# Revisiting the accretion disc spectra of Dwarf Novae and Novalike variables:
implications for the standard disc model

Gabriella Zsidi,1,2,3 C. J. Nixon,1 T. Naylor4 and J. E. Pringle5
1School of Physics and Astronomy, Sir William Henry Bragg Building, Woodhouse
Ln., University of Leeds, Leeds LS2 9JT, UK
2Konkoly Observatory, HUN-REN Research Centre for Astronomy and Earth
Sciences, Konkoly-Thege Miklós út 15-17, 1121 Budapest, Hungary
3CSFK, MTA Centre of Excellence, Budapest, Konkoly Thege Miklós út 15-17,
1121, Hungary
4Department of Physics and Astronomy, University of Exeter, Stocker Road,
Exeter EX4 4QL, UK
5Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge,
CB3 0HA, UK E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

Accretion discs are fundamental to much of astronomy. They can occur around
stars both young and old, around compact objects they provide a window into
the extremes of physics, and around supermassive black holes in galaxy centres
they generate spectacular luminosities that can outshine the entire galaxy.
However, our understanding of the inner workings of accretion discs remains
far from complete. Here we revisit a conundrum in the observations of some of
the simplest accreting systems; the Cataclysmic Variables (CVs). The high-
accretion-rate states of (non-magnetic) CVs can be divided into the short-
lived outbursts ($\sim$ a week) typical of dwarf novae (DNe) and the long-
lived (and sometimes perpetual) high states of nova-like (NL) CVs. Since both
sorts of high-state occur in approximately steady-state accretion discs with
similar properties and accretors, we would expect them to display similar
spectral energy distributions. However, previous analyses based on UV spectra
from the International Ultraviolet Explorer have shown that their spectral
energy distributions are different. We perform a re-analysis of the data using
up to date calibrations and distance (and thus dereddening) estimates to test
whether this difference persists and whether it is statistically significant
over the sample. We find that it does persist and it is statistically
significant. We propose routes to investigating this discrepancy further and
discuss the implications this has for other accreting systems, such as X-ray
binaries, Active Galactic Nuclei and protoplanetary discs.

###### keywords:

accretion, accretion discs – stars: dwarf novae – (stars:) novae, cataclysmic
variables – ultraviolet: stars

††pubyear: 2024††pagerange: Revisiting the accretion disc spectra of Dwarf
Novae and Novalike variables: implications for the standard disc
model–Revisiting the accretion disc spectra of Dwarf Novae and Novalike
variables: implications for the standard disc model

##  1 Introduction

Accretion discs occur in many astrophysical systems as a consequence of the
conservation of angular momentum in material that is falling due to the
gravitational pull of an accretor (Lynden-Bell, 1969; Pringle & Rees, 1972;
Shakura & Sunyaev, 1973; Pringle, 1981). Discs form around protostars in star
forming regions (e.g. Armitage, 2011), around main-sequence stars (e.g.
Rivinius et al., 2013), due to mass transfer in binary systems (e.g. Warner,
1995) and around the supermassive black holes in galaxy centres (AGN; Lynden-
Bell, 1969).

The basics of the standard model of accretion discs are laid out in Pringle &
Rees (1972) and Shakura & Sunyaev (1973), and summarised in Pringle (1981).
The disc matter orbits the central object in a sequence of planar and circular
orbits until arriving at the inner boundary, the location of which is set by
the physics of the central object and, in some cases, the properties of the
disc itself. At each radius in the disc, the material must give up some
orbital energy to move to a smaller radius. This is effected by viscous
torques within the disc that cause angular momentum to be transferred outwards
and mass inwards, and the net energy transfer results in heating of the disc
material. This heat is radiated from the disc surface and this is what we see
from the accretion disc.

A basic assumption in the standard model is that the accretion energy is
radiated locally from the disc in the region in which it is produced. If this
is the case, and assuming a steady supply of mass from large radii and that
the disc surface radiates as a blackbody, the disc temperature profile takes a
simple form that is independent of the details of the viscous process that
produces the energy dissipation. The temperature profile is (e.g. Pringle,
1981)

$T(R)=\left\\{\frac{3GM{\dot{M}}}{8\pi R^{3}\sigma_{\rm
SB}}\left[1-\left(\frac{R_{\rm in}}{R}\right)^{1/2}\right]\right\\}^{1/4}$ (1)

where ${\dot{M}}$ is the accretion rate on to the central object of mass $M$
and radius $R_{\rm in}$, and $\sigma_{\rm SB}$ is the Stefan-Boltzmann
constant. An implication of this simple equation is that, for a central
accretor of the same mass and radius accreting at the same rate, we should
expect to see basically the same spectrum emitted from the disc.

In many cases the environment and dynamics of accretion flows are complex and
this can make direct comparisons between models and observational data
difficult. However, as our understanding of the inner workings of accretion
discs is not yet complete, it is important to test the basic assumptions that
are made in disc models where we can. To keep things as simple as possible, we
need to find systems where the accretion flow is approximately time-steady (at
least in the inner regions where most of the energy is produced), with simple
central objects, and surface temperatures in the regime where we have a good
understanding of the emitted spectra. These properties are satisfied by some
cataclysmic variable (CV) stars, particularly the dwarf novae (DNe) and the
novalike variables (NLs).111Here we restrict our attention to the “non-
magnetic” cases, where there is no evidence that the central object hosts a
sufficiently strong magnetic field as to interfere with the accretion flow.

CVs are close binaries, in which material is transferred from a Roche-lobe
filling secondary to a white dwarf primary through the L1 point (Warner,
1995). For non-magnetic CVs, where the magnetic field strength of the white
dwarf is insufficient to significantly perturb the accretion flow, the
transferred material forms an accretion disc that extends down to the surface
of the white dwarf (Lubow & Shu, 1975). The subclass of CVs known as DNe
undergo a thermal-viscous instability cycle, during which the accretion disc
displays two different accretion regimes: (1) quiescent states in which the
disc viscosity parameter $\alpha$ is low and mass accumulates in the disc over
time and (2) outburst states in which the disc viscosity parameter is high and
mass drains on to the central white dwarf (e.g. Meyer & Meyer-Hofmeister,
1981; Faulkner et al., 1983; Smak, 1984; Warner, 1995; Lasota, 2001).222In the
quiescent states we do not have a good handle on the magnitude of the Shakura
& Sunyaev (1973) viscosity parameter $\alpha$ with estimates of
$\alpha\lesssim 0.01$. In the outburst states $\alpha$ is measured to be
around 0.3 (e.g. Kotko & Lasota 2012; see the summaries in King et al. 2007;
Martin et al. 2019). During the outbursts, which typically last for around a
week, the disc dominates the electromagnetic output of the system. NLs do not
exhibit DNe outbursts,333They can exhibit brightening events or flares, but
these are mostly attributed to mass transfer variations from the companion
star (see, e.g., Zsidi et al. 2023 and references therein). and instead
exhibit luminosities that are commensurate with the outburst states of the DNe
for extended periods of time. The accepted explanation for the difference
between these two sub-classes of CV is that the NLs have a high enough mass
transfer rate from the companion star to exist permanently in the high state,
whereas the DNe, having smaller mass transfer rates, must suffer the limit
cycle behaviour induced by the thermal-viscous instability. Naively one may
therefore expect the NLs to appear bluer than the DNe as the disc temperature
grows slightly with increasing ${\dot{M}}$. But it is also possible that, as
the NLs can be intrinsically brighter, we may typically observe them at larger
distances making them more vulnerable to interstellar reddening.

To explore the differences between the DN and NL classes, la Dous (1991)
undertook a statistical analysis of the ultraviolet (UV) spectra of 32 DNe and
23 NLs using data from the International Ultraviolet Explorer (IUE) satellite,
and found that the NLs appear to be somewhat redder444By “redder” we mean a
larger ratio of long-wavelength to short-wavelength flux. than the DNe in
outburst. More recently, Hamilton et al. (2007) fit the IUE DNe outburst
spectra with the synthetic accretion disc spectra of Wade & Hubeny (1998),
finding that the synthetic spectra provide a good fit for almost all of the
DNe outburst spectra. From this we can surmise that the standard disc model is
a good description of the accretion process in DNe. However, the same cannot
be said for NLs. Following initial studies on individual systems by Wade
(1984, 1988), la Dous (1991) used a large sample of CVs to show that IUE
spectra of NLs are redder than the DNe outburst spectra, and thus are not well
fit by the synthetic models (see also Puebla et al., 2007). Subsequent models
that provide a good fit with NL spectra (from IUE and also FUSE and HST) are
provided by Long et al. (1994), Linnell et al. (2007a, b), Linnell et al.
(2010) and Godon et al. (2017). These models artificially increase the inner
radius of the disc to several times the size of the white dwarf; this has the
effect of reducing the peak temperature of the inner disc, making the spectrum
redder as a result (some models artificially cap the temperature at a maximum
value or change the radial power-law of the temperature profile instead of
moving the inner boundary). In agreement with la Dous (1991), Godon et al.
(2017) explore the disc-dominated IUE spectra of CVs and find that NL spectra
are typically redder than the DNe in outburst.

In this paper, we aim to build on the analyses presented by la Dous (1991) and
Godon et al. (2017), by performing a reanalysis of the IUE spectra taking
account of updated calibrations of the IUE pipeline, updated distances that
allow for more accurate dereddening of the spectra, and to perform statistical
tests of the samples to firmly establish the significance of the difference in
the colour of the IUE spectra of DNe and NLs. In Section 2 we discuss the IUE
data, sample selection and dereddening calculations. In Section 3 we present
our analysis of the colours of DNe and NLs. In Section 4 we present discussion
of the results and we conclude and provide implications for other systems in
Section 5.

##  2 Archival IUE observations

In this work we use data from IUE, which carried out observations between 1978
January and 1996 September. The spacecraft was equipped with a short-
wavelength prime (SWP: $1150-2000\,\AA$), a long-wavelength prime (LWP:
$1850-3400\,\AA$) and a long-wavelength redundant camera (LWR:
$1850-3400\,\AA$), and could provide both high- and low-dispersion spectra
with a resolution of 0.2 $\AA$ and 6 $\AA$. In practice, the LWR provided
better quality data and thus was generally used in preference to the LWP. To
maintain a consistent dataset we consider only the low-dispersion spectra as
these are more numerous.

The original spectroscopic image processing system, IUESIPS, was released in
1977, and it removed instrumental signatures and provided wavelength and flux
calibration. This data reduction pipeline has, however, evolved throughout the
years as users discovered certain shortcomings. This introduced
inconsistencies to the IUE data base (Nichols & Linsky, 1996). After IUE was
deactivated in 1996, the entire IUE database was reprocessed in order to
compile a high-quality and uniformly processed IUE Final Archive. These
reprocessed data products are called NEWSIPS (New Spectroscopic Image
Processing System; Nichols et al., 1994; Nichols & Linsky, 1996). New image
processing techniques were developed with this procedure, which included e.g.
the improvement of the photometric calibration, new models for flux
calibration, and also uniform wavelength calibration. The new procedure also
resulted in an increase in the S/N of $10-50$ per cent for the low dispersion
spectra. A detailed comparison between the IUESIPS and the NEWSIPS data
products with example spectra can be found in Nichols & Linsky (1996).

The NEWSIPS data showed significant improvement compared to the data products
from the IUESIPS pipeline. However, comparison between IUE low-dispersion
data, models, and Hubble Space Telescope (HST) observations revealed some
problems with the absolute flux calibration, the thermal corrections, and
time-dependent sensitivity (Massa & Fitzpatrick, 2000). In certain cases, this
could result in discrepancies as large as 10%. Massa & Fitzpatrick (2000)
provide the correction files and corresponding IDL
routine.555https://archive.stsci.edu/iue/contrib.html In Fig. 1, we show an
example spectrum of TT Ari using the IUESIPS pipeline (employed by la Dous
1991), using the NEWSIPS data, and a “Corrected NEWSIPS” which include the
Massa & Fitzpatrick (2000) corrections.

Figure 1: SWP10131 (top) and SWP10614 (bottom) spectra of TT Ari. The black
lines show the IUESIPS spectra, the orange lines show the NEWSIPS spectra, and
the green lines show the NEWSIPS spectra after applying the Massa &
Fitzpatrick (2000) correction. This shows that the UV spectral slope can be
affected by the new IUE data reduction pipeline (NEWSIPS) and, to a lesser
extent, the Massa & Fitzpatrick (2000) correction.

The target acquisition of IUE was carried out by the Fine Error Sensor (FES),
which also recorded brightness information of the observed system. The raw FES
counts are often available from the file headers for bright sources, which can
be used to estimate the visual brightness of the system (FES magnitudes)666See
section 4.7. of the IUE observing guide for further details:
https://archive.stsci.edu/iue/instrument/obs_guide/ using the calibration of
Pérez (1991). When using this calibration we ignore the colour corrections as
they are small for DNe in outburst and NLs which have $B-V\approx 0$ (Warner,
1995).

###  2.1 Sample selection

Our starting point was the sample presented in la Dous (1991). la Dous (1991)
studied the IUE spectra obtained between 1978 and 1987, and her sample
included 21 DNe and 11 NLs, whose inclination is $i\leq 60^{\circ}$ (when
using the updated system parameters) and both short and long wavelength
spectra were available. We collected updated system parameters for each
object, which we list in Tables 1 and 2. In order to restrict our attention to
systems in which an accretion disc is present, clearly observable, and extends
to the white dwarf surface, we only include non-magnetic systems and those
with low inclination ($i\leq 60^{\circ}$).

We discarded EK TrA from our sample due to its unlikely large distance of
6650${}^{+2820}_{-2199}$ pc (Bailer-Jones et al., 2021), and SW UMa due to its
likely magnetic nature (Shafter et al., 1986; Dubus et al., 2018). We did not
find contemporaneous photometric observations for V794 Aql, but its faint
spectra suggest that it was probably in low state. There were also no
available light curve for V425 Cas from the American Association of Variable
Star Observers (AAVSO) data base, but Szkody (1985) reports that the system
was in a mid-state during a decline from maximum to minimum. For these
reasons, we removed V794 Aql and V425 Cas from our sample. IUE carried out
further observations until 1997, which allows us to increase the number of
objects in our sample. When selecting further sources from the IUE archive, we
followed the selection criteria by la Dous (1991). Namely, we cross-correlated
the JD of the IUE spectra with light curves in order to ensure that the
sources we include are DNe in outburst and NLs in high state. The AAVSO data
base provides the most complete collection of light curves from the times of
the IUE observations but, unfortunately, these light curves do not always
provide contemporaneous information on the source brightness. We only included
those new DNe, which we are confident are in outburst, and this gave us an
additional 6 DNe. The light curves of the NLs are very sparse, and we
therefore attempted to use the FES counts to estimate the system brightness.
Many of these systems were too faint in the optical for the FES and the
acquisition was made in ‘Blind Offset’ mode, which means that no FES counts
were recorded for the target. Other objects were either in low state or are
high inclination eclipsing systems, so our search for new NLs matching our
criteria did not return any additional NLs. Therefore, our total sample
consists of 28 DNe and 8 NLs, as V426 Oph is classified as a DNe according to
more recent catalogues (e.g. Ritter & Kolb, 2003). In all cases, we used the
low dispersion large aperture observations, and we chose datasets where the
typical time lag between the observations carried out by the short-wavelength
(SWP) and the long-wavelength (LWP or LWR) camera is approximately half an
hour (or in rare cases up to a few hours).

###  2.2 Distances and Dereddening

We determined individual extinctions to each CV using the extinction maps of
Vergely et al. (2022) which are derived using measured extinctions and Gaia
parallaxes of about 35 million stars. This required distances to each CV, for
which we used the geometric distances given by Bailer-Jones et al. (2021)
based on the Gaia parallaxes. We do not use the photogeometric distances since
the prior for these distances includes assumptions about the absolute
magnitude of the star which are inappropriate for peculiar objects such as
CVs. It is normal when using Gaia parallaxes to exclude stars where the
astrometric fits to the individual observations are poor, expressed as a high
renormalized unit weight error (RUWE, see for example, Fabricius et al. 2021).
However we did not apply any such quality cut because it appears that CVs have
anomalously poor summary statistics for the quality of their astrometric fits,
which are not indicative of poor parallaxes (see Section 2.1 of Pala et al.,
2020). We suspect this is due to their variability.

Using these distances we obtained extinctions to each CV by integrating
through the 10 pc resolution Vergely et al. (2022) data cubes from the Sun to
the CV. The Vergely et al. (2022) extinction cubes are created using the
Fitzpatrick et al. (2019) extinction law, which is qualitatively different
from previous work in that it works in monochromatic extinctions rather than
extinctions over a given filter bandpass. Hence integrating through the
Vergely et al. (2022) extinction cubes yields the extinction at 550 nm (only
roughly equivalent to $A_{V}$) which we then converted to extinction at a
given wavelength using the standard Galactic value of $R(55)$ (again the rough
equivalent $R_{V}$) of 3.02. Four objects (RZ Gru, VY Scl, VZ Scl and DW UMa)
are all sufficiently far from the Galactic Plane that they lie above or below
the $\pm$400 pc thickness of the Vergely et al. (2022) models, but given the
small scale height of the Galactic dust layer (around 100 pc, e.g. Jones et
al., 2011) and their small extinctions (0.04 mags) integrating out to the edge
of the model should be sufficient for our purposes.

Table 1:  List of DNe analysed in this work. The table includes the object
names, revised extinction values, the representative SWP and LWR/LWP spectrum
spectra, inclination, DNe subtype, distance, orbital period, FES magnitude,
outburst state, the time since the start of the outburst (where known), and
the brightness state of the source at the time the spectra were taken.

Object | $A_{55}$ | SWP | LWR | LWP | Inclination | Subtype | Distance [pc] | $P\,(\textrm{d})$ | FES | Outburst state | $\Delta t\,(\textrm{d})$ | Brightness state
---|---|---|---|---|---|---|---|---|---|---|---|---
RX And | 0.054 | 17673 | 13931 | - | 51.0$\pm$9∘ | ZC | 196.60 | 0.209893 | 11.1 | Peak outburst | 0 | Peak brightness
AR And | 0.104 | 18877 | 14885 | - | 40.0$\pm$10∘ | UG | 409.59 | - | 12.4 | Early decline | 1 | Peak$-0.5$ mag
VY Aqr | 0.043 | 21720 | - | 02367 | 41.0∘ | WZ | 141.13 | - | 11.6 | Decline | 6 | Peak$-1$ mag
Z Cam | 0.028 | 18630 | 14697 | - | 57.0$\pm$11∘ | ZC | 213.50 | 0.289841 | 10.0 | Peak / early decline | 1 | Peak$-0.5$ mag
SY Cnc | 0.055 | 08082 | 07051 | - | 56.7$\pm$20∘ | ZC | 401.25 | 0.382375 | 11.5 | Decline | 6 | Peak$-0.5$ mag
YZ Cnc | 0.040 | 03727 | 03308 |  | 38.0$\pm$3∘ | SU | 233.49 | 0.0868 | 11.9 | Early decline | 6 | Peak$-$1 mag
HL CMa | 0.175 | 19873 | 15857 | - | 45.0$\pm$10∘ | ZC | 293.54 | 0.216787 | 12.0 | Outburst | - | -
WW Cet | 0.044 | 10664 | 09373 | - | 54.0$\pm$9∘ | ZC | 218.06 | 0.1758 | 12.2 | Outburst | - | -
SS Cyg | 0.050 | 01788 | 01675 | - | 37.0$\pm$5∘ | UG | 112.35 | 0.27513 | 8.6 | Wide outburst | 5777Number of days since rise. | Peak brightness
AB Dra | 0.253 | 17619 | 13886 | - | 40.0$\pm$10∘ | UG | 396.61 | - | 12.7 | Late rise | -1 | Peak$-0.5$ mag
VW Hyi | 0.019 | 24307 | - | 04652 | 60.0$\pm$3∘ | SU | 53.74 | 0.074271038 | 9.0 | Super outburst | 2888Number of days since rise. | Peak brightness
WX Hyi | 0.036 | 07509 | 06500 | - | 40.0$\pm$10∘ | SU | 228.42 | 0.074813 | 11.6 | Super outburst | 0 | Peak brightness
X Leo | 0.072 | 15951 | 12282 | - | 41.0$\pm$10∘ | UG | 438.11 | - | 12.6 | Outburst | - | -
AY Lyr | 0.088 | 09342 | 08098 | - | 41.0$\pm$10∘ | SU | 415.55 | - | 13.2 | Super outburst999With precursor outburst. | 5 | Peak brightness
CY Lyr | 0.307 | 21030 | 16779 | - | 60.0$\pm$10∘ | UG | 468.33 | - | 13.2 | Peak outburst | 2 | Peak brightness
CZ Ori | 0.165 | 07358 | 06376 | - | 18.0$\pm$10∘ | UG | 499.97 | 0.2189 | 12.4 | Decline | 4 | Peak$-0.5$ mag
RU Peg | 0.105 | 15062 | 11594 | - | 41.0$\pm$7∘ | UG | 271.30 | 0.3746 | 11.2 | Mid-decline | 7 | Peak$-1.5$ mag
TZ Per | 0.629 | 17643 | 13907 | - | 56.7$\pm$20∘ | ZC | 456.95 | - | 12.9 | Outburst | 6 | Peak brightness
KT Per | 0.287 | 17712 | 13968 | - | 60.0$\pm$10∘ | UG | 244.66 | 0.162656 | 12.5 | Peak outburst | 0 | Peak brightness
UZ Ser | 0.601 | 17633 | 13901 | - | 50.0$\pm$10∘ | ZC | 308.83 | - | 12.7 | Outburst | - | -
VW Vul | 0.425 | 18875 | 14883 | - | 41.0$\pm$10∘ | ZC | 542.92 | - | 13.9 | Outburst | - | -
AH Her | 0.096 | 41939 | - | 20704 | 46.0$\pm$3∘ | ZC | 327.91 | 0.258116 | 12.3 | Early decline101010Outburst with slow rise. | 3 | Peak$-0.8$ mag
DX And | 0.312 | 37689 | - | 16844 | 45.0$\pm$12∘ | UG | 585.37 | 0.440502 | 12.6∗ | Wide outburst | -1 | Peak$-0.3$ mag
IR Gem | 0.054 | 38524 | - | 17696 | 15.0$\pm$10∘ | SU | 257.83 | 0.0684 | 12.3 | Peak outburst | 3 | Peak brightness
SU UMa | 0.040 | 34824 | - | 14532 | 43.0$\pm$10∘ | SU | 220.14 | 0.07635 | 12.8 | Early decline | <1 | Peak$-0.5$ mag
TT Crt | 0.077 | 47801 | - | 25677 | 60.0$\pm$10∘ | UG | 532.41 | 0.26842 | … | Early decline? | - | -
V1159 Ori | 0.303 | 56781 | - | 31957 | 35.0$\pm$10∘ | SU | 351.17 | 0.062178 | … | Outburst | 11 | Peak brightness
V426 Oph111111V426 Oph was classified as a NL in la Dous (1991), but according to more recent catalogues (e.g. Ritter & Kolb, 2003) based on more recent observations, this object is a DN. | 0.243 | 29546 | - | 09427 | 59.0$\pm$6∘ | ZC | 190.01 | 0.2853 | 11.7 | Outburst | - | Peak brightness

Notes. The distances are the geometric distances from the Bailer-Jones et al.
(2021) catalogue based on Gaia measurements. The FES magnitudes flagged with
an asterisk indicate that the brightness estimated from the FES counts differs
by more than 0.5 mag from the brightness suggested by the AAVSO light curve.
Where an FES magnitude is not available this is indicated with an ellipsis.
The DNe subtypes are: U Gem (UG), SU UMa (SU), Z Cam (ZC), WZ Sagittae (WZ).
7Number of days since rise.
8Number of days since rise.
9With precursor outburst.
10Outburst with slow rise.
11V426 Oph was classified as a NL in la Dous (1991), but according to more
recent catalogues (e.g. Ritter & Kolb, 2003) based on more recent
observations, this object is a DN.

Table 2:  List of NLs analysed in this work. The table includes the object
names, revised extinction values, the representative SWP and LWR/LWP spectra,
inclination, NL subtype, distance, orbital period, and the FES magnitude.

Object | $A_{55}$ | SWP | LWR | LWP | Inclination | Subtype | Distance [pc] | $P\,(\textrm{d})$ | FES
---|---|---|---|---|---|---|---|---|---
HL Aqr | 0.199 | 23325 | - | 03647 | 18.0∘ | UX | 567.09 | 0.13557 | 13.4
TT Ari | 0.152 | 06278 | 05448 |  | 30.0$\pm$10∘ | VY | 247.02 | 0.13755 | 10.9
KR Aur | 0.213 | 14734 | 11299 | - | 38.0∘ | VY | 458.89 | 0.162772 | 13.1
KQ Mon | 0.105 | 16267 | 12516 | - | <60∘ | UX | 612.67 | - | 12.6
VY Scl | 0.042 | 32594 | - | 12365 | 30.0∘ | VY | 595.61 | 0.2323 | 12.7
RW Sex | 0.060 | 07500 | 06494 | - | 34.0$\pm$6∘ | UX | 221.63 | 0.245145 | 10.6
IX Vel | 0.027 | 18578 | 14655 | - | 57.0$\pm$2∘ | UX | 90.31 | 0.193927 | 9.3
BZ Cam | 0.214 | 21251 | 16940 | - | 15.0∘ | VY | 369.64 | 0.15353 | 12.6

Notes. The distances are the geometric distances from the Bailer-Jones et al.
(2021) catalogue based on Gaia measurements. Where an FES magnitude is not
available this is indicated with an ellipsis. The NL subtypes are UX UMa (UX)
and VY Scl (VY). Novalikes may also be classified based on analysis of the
emission lines in their spectra; from the above list HL Aqr, TT Ari, and KQ
Mon have been classified as also belonging to the SW Sex (SW) subtype (see
Mizusawa et al. 2010, Ritter & Kolb 2003 and Wolfe et al. 2013 respectively).

Figure 2: The light curve and the spectra of SS Cyg. The top panel shows the
AAVSO light curve with the time of the spectroscopic observation indicated
with a vertical line. The bottom panel shows the IUE SWP and LWR spectra of SS
Cyg. The grey shaded regions indicate the bands that we chose to obtain
$C_{\textrm{short}}$ and $C_{\textrm{long}}$.

###  2.3 Final sample

After collecting the NEWSIPS IUE spectra from MAST for the extended sample, we
applied the corrections determined by Massa & Fitzpatrick (2000), and we
dereddened each spectrum using the Fitzpatrick et al. (2019) extinction law.
All sources, CV class and subclass,121212We provide the subclasses based on
the Ritter & Kolb (2003) catalogue, which was last updated in 2011. IUE
spectrum ID, extinction values, inclination, Gaia distances, orbital periods,
and the FES magnitudes are listed in Tables 1 and 2. The FES magnitudes were
calculated primarily based on the FES counts on the target (‘CTS’), which is
included in the headers of the SWP fits files. When either the CTS or the
‘focus step’ was not available in the SWP files, we estimated the FES
magnitude from the LWR or the LWP files. We note that the FES magnitudes
provide only an estimate on the system brightness. We indicated with asterisks
those FES magnitudes, which differ by more than 0.5 mag from the brightness
estimated from the AAVSO light curves.

##  3 Results

###  3.1 Colours – definition

To study and compare the UV spectra of DNe and NLs, we analysed the dereddened
spectra in the following way. We first smoothed them by calculating the mean
flux in a 9 pixel wide sliding window (box filter). We then chose three bands
centred on $1460\,\AA$ with a width of $80\,\AA$, $1765\,\AA$ with a width of
$30\,\AA$, and $2630\,\AA$ with a width of $60\,\AA$. We chose these central
wavelengths and band widths so that we can minimize the contamination of any
absorption or emission lines in the spectra, and so that they are also far
enough from the edges where the spectra tend to be noisier. We chose two bands
in the SWP spectra, because the effect we want to examine is expected to be
more significant at shorter wavelength. We show an example spectrum with the
selected bands in the bottom panel of Fig. 2. We integrated the flux in each
band for all of the objects in our sample. As a final step, we defined a short
and a long wavelength colour by taking the negative logarithm of the ratio of
the neighbouring bands, i.e.:

$C_{\rm
short}=-\log\left(\frac{\int_{\rm\lambda=1420\AA}^{1500\AA}f_{\rm\lambda}\,\rm
d\lambda}{\int_{\rm\lambda=1750\AA}^{1780\AA}f_{\rm\lambda}\,\rm
d\lambda}\right)=-\log\left(\frac{F_{1460}}{F_{1765}}\right)$ (2)

and

$C_{\rm
long}=-\log\left(\frac{\int_{\rm\lambda=1750\AA}^{1780\AA}f_{\rm\lambda}\,\rm
d\lambda}{\int_{\rm\lambda=2615\AA}^{2645\AA}f_{\rm\lambda}\,\rm
d\lambda}\right)=-\log\left(\frac{F_{1765}}{F_{2630}}\right).$ (3)

Figure 3: Colour-colour, $C_{\rm short}$ vs. $C_{\rm long}$, plot for the DNe
and the NLs. The red symbols are the NLs. The blue symbols are the DNe, with
triangles marking the Z Cam systems (which can show prolonged standstills),
crosses marking the SU UMa systems (which can show superoutbursts) and circles
marking the remaining DNe subtypes. We have indicated the size of the
representative errorbar with the black cross in the top left corner of the
figure. As described in the text the dominant source of scatter in this plot
is intrinsic astrophysical variations of the sources which appears to be at
the $\approx 0.02$ level in these quantities, with uncertainties in the data
analysis at the level of $\approx 0.01$. The intensity of the colour indicates
the inclination of the system as shown in the colour bar (see Tables 1 & 2 for
details). This figure shows that the DNe and NLs are not drawn from the same
parent population.

The uncertainties in the flux ratios defined in this way are small.
Propagating the uncertainties from NEWSIPS through our analysis gives median
uncertainties of 0.005 in $C_{\rm short}$ and 0.012 in $C_{\rm long}$, with a
largest in either band being 0.018. These are comparable to, but smaller than,
the differences we obtain between repeated measurements of the nova-likes IX
Vel and TT Ari. The RMS between the $C_{\rm short}$ values of the 28 IX Vel
low-resolution large-aperture SWP spectra is 0.011. For the 37 high-state
spectra of TT Ari the RMS is 0.016.

Figure 4: Histogram of the $C_{\textrm{short}}$ values for the DNe (blue bars)
and NLs (red bars). It is clear that the two distributions are different.
Figure 5: Histogram of the $C_{\textrm{long}}$ values for the DNe (blue bars)
and NLs (red bars). Again it can be seen that the two distributions are
distinct, but with less clarity than in Figure 4.

###  3.2 Initial comparisons

In Fig. 3, we plot $C_{\rm short}$ against $C_{\rm long}$ in order to provide
an initial comparison of the UV spectral slopes of the DNe and the NLs. The
DNe and NLs are shown in different colours. It is immediately evident that the
two sets of objects do not appear to have been drawn from the same parent
population. To make this more apparent we compare in Fig. 4 histograms of the
$C_{\rm short}$ for the DNe and the NLs, and in Fig. 5 the same for the
$C_{\rm long}$. It is clear from the comparisons of the $C_{\rm short}$
distributions (Fig. 4), and to a lesser extent the $C_{\rm long}$
distributions (Fig. 5) that the UV spectra of the NLs in the high state are
generally redder than the DNe in outburst. Although it should be noted that
there is some overlap in the two sets of distributions. Note, also, that the
$C_{\rm short}$ and the $C_{\rm long}$ are not independent datasets.

Because of the paucity of data points (in particular for the NLs) it is not
appropriate to make use of statistical tests (for example $\chi^{2}$) that
rely on the assumption of underlying Gaussian distributions or approximations
thereto (for example by application of the central limit theorem).
Furthermore, as explained in Section 3.1, the precise level of the
uncertainties for individual datapoints is unclear. We therefore make use of
non-parametric statistics, and in particular the Kolmogorov-Smirnov (KS) test.
Application of the 1D two sample KS test to the data in Fig. 4 ($C_{\rm
short}$) shows that the probability that the null hypothesis (that both
samples are drawn from the same parent population) being valid is less than or
equal to 1.2 per cent. The same test applied to the data in Fig. 5 ($C_{\rm
long}$) yields a probability of less than or equal to 3.8 per cent. Typically,
the threshold for not rejecting the null hypothesis that the two samples are
drawn from the same distribution is 5 per cent. It seems clear therefore that
the conclusion of la Dous (1991), that there is a significant difference
between the UV spectra of DNe in outburst and of NLs in the high state, is
upheld by our more detailed analysis which makes use of the most recent
instrumental calibrations and also takes careful account of likely
differential reddening between the DNe and the (on average) more distant NLs.

Figure 6: This figure shows the same plot as that in Fig. 3, but here with the
names of the systems included for reference.

###  3.3 A more detailed comparison

We here briefly look at the dataset in more detail. We construct a modified
version of the $C_{\rm short}$ vs. $C_{\rm long}$ plot (Fig. 3) in which we
indicate the position of each object together with the name of the system
(Fig. 6). We can see from this plot that most of the DNe that were outliers
(i.e. too red) in the study of la Dous (1991) or were not fitted well (again
largley because they were too red compared with the models) according to
Hamilton et al. (2007) are the ones whose positions lie in the region of
overlap with the NL sample. These systems are the following: AB Dra, WW Cet,
RX And, VW Vul, WX Hyi, YZ Cnc, Z Cam, DX And. There is also one source, WX
Hyi, which lies in a region of Fig. 6 that would suggest that good fit should
be possible, but Hamilton et al. (2007) find a poor fit. Comparing the IUE
spectra used by Hamilton et al. (2007) (SWP23952) and the one used here
(SWP07509) we find that the latter is at a significantly higher flux level
(around a factor of 5 in the wavelength range $1420-1500$ Å) and thus we
expect that the spectrum used by Hamilton et al. (2007) for this source was
not near the peak of the outburst.

Puebla et al. (2007) modelled the UV spectra of a sample of CVs, which
included 23 nova-like variables. They found that accretion-model spectra could
not reproduce the observed spectra of the NLs, and concluded that almost all
the models produce a bluer spectrum than the observations. This was the case
despite allowing the system parameters ($M$, $\dot{M}$ and $i$) to be free
parameters, which could lead to lower $\dot{M}$ and would mean that NLs also
undergo disc instability. However, observations do not show disc instability
in NLs. In addition they tried allowing the distance to be a free parameter in
order to match the flux level of the observations. It was the most poorly
constrained parameter in their modelling procedure, and more accurate distance
estimates now available from the Gaia catalogue, reveal that the distance
estimates from synthetic spectral fitting are lower by a factor of up to six
than those measured by Gaia (Gilmozzi & Selvelli, 2024). Our results,
according to which the NLs spectra are redder than the DNe spectra, also
suggest that the standard disc model will not reproduce the observed UV
spectra of NLs. This finding is in agreement with Puebla et al. (2007), who
proposed that a revision of the temperature profile is needed to model the UV
spectra more accurately, and this is particularly important for the innermost
part of the disc. Using the most up-to-date Gaia distances will allow more
accurate estimates of the system parameters, e.g. those estimated by Gilmozzi
& Selvelli (2024), and new efforts can be made to improve the modelling of the
NL spectra.

Godon et al. (2017) studied the entire IUE CV sample, including both magnetic
and non-magnetic sources. They determined the slope of the dereddened UV
spectra using the wavelength range between ${\sim}1600\,$Å and $3200\,$Å, and
measured the slope on a log-log scale. They found that the UV slope of the
different disc-dominated CV subtypes decreases in the following order: SU UMa,
U Gem, Z Cam type DNe, then UX UMa, and VY Scl type NLs. The location in Fig.
6 of the four UX UMa type NLs (HL Aqr, KQ Mon, RW Sex, IX Vel) and the four VY
Scl type NLs (TT Ari, KR Aur, VY Scl, BZ Cam) in our sample reveals that our
results are not inconsistent with the findings of Godon et al. (2007), though
our sample is small. However, we must note that Godon et al. (2007) studied
the entire CV sample of IUE (including systems with both low and high
inclination), therefore, their results should not be discussed only in the
context of intrinsic differences in accretion discs. They also note that the
observed UV continuum emission of the SW Sex type systems is probably affected
by their high inclination. In comparison, we focused on systems in which the
disc dominates the UV spectrum (i.e. $i\leq 60^{\circ}$), therefore, the
differences we observe between the different CV classes and subclasses can be
attributed to differences in their discs.

In Fig. 6 we have distinguished the different types of DNe. In particular we
note here the ten Z Cam type objects which are included in the analysis. These
objects exhibit both DNe and NL-like characteristics; the continuous series of
outbursts, which resemble those of standard DNe, are occasionally interrupted
by the decline from outburst being halted by a long standstill, during which
the the system brightness remains constant for a few weeks/months at a
brightness level of around $0.5-1$ mag fainter that the peak outburst level.
The standstills seem to be analogous to the behaviour of NLs which remain in a
high state for long periods (cf. Meyer & Meyer-Hofmeister, 1983).

Examining Fig. 6, we can see that the Z Cam stars whose colours overlap the
majority of the other DNe are SY Cnc, HL CMa, TZ Per, AH Her and UZ Ser
(although we note that UZ Ser is the bluest source by a significant margin).
From inspection of the AAVSO lightcurves (not depicted), we can be confident
that all of these sources are near the maximum of DNe outbursts at the time
the IUE spectra were taken, with perhaps the exception of TZ Per and HL CMa
for which short standstills cannot be ruled out. Z Cam is on the boundary
between the DNe group and the NL group, and from the AAVSO lightcurve appears
to be in the DNe phase at the time the IUE spectrum was taken. VW Vul has a
similar $C_{\rm long}$ to the DNe, but a larger $C_{\rm short}$; from the
AAVSO lightcurve the state of VW Vul is inconclusive, but may be during a
standstill (although at a slightly higher flux level than previous standstills
which may indicate a DNe state). V426 Oph has the opposite to VW Vul with a
similar $C_{\rm short}$ to the DNe and a similar $C_{\rm long}$ to the NLs;
similar to other sources above the optical lightcurves indicate that V426 Oph
was probably in a DNe outburst state, but there insufficient data to be
definitive. The Z Cam stars with colours most discrepant from the DNe group
are WW Cet and RX And. For RX And we note that there is another IUE spectrum
taken a day later that has an increased flux level (around 25 per cent higher)
and appears bluer indicating that the spectrum we used may not have been near
the peak of the outburst. For WW Cet, the AAVSO lightcurve is inconclusive and
it is possible that the source was in a standstill state.

Given the possibility that some of the Z Cam stars were not necessarily in the
DNe outbursting state, we also performed the 1D KS test using a reduced DNe
sample excluding the Z Cam objects. This resulted in even smaller
probabilities for the validity of the null hypothesis that the two samples
(DNe and NLs) were drawn from the same parent population, being less than or
equal to 0.3 per cent for $C_{\rm short}$ and less than or equal to 1.2 per
cent for $C_{\rm long}$. This result highlights that Z Cam stars may be an
important class of stars for testing the evolution of the UV spectra with
time, particularly through a prolonged standstill.

##  4 Discussion

We have reanalysed the UV spectra obtained by IUE for a sample of non-magnetic
CVs comprising DNe in outburst and NLs that are sufficiently face-on to the
line-of-sight that we have a good view of the accretion disc. Taking account
of the NEWSIPS IUE data processing, the Massa & Fitzpatrick (2000)
corrections, and updated distance and interstellar reddening estimates, we
find that – in agreement with the original analysis of la Dous (1991) and a
more recent analysis by Godon et al. (2017) – the NLs have disc spectra that
are significantly redder than the DNe. This is consistent with the extensive
modelling efforts of the community that consistently find that (1) the
outburst spectra of DNe can be fit by the standard disc model (Hamilton et
al., 2007) and (2) the spectra of NLs cannot be fit by the standard model
(Wade, 1988; Long et al., 1994; Linnell et al., 2007a, b; Puebla et al., 2007;
Godon et al., 2017). Our analysis substantiates that this difference cannot be
accounted for by either interstellar reddening or statistical effects. We are
therefore able to conclude that there is an intrinsic physical difference
between these two very similar disc systems. Understanding this difference is
likely to lead to a deeper understanding of the accretion process in general.

We note that, because the DNe in outburst are well-fit by the standard model,
it is natural to assume that it is the NLs that are the “outliers” and must
harbour some additional or alternative physical disc process. There is of
course no guarantee that this assumption is correct. In this context we here
draw attention to a couple of complementary ideas which may contain the germs
of an explanation.

Nixon & Pringle (2019) note that a physical difference between DNe in outburst
and NLs is the time spent in the high accretion state (see also Tout et al.,
1993). They present a model in which the disc, on entering the high accretion
state, initially has only small-scale surface magnetic fields allowing for a
local viscosity (i.e. the standard model). On timescales longer than the DNe
outbursts, i.e. longer than at least a week or so, the lengthscale of the
surface magnetic field structures grows, allowing for non-local transport of
the accretion energy. Nixon & Pringle (2019) suppose this occurs in the
surface layers of the disc with mass transported along poloidal field lines to
larger radii in what they call a Magnetically Controlled Zone (MCZ). They
present simple calculations that show how this might affect the disc
temperature profile, and are able to redistribute the energy in the disc in a
way that may account for the difference in the observed spectra of DNe and
NLs.

There have long been concerns about modelling accretion disc spectra (see for
example Pringle, 1992), and in particular at what height in the disc the
accretion energy is likely to be generated. The standard model, referred to
above, makes the assumption that most of the accretion energy is deposited at
high optical depth, on the grounds that most of the energy is likely to be
deposited where most of the mass resides (Shakura & Sunyaev, 1973). However,
this may well not be the case especially if the main viscous process is
magnetic, since in that case the accretion energy appears in the form of
magnetic energy which can then be dissipated elsewhere. Drawing on the ideas
of Kriz & Hubeny (1986), Shaviv & Wehrse (1986), Hubeny (1990) and Puebla et
al. (2007) (see also the discussion by Pringle 1992),131313Similarly, Hassall
(1985) experiments with models that include an optically thin component
finding that it can be a dominant component in the modelling of EK TrA. Hubeny
& Long (2021) provide models that can provide better fits to the NL spectra in
which the accretion energy is deposited as a function of height through the
disc, rather than all in the optically thick disc midplane region. This leads
to emission from optically thin regions, which results in both a better
representation of the UV slope and to filling in the Balmer jump, the depth of
which is usually over-predicted in the models. The drawback with the modelling
of Hubeny & Long (2021) is that it does not provide any physical motivation
for why the energy deposition would occur predominantly in the mid-plane in
the DNe outbursts (as implied by the standard model fits) and predominantly at
larger vertical heights in the NLs (their modelling of the NL IX Vel finds
that 99 per cent of the energy is dissipated at the disc surface).

It is plausible that the modelling of Hubeny & Long (2021) provides a good
description of the observational implications of the Nixon & Pringle (2019)
model. This might be expected because the Nixon & Pringle (2019) model uses
large scale magnetic field structures (in the NLs) to expel mass and deposit
it at larger radii resulting in additional deposition of energy locally to the
larger radii (the change in UV slope is effected by the lack of additional
energy deposition at small radii in the disc for which there is no smaller
radii to generate the surface mass flow). How the additional energy deposited
by the surface flow is distributed through the vertical column of the disc
depends on the mass flow rate in the MCZ and the surface density and
temperature structure at the location at which mass is deposited. It is
plausible that in some conditions the expelled matter landing on the disc at
larger radii results in shocks that are confined to the surface layers, which
could be commensurate with the model of Hubeny & Long (2021) for IX Vel.

We note that Hartley et al. (2002) also proposes the presence of a large scale
magnetic field in the NL discs, but they came to this conclusion by following
a different approach. Their time-resolved spectroscopic observation of the
nova-like variable IX Vel do not reveal significant variability on short
timescales (from $\sim$10 to $\sim$100 s), and they conclude that the line-
driven wind cannot be responsible for the observed level of mass-loss rate in
this system. They propose that a large scale ordered magnetic field could be
responsible for the observed level of mass-loss rate, if the streamline
geometry required by the magnetic field controls the mass-loss rate in these
systems. However, the large scale magnetic field that would be required to
drive the wind need not to be the same as the one required for the MCZ, and we
have not yet been able to address whether these are part of the same magnetic
field structure.

###  4.1 Investigating the time-dependence of the spectra from TT Ari

If these ideas are in any way relevant, it would suggest that the differences
found by la Dous (1991), and substantiated here, between the UV spectra of DNe
in outburst and NLs in the high state are caused by some difference in the
magnetic structures in these discs, which thereby give rise to differences in
where and how the accretion energy is deposited and turned into radiation. The
obvious difference between the discs in outbursting DNe and those in NLs is
the length of time that the discs have been in the high state (Tout et al.,
1993; Nixon & Pringle, 2019). In this picture the time taken for the disc to
reach a steady state (the viscous timescale) is less than the time taken for
the magnetic configuration to reach an equilibrium structure.

Assuming the major difference in the spectra arises from the time spent in the
high accretion state, and that this leads to magnetic structures as proposed
by Nixon & Pringle (2019), then we can estimate the timescales we might expect
in the following manner. The radius of a $0.6M_{\odot}$ white dwarf is $R_{\rm
wd}\approx 0.01R_{\odot}$, and the orbital timescale at a radius of, say,
$10R_{\rm wd}$ is $2\pi\sqrt{R^{3}/GM}\approx 300$ s. The magnetic dynamo
timescale is typically ten times the orbital timescale (e.g. Tout & Pringle,
1992; Stone et al., 1996), and so at $10R_{\rm wd}$ this is $\approx 3000$ s,
and so we expect the generation of small-scale fields with lengthscale $\sim
H$ at the disc surface on timescales of order an hour. If there is a diffusive
process to transfer magnetic field structures from size of order $H$ to size
of order $R$, then the number of steps required in such a process is around
$(R/H)^{2}$ (King et al., 2004), and here with $H/R\approx 0.03$ we have
$(R/H)^{2}\approx 10^{3}$. Thus, we estimate that the total time required to
set up organised field on a scale of order $10R_{\rm wd}$ might be $\sim
3\times 10^{6}$ s, which is about a month. This is consistent with the split
between DNe outbursts which last about a week and show standard disc spectra
and the NLs which have been in the high accretion state for times much longer
than a month and have spectra that require modification from the standard
model. Even VY Scl type NLs, which spend some time in low accretion state,
would have enough time to form a large scale magnetic field a few weeks/months
after they returned to the high state. This is also supported by our result
that all of the NL variables are redder than the average dwarf nova (Fig. 6).
However, if we were to obtain UV spectra of a VY Scl type system immediately
after it has reached the maximum brightness, we might observe a bluer spectrum
than a few months later.

With such ideas in mind, Tout et al. (1993) studied the IUE spectra of TT Ari,
a VY Scl type NL. These type of CVs spend extended time in a high state, which
is occasionally interrupted by periods when the system brightness drops by
several magnitudes. These low states may last several weeks to months. With
the aim of examining the temporal evolution of a disc going through a cycle of
high state and low state and then returning to the high state again, Tout et
al. (1993) measured the slope of the UV spectra by defining a ‘blueness’
parameter ($B_{\rm Q}$) in the following way141414Note that our $C_{\rm
short}$ and $C_{\rm long}$ definitions are distinct, being both logged and
inverted in comparison.,

$B_{\rm Q}=\frac{\int_{\rm\lambda=1318.25\AA}^{1380.38\AA}f_{\rm\lambda}\,\rm
d\lambda}{\int_{\rm\lambda=1819.7\AA}^{1905.46\AA}f_{\rm\lambda}\,\rm
d\lambda}.$ (4)

Their results, which we have reproduced in the top panel of Fig. 7, suggests
that the UV spectrum of TT Ari is ‘bluer’, and therefore hotter, after the
system returned to the high state from the brightness minimum (see black
symbols in the top panel of Fig. 7).

Table 3: Observing log of the TT Ari spectra we used in Fig. 7.  SWP | Date and time of observation | $\textrm{JD}-2\,440\,000$
---|---|---
10041 | 1980-09-07 09:27:17 | 4489.89
10042 | 1980-09-07 10:54:22 | 4489.95
10043 | 1980-09-07 12:17:58 | 4490.01
10044 | 1980-09-07 13:36:28 | 4490.07
10131 | 1980-09-14 23:38:56 | 4497.49
10614 | 1980-11-17 15:25:37 | 4561.14
10944 | 1981-01-02 04:00:50 | 4606.67
10945 | 1981-01-02 06:21:50 | 4606.77
10946 | 1981-01-02 07:39:05 | 4606.82
11035 | 1981-01-12 08:40:37 | 4616.86
13398 | 1981-03-03 04:33:59 | 4666.69
Figure 7: Top panel: The blueness parameter defined by Tout et al. (1993) for
various TT Ari spectra, with black crosses corresponding to the IUESIPS data
used by Tout et al. (1993) and green circles corresponding to the Corrected
NEWSIPS data used here. The tick marks on the top axis show the beginning of
each month. This figure shows that the re-analysis of the TT Ari data yields
no significant change in the colour for these spectra. Bottom panel: The same
as the top panel, but this time showing the $C_{\rm short}$ (filled circles)
and $C_{\rm long}$ (open circles) UV colours we employ in this paper. Again,
there is no strong evidence for the evolution of the colours in these spectra.

However, with the aim of examining the impact of the new data reduction
procedure on the scientific results, we re-analysed the IUESIPS low dispersion
spectra used by Tout et al. (1993) and employed the data analysis methodology
outlined in Section 2 to produce the Corrected NEWSIPS low dispersion SWP
spectra of TT Ari. We summarised the log of observations in Table 3. As the
bottom panel of Fig. 1 suggests, the new data reduction introduced changes in
some of the spectra. Therefore, we calculated the $B_{\rm Q}$ blueness
parameter as defined above using both the old and the new datasets. The
results suggest, in contrast to the analysis of Tout et al. (1993) on the
IUESIPS data, that there was no significant change in the blueness parameter
for the Corrected NEWSIPS data over the period analysed by Tout et al. (1993).

We must note, however, that the minimum Tout et al. (1993) studied was very
short and not well defined. Information on the system brightness from this
period of time is available from the AAVSO database and we plot it in Fig. 8.
These optical photometric data suggest that the system was indeed in a fainter
state when some of the IUE observations were carried out, but it is unclear
from these data when the minimum started. Furthermore, this minimum was not as
deep as other low states of TT Ari. The FES magnitudes also confirm that the
system was in a fainter state ($\sim$14 mag) for one of the IUE observations,
but not as faint as during a regular low state ($\sim$16 mag). Additionally,
in the bottom panel of Fig. 7 we provide the $C_{\rm short}$ and $C_{\rm
long}$ values for the Corrected NEWSIPS spectra over this period and find that
vary in the range expected for the NLs based on the distributions in Fig. 3.
All in all, based on our re-analysis of the tentative results presented by
Tout et al. (1993) and our interpretation of the lightcurve of TT Ari around
the time these spectra were taken, we can conclude that at present there is no
significant data pointing towards an evolution of the colour of the spectra in
these sources on a specific timescale. This question will need to be revisited
in the future.

Figure 8: Top: The long-term optical (AAVSO) lightcurve of TT Ari for the
entire IUE observing period. The black vertical lines indicate the epochs at
which SWP spectra were obtained, the red dotted lines mark the epochs of the
LWP/LWR spectra. We plotted the red dotted line over the black solid line for
the epochs when both SWP and LWP/LWR spectra were taken. The gray shaded area
shows that region of the lightcurve which is enlarged in the bottom panel.
Bottom: The AAVSO lightcurve highlighting the time at which the spectra in
Fig. 7 were taken. As can be seen in the figure the minimum prior to the time
at which the spectra were taken is both shallow and short, and the recovery
did not reach the long-term high state. As such it is clear that this was not
a good minimum on which to attempt this analysis.

##  5 Conclusions and Implications

Through a reanalysis of the UV spectra of non-magnetic DNe in outburst and NLs
from the IUE satellite we have confirmed the findings of la Dous (1991), and
others, that the NL spectra are significantly redder than the DNe in outburst.
This is an important discrepancy in our understanding of accretion discs. A
physical explanation for this difference may generate significant
understanding of the accretion process in a wide range of systems. We have
provided discussion on the physical mechanisms that may be responsible and
suggestions for future observations that may help to disentangle this mystery.

If, as we suggest, strong magnetic field structures are able to grow on the
surfaces of the NL discs on the timescales indicated above, then this has
significant implications for the behaviour of discs in other accreting sources
and the application of the standard model to them. From the arguments in
Section 4.1, the timescale to substantially modify the accretion process
through growth of surface magnetic field structures is $t_{\rm
MCZ}\sim(R/H)^{2}t_{\rm dynamo}$, where again we shall take $t_{\rm
dynamo}=ft_{\rm orb}$ with $f\approx 10$ and, as above, estimate the timescale
at 10 stellar radii.151515Note that the timescale here is
$\propto(R/H)^{2}t_{\rm orb}$ and scales in a way that is similar to the
standard viscous timescale. So if the process is somehow not related to the
disc dynamo, but instead nonetheless related to the viscous timescale, the
numbers given here are still relevant. We consider discs in X-ray binaries,
AGN, and protostellar/protoplanetary discs.

The implications for X-ray binaries are clear as these discs are broadly
similar to those in CVs, but with substantially reduced inner radii, and thus
timescales in the inner regions, due to the more compact accretor. In this
sense the inner regions of the disc are much “older” in terms of number of
orbits than the CV discs. The disc thicknesses are also similar with
$H/R\approx 0.01$. Here we have, as an illustration, a central black hole
accretor of mass, say, $10M_{\odot}$ and thus gravitational radius $R_{\rm
g}=GM/c^{2}\approx 15$ km. The orbital timescale at $10R_{\rm g}$ is therefore
$\approx 0.01$ s, and therefore the timescale $t_{\rm MCZ}$ is about a quarter
of an hour.

For active galactic nuclei (AGN) the time- and length-scales are significantly
larger, but the systems are also much longer lived. Scaling to a central black
hole of mass $10^{7}M_{\odot}$ we have a gravitational radius of $\approx 0.1$
au and an orbital timescale at $10R_{\rm g}$ of just under 3 hours. Employing
the AGN structure derived by Collin-Souffrin & Dumont (1990), from which we
have $H/R\approx 0.002$ we get the timescale as $t_{\rm MCZ}\sim 10^{3}$ yrs.
As the lifetime of AGN discs is around $10^{5}-10^{6}$ yrs (e.g. Schawinski et
al., 2015; King & Nixon, 2015), $t_{\rm MCZ}$ in the inner disc regions is
only a small fraction of the typical disc lifetime.

For protoplanetary discs we have $H/R\approx 0.1$ (e.g. Bell et al., 1997), a
central star of mass $\sim 0.5M_{\odot}$ and radius of, say, $2R_{\odot}$. In
this case, the orbital timescale at 10 stellar radii is around 2 weeks. The
dynamo timescale is then around half a year, and the timescale $t_{\rm MCZ}$
is tens of years. Again, this constitutes a small fraction of the $\sim$ Myr
lifetime of these discs. However, the ionisation structure of protoplanetary
discs is complex, with large regions of the disc attaining sufficiently low
levels of ionisation that magnetic activity can be shut off. We might
therefore expect the effect to manifest most strongly in the innermost regions
of protoplanetary discs, on scales $\lesssim 0.1-1$ au, which are thought to
remain sufficiently ionised by the radition field from the central star for
magnetic activity to operate continuously (e.g. Armitage, 2011). Additionally,
while protoplanetary discs are thought to spend most of their lifetimes in a
low-viscosity state, at least some objects spend time in high viscosity states
known as FU Ori outbursts. These can last longer than decades, and so may
exhibit evolution in their surface magnetic structures similar to that
inferred for the NLs.

In summary, if the dichotomy between the spectra of the DNe in outburst and
the NLs is determined by the timescale for which these sources have been in
the high accretion (fully ionised) state, and the critical timescale scales in
the way we have described above, then almost all other discs in astronomy
behave like the NLs rather than the DNe.161616It is worth noting that the
observations of Her X-1 provided by Kosec et al. (2020) provide evidence for
the gravitationally bound, surface outflows postulated by the Nixon & Pringle
(2019) model (see Nixon & Pringle, 2020). The conclusion of this line of
thought is that almost all fully-ionised discs in the universe are non-
standard in this manner. The discrepancy in spectra between DNe in outburst
and NLs is therefore a significant unsolved question in our understanding of
accretion discs, and addressing this question is likely to yield valuable
information on the dynamical processes that govern disc evolution in a wide
range of astrophysical systems.

Finally, we suggest that some useful next steps are:

  1. 1.

To connect the Nixon & Pringle (2019) and Hubeny & Long (2021) modelling
efforts to provide predictions for the time evolution of the UV spectra of CV
discs upon (re)entering the high accretion state.

  2. 2.

To obtain improved UV spectra, both in terms of quality and quantity, of both
DNe in outburst and NLs in the high state to improve the statistics presented
here.

  3. 3.

And to catch discs that either enter a prolonged outburst (including the
superoutbursts exhibited by SU UMa stars and the standstills exhibited by Z
Cam stars) or that are recovering from a deep minimum (the VY Scl type stars)
to explore the time-dependence of the UV spectra.

For example, Hubble Space Telescope spectra would be highly valuable as they
would facilitate the analysis at shorter wavelengths than has been possible
with the IUE data here. From Figs. 9 & 15 of Linnell et al. (2007b) and Fig. 4
of Nixon & Pringle (2019), we anticipate that the effect is more pronounced at
shorter wavelengths for which the emission originates from the central regions
of the disc.

## Acknowledgements

GZs and CJN acknowledge support from the Leverhulme Trust (grant number
RPG-2021-380). CJN acknowledges support from the Science and Technology
Facilities Council (grant number ST/Y000544/1). This work has made use of data
from the European Space Agency (ESA) mission Gaia
(https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and
Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium). This research is based
on observations made with the International Ultraviolet Explorer, obtained
from the MAST data archive at the Space Telescope Science Institute, which is
operated by the Association of Universities for Research in Astronomy, Inc.,
under NASA contract NAS 5–26555. We acknowledge with thanks the variable star
observations from the AAVSO International Database contributed by observers
worldwide and used in this research.

## Data Availability

The data underlying this article are available in the MAST and the AAVSO
archives.

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