., Bojicic, I. S. & Parker, Q. A. A catalogue of integrated Hα​fluxes for
1258 galactic planetary nebulae. Mon. Not. R. Astron. Soc. 431, 2–26 (2013).
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6. Letter RESEARCH
Extended Data Figure 1 | Visible and X-ray light curves of Nova Scorpii
1437. a–c, Phase plots of Nova 1437 photometry in V, g′​and B bands. The
observations are phased with the orbital period (Porb), using the ephemeris
of equation (1). d, Phase plot of g′​photometry, after subtracting all
variability related to the orbital motion (see Methods). The observations
were phased with the spin period (Pspin) of the white dwarf, using the
ephemeris of equation (2). e, As for panel d, but with points binned with a
bin size of 0.025 ×​ Pspin. f, Chandra observations in the 0.2–10 keV band,
phased with the spin period of the white dwarf using the ephemeris from
equation (2). All error bars are 1σ.
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7. LetterRESEARCH
Extended Data Figure 2 | The phase dispersion minimization (PDM) statistic as a function of frequency for the light curves of Nova
Scorpii 1437. This PDM plot (see text) allows us to determine the binary orbital frequency to be 1.8725691 per day, corresponding to an orbital
period of 0.5340257 days.
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8. Letter RESEARCH
Extended Data Figure 3 | Light curves of Nova Scorpii 1437, centred on eclipses. Measurements are shown in V, g′​and B filters, and the error
bars are 1σ.
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9. LetterRESEARCH
0
0.1
0.2
0
0.5
1
4000 4500 5000 5500 6000 6500
1
2
3
4650 4800
1.5
2
Extended Data Figure 4 | SALT-based spectra of the nova shell, the
old nova and a synthetic spectral standard. Top, SALT spectra for the
brightest region on the nova shell (southeast of the cataclysmic variable),
with the main emission lines identified. Note the strong lines of [S ii]
6,716 Å and 6,731 Å, and of [N ii] at 6,548 Å and 6,583 Å. The y-axis
shows the flux of the relevant emission line with respect to the flux of the
Hα​line. Bottom, SALT spectrum of the cataclysmic variable taken on 23
September 2016, with the synthetic spectrum of a K3 V star overlaid (with
effective temperature Teff =​ 4,750 K, gravitational acceleration log g =​ 4.5,
and Solar composition), reddened with A(V) ≈​ 1. The insert shows the
emission profiles of Hβ​as well as those of He ii and the Bowen C iii-N iii
blend.
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10. Letter RESEARCH
Extended Data Figure 5 | The radial-velocity curve of Nova Scorpii
1437. The radial velocities (VHEL, where ‘HEL’ is ‘heliocentric’) were
obtained by measuring the wavelength differences between the 20
strongest absorption features in the observed spectrum and those in the
synthetic spectrum of Fig. 4. The systemic velocity is −​46 km per second,
and the error bars are 1σ. The solid curve corresponds to the secondary
star’s inferior conjunction occurring at mid-eclipse. The (better-fitting)
dashed curve corresponds to the inferior conjunction that occurs 0.035
orbital periods after the eclipse. See text for details. Tconj and Tecl, time of
conjunction and time of eclipse.
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