text:
\section{Introduction}
Cataclysmic Variables (CVs) are interacting binary systems, each consisting of a white dwarf accreting material from a Roche-lobe-filling companion. The material flowing from the companion forms an accretion disk around the white dwarf before being dumped on its surface.
In the case of a highly-magnetized white dwarf ($B > 10^{6}$ G), the magnetic field of the white dwarf truncates the inner regions of the disk or even completely prevents it from being formed. In this case the material follows the magnetic field lines onto the surface of the white dwarf. These are known as magnetic CVs (see \citealt{Warner_1995} for a review).
CVs experience several types of cataclysmic events and thermonuclear explosions, hence the name. One of these events is called a dwarf nova (DN) and is a viscosity-induced instability in the accretion disk, resulting in a temporary increase in mass transfer rate and heating the whole disk. DNe manifest as relatively low amplitude outbursts, typically $\Delta m \sim 2 - 5$ mag up to $\sim$ 9\,--\,10 mag in some extreme cases \citep{Kawash_etal_2021a}.

While material builds up on the surface of the white dwarf through secular accretion, the pressure and density increases in its surface layers. Once a critical density is reached, a thermonuclear runaway is triggered on the surface of the white dwarf, leading to an increase in the brightness of the system by up to 15\,mag or more in a matter of a few days \citep{Starrfield_etal72,Yaron_etal_2005}. These events are known as classical novae (see \citealt{Gallaher_etal_1978,Bode_etal_2008,Woudt_Ribeiro_2014,Della_Valle_Izzo_2020,Chomiuk_etal_2020} for reviews) and their recurrence timescale is typically $\gtrsim$ thousands of years \citep{Yaron_etal_2005}. In some cases, novae recur on shorter timescales -- short enough to be recorded more than once. These are called recurrent novae and these systems are usually characterized by a high mass transfer rate, often due to the presence of an evolved secondary \citep{Schaefer_2010}.

After a classical nova event, the mass transfer rate is expected to be high enough to keep the disk in a hot, ionized state, temporarily preventing DNe in the system.
DNe are only expected thousands of years later, when the mass transfer rate becomes low enough for the disk to cool and again become susceptible to disk instability events. This is implied by the hibernation scenario of CVs \citep{Shara_etal_1986}, which suggests that CVs go through cycles of low and high states of mass transfer rate. In this scenario, the states are mostly determined by the irradiation of the secondary by a nova event (which increases the mass transfer rate) and the growing separation between the two stars (which decreases the mass transfer rate). Only a few CV systems have shown DN outbursts after classical nova eruptions;
e.g.,  GK~Per \citep{Sabbadin_Bianchini_1983,Bianchini_etal_1986}, V1017~Sgr \citep{Sekiguchi_1992,Salazar_etal_2017}, and V446~Her \citep{Honeycutt_etal_1995,Honeycutt_etal_2011}.

\subsection{V1047~Cen -- the 2019 outburst}
V1047~Cen (Nova Cen 2005) was discovered as a Galactic transient on 2005 September 1.03 at a visual magnitude of around 8.5 (left panel of Figure~\ref{Fig:LC_comp}) and later classified spectroscopically as a classical nova eruption \citep{Liller_2005}.
\citet{Walter_etal_2012} reported spectroscopic follow up taken approximately 5 and 7 days after discovery. The spectra showed
typical lines of Balmer, \eal{Fe}{II}, and \feal{O}{I}, which are characteristic of a classical nova near optical peak.
The Balmer lines showed multiple P Cygni absorptions with blueshifted velocities of around 750 and 1800\,km\,s$^{-1}$.

Otherwise, V1047 Cen was not extensively observed during the 2005 eruption and
little else is known about the classical nova eruption.
Archival observations of the system taken by the Inner Galactic Disk with MIPS (MIPSGAL; \citealt{Carey_etal_2009}) Survey in the 24 and 70 micron bands, yield a 24-micron average magnitude of 0.72 $\pm$ 0.02\,mag \citep{Gutermuth_etal_2015}.
The MIPSGAL survey data were collected between September 2005 and October 2006.
Additional archival observations by the VISTA Variables in the Via Lactea (VVV) survey between March 2010 and August 2011, show the nova fading from $K$=12.84 to $K$=13.11 \citep{Saito_etal_2012}.
The Neil Gehrels Swift Observatory (hereafter \textit{Swift}; \citealt{Gehrels_etal_2004}) observed the classical nova eruption between 2005 and 2008 \citep{Ness_etal_2007}. The observations obtained in November 2005 and January 2006 resulted in detections of hard X-rays from strongly absorbed shock emission, which are typically seen in novae during the early weeks/months of the eruption (see, e.g., \citealt{Mukai_etal_2008,Schwarz_etal_2011,Gordon_etal_2021}). The \textit{Swift} observations obtained in 2008 led to a non-detection.

Fourteen years after the 2005 nova eruption, \citet{2019TNSTR.985....1D} reported the discovery of an astronomical transient AT2019hik/Gaia19cfn possibly associated with V1047~Cen on 2019 June 11.6, with a discovery magnitude of 16.2 in the $G$-band. Based on regular monitoring by the Optical Gravitational Lensing Experiment (OGLE; \citealt{Udalski_etal_2015}) survey, \citet{ATel_12876} confirmed that the FK5 J2000 equatorial coordinates of the transient ($[\alpha, \delta]$ = [13$^{\mathrm{h}}$20$^{\mathrm{m}}$49$^{\mathrm{s}}$\!.78, --62$^{\circ}$37$'$50\arcsec\!\!.6]) were consistent with those of V1047~Cen and found that the re-brightening of the system started as early as 2019 April 6.11 (HJD 2458579.61; right panel of Figure~\ref{Fig:LC_comp}). This date is considered as the outburst start ($t_0$) throughout this paper. \citet{ATel_12876}
noted that the slow re-brightening of the system is inconsistent with a recurrent nova eruption. The re-brightening of the system triggered follow-up observations across the electromagnetic spectrum. \citet{ATel_12885} reported optical spectroscopy, which showed DN outburst spectral features superimposed on spectral features of a classical nova nebula. \citet{Geballe_etal_2019} reported infrared spectroscopy of the then ongoing outburst, during its first 160 days, concluding that the event was possibly a DN outburst. However, upon further follow-up, the 2019 outburst of V1047~Cen seems to be a more exotic phenomenon.

Here we report on multi-wavelength follow-up of the 2019 outburst of V1047~Cen spanning the spectrum from X-ray to radio.
In Section~\ref{Obs} we present the observations and data reduction. In Section~\ref{Res} we show our results, while in Section~\ref{Disc} we offer discussion about the nature of the event and its peculiar observational features. Our conclusions are given in Section~\ref{Conc}.

\begin{figure*}
\begin{center}
  \includegraphics[width=0.49\textwidth]{nova_2005.pdf}
  \includegraphics[width=0.49\textwidth]{outburst_2019.pdf}
\caption{A direct comparison between the visual light curve of the 2005 classical nova eruption from AAVSO (\textit{left}) and the OGLE $I$-band light curve of the 2019 outburst (\textit{right}). The remarkable difference between the two light curves indicates that the two events are of a completely different nature.}
\label{Fig:LC_comp}
\end{center}
\end{figure*}

\section{Observations and data reduction}
\label{Obs}
\subsection{Optical and near-IR photometry}
V1047~Cen has been observed by the OGLE sky survey \citep{Udalski_etal_2015} since May 2013, several years before the outburst, as part of the OGLE Galaxy Variability Survey (GVS). All data were taken in the $I$ band with an exposure time of 30~s, and they were reduced and calibrated using the standard OGLE pipeline \citep{Udalski_etal_2015}. A sample of the observations is listed in Table~\ref{table:OGLE_log}.

Optical photometry between days 76 and 437 was performed in the $BVRI$ bands by several observers from the American Association of Variable Star Observers (AAVSO; \citealt{Kafka_2020}). The bulk of the multi-band photometry comes from two observers, Gordon Meyer and Josch Hambsch. A sample of the observations is listed in Table~\ref{table:AAVSO_log}.

We obtained SMARTS {\it Andicam} photometry in the $J, H$, and $K_s$
bands on 29 nights between days 80 and 117.
Data reduction is described in \cite{Walter_etal_2012}. A sample of the observations is listed in Table~\ref{table:SMARTS_log}.
The {\it Andicam} instrument was retired on 2019 August~1 (day 117). We also make use of IR photometry from the enhanced Wide-field Infrared Survey Explorer (NEOWISE; \citealt{Mainzer_etal_2011}) covering only three epochs during the outburst in the $W1$ (3.35\,$\mu$m) and $W2$ (4.60\,$\mu$m) bands. A log of the observations is listed in Table~\ref{table:NEOWISE_log}.

The field of V1047~Cen was observed by the Transiting Exoplanet Survey Satellite (\textit{TESS}; \citealt{Ricker_etal_2015}) during sector 11, which covers the early rise of the 2019 outburst between days 17 and 44. We used the open-source tool ELEANOR \citep{Feinstein_etal_2019} to extract light curves from the \textit{TESS} full-frame images,
opting to utilize the corrected flux light curve and to only include data that are not associated
with a quality flag.

All the data will be available as online material. All multi-band photometry is calibrated using the
\textit{Vega} system zero points.

\begin{figure*}
\begin{center}
  \includegraphics[width=\textwidth]{multi_plot.pdf}
\caption{Multi-band light curves of the 2019 outburst of V1047~Cen. The \textit{top panel} shows optical and IR photometry from AAVSO ($BVRI$), OGLE ($I$), SMARTS ($JHK$), and NEOWISE ($W1$ and $W2$) photometry. The \textit{middle panel} shows \textit{Swift} UVOT UV data, and the \textit{bottom panel} shows MeerKAT L-band radio flux at 1.28 GHz. Colors and symbols are coded as indicated in figure legends.}
\label{Fig:optical_UV}
\end{center}
\end{figure*}

\subsection{Optical and IR spectroscopy}

We obtained optical and IR spectroscopic observations of the 2019 outburst between days 74 and 643 using a diverse group of telescopes and instruments. A log of the spectral observations is presented in Table~\ref{table:spec_log}.

We obtained spectra using the High Resolution Spectrograph (HRS; \citealt{Barnes_etal_2008,Bramall_etal_2010,Bramall_etal_2012,Crause_etal_2014}) and the Robert Stobie Spectrograph (RSS; \citealt{Burgh_etal_2003}; \citealt{Kobulnicky_etal_2003}) mounted on the Southern African Large Telescope (SALT; \citealt{Buckley_etal_2006,Odonoghue_etal_2006}) in Sutherland, South Africa. HRS was used in the low resolution LR mode, yielding a resolving power $R \approx$ 14,000 over the range 4000\,--\,9000\,\AA. The primary reduction of the HRS spectroscopy was conducted using the SALT science pipeline \citep{Crawford_etal_2010}, which includes over-scan correction, bias subtraction, and gain correction. The rest of the reduction was done using the MIDAS FEROS \citep{Stahl_etal_1999} and $echelle$ \citep{Ballester_1992} packages. The reduction procedure is described by \citet{kniazev_etal_2016}. RSS was used in long-slit mode with the 1.5 arcsec slit and the PG900 grating, resulting in a resolving power $R \approx 1500$. The spectra were first reduced using the  PySALT pipeline \citep{Crawford_etal_2010}, which involves bias subtraction, cross-talk correction, scattered light removal, bad pixel masking, and flat-fielding. The wavelength calibration, background subtraction, and spectral extraction were done using the Image Reduction and Analysis Facility (IRAF; \citealt{Tody_1986}).

We also carried out low- and medium-resolution optical spectroscopy using the Goodman spectrograph \citep{Clemens_etal_2004} on the 4.1\,m Southern Astrophysical Research (SOAR) telescope located on Cerro Pach\'on, Chile. The observations were carried out in two setups: one setup using the 400~l\,mm$^{-1}$ grating and a 0.95\arcsec\ slit, yielding a resolving power $R \approx$ 1000 over the wavelength range 3820--,850\,\AA. Another setup was used with a 2100~l\,mm$^{-1}$ grating and a 0.95\arcsec\ slit, yielding a resolving power $R \approx$ 5000 in a region centered on H$\alpha$ that is 570\,\AA\ wide. The spectra were reduced and optimally extracted using the \textsc{apall} package in IRAF.

Four high-resolution spectra were also obtained using the Chiron fiber-fed echelle spectrograph \citep{Tokovinin_etal_2013} mounted on the CTIO 1.5m telescope.
Integration times were 20 minutes, with three integrations per night summed for 1 hour net exposure time. All spectra were taken in ``fiber mode'', with
4x4 on-chip binning yielding a resolving power $\lambda/\delta\lambda \approx$27,800. The data were reduced using a pipeline coded in IDL (Walter 2017)\footnote{\url{http://www.astro.sunysb.edu/fwalter/SMARTS/CHIRON/ch_reduce.pdf}}.

\subsection{Infrared spectroscopy}

Near-infrared spectra of V1047 Cen were obtained at the Gemini South Telescope on Cerro
Pachon in Chile on days 223 and 230
using the facility instrument FLAMINGOS-2, for program GS-2020A-Q-201. The
observations are summarized in Table~\ref{table:IR_spec},. The 0.36'' wide slit and R3000 grism were employed to
obtain spectra at resolving powers, $R$, of up to 3000 in portions of the $J, H,$ and $K$ windows. The JH R1200 grism was used with the same slit to obtain a spectrum covering 0.89-1.75 $\mu$m. Note
that for each grism there is considerable variation in $R$ across each wavelength interval. (see Ref.\footnote{\url{http://www.gemini.edu/instrumentation/flamingos-2/components\#Grisms}}). Data reduction
employed standard near-infrared techniques utilizing both Gemini/IRAF and Figaro commands \citep{Shortridge_etal_1992}. Flux calibrations, derived from spectra of the standard stars listed in
Table~\ref{table:IR_spec} and are accurate to $\pm$30\%.

NASA SOFIA \citep[Stratospheric Observatory for Infrared Astronomy;][]{Young_etal_2012} airborne observations of V1047~Cen were obtained with the the Faint Object InfraRed  CAmera \cite[FORCAST;][]{Herter_etal_2018}, the dual-channel mid-infrared imager and grism spectrometer operating from 5 to 40~\micron{} on two separate, consecutive flight series on 2019 July 02.542 UT
(F0589) 03.559 UT
(F0590), days 88 and 89, using the G111 and G227 grisms with the instrument configured using a long-slit (4\farcs7 $\times$ 191\arcsec), yielding a spectral resolving power $R = \lambda/\Delta\lambda \sim$ 140--300. the G111 grating was used to observe V1047~Cen both nights. Standard pipeline processed and flux calibrated data \citep[for details of the reduction processes see][]{Clarke_etal_2015} were retrieved for science analysis from the Infrared Processing and Analysis Center (IPAC) Infrared Science Archives (IRSA). The data products contain a computed atmospheric transmission model appropriate for the flight altitudes which were used to mask-out spectral points in the observed spectral energy distributions where the atmospheric transmission was $\lesssim$ 70\%. Within the statistical errors, no variations between the G111 data sets obtained on the two different nights were noted, hence these data were averaged into a single resultant spectrum.

\subsection{MeerKAT observations}
We observed the field of V1047 Cen with the MeerKAT radio interferometer \citep{2016mks..confE...1J} 15 times between days 277 and 700.
Observations were taken with the 64 antenna array at a central frequency of $1.28\,\rm{GHz}$ with a $856\,\rm{MHz}$ bandwidth. Each observation consisted of 15 minutes integration on the field of V1047 Cen with two minutes on the phase calibrator J1424--4913 before and after. J1939--6342 was used to set the flux scale and calibrate the bandpass response of the instrument. Data were reduced using the OxKAT (see \citealt{Heywood_2020oxkat} for details) reduction scripts, which include recipes for both phase reference and amplitude and phase self-calibration. The typical noise in a region of our images without obvious emission is 35\,$\mu\rm{Jy}$/beam. A log of the MeerKAT observations is given
in Table~\ref{table:radio_log}

\subsection{\textit{Swift} observations}
Observations of V1047 Cen with \textit{Swift} commenced on 2019 June 24, 79 days after the re-brightening start. The UV/Optical
Telescope (UVOT; \citealt{Burrows_etal_2005,Roming_etal_2005}) detected emission across all three UV filters ($uvw1$: central wavelength
of 2600 \AA; $uvm2$: 2246 \AA; $uvw2$: 1928 \AA; a sample of the observations is listed in Table~\ref{table:Swift_log}). No X-ray emission was
detected with the X-ray Telescope (XRT), however. A second observation was performed a month later (day 109), followed by approximately weekly observations from day 123 until day 420 after the start of the outburst.
Throughout this
time, a variable UV source was detected, while no individual observations
showed significant X-ray emission. Co-adding the full $\sim$50 ks of XRT data, a faint X-ray source with a count rate of 2.8$^{+1.3}_{-1.1}$~$\times$~10$^{-4}$ count~s$^{-1}$ was detected.
However, with only 16 counts in the source extraction region, no spectral
analysis can be sensibly performed.

\section{Results}
\label{Res}
\subsection{Optical/UV light curves}

The AAVSO ($BVRI$), OGLE $I$, and \textit{Swift}  UV light curves of V1047 Cen are presented in Figure~\ref{Fig:optical_UV}. The data
reveal peculiar behavior,
with the outburst lasting for around 400 days. There is a general trend of increased brightness between days 0 and $\sim$ 310, before the brightness of the system starts decreasing. Throughout the 400 day outburst, the light curves show variability of the order of a few days and with amplitude variations of $\lesssim$ 1\,magnitude.
Based on the OGLE data the amplitude of the outburst exceeded
6 magnitudes (Figure~\ref{Fig:LC_comp}).

The evolution of the optical broadband colors $(B-V)_0$, $(V-R)_0$, and $(R-I)_0$ is presented in Figure~\ref{Fig:color_evolution}. We applied a reddening correction based on the $E(B-V)$ extinction derived in Section~\ref{red_dist}. The colors show random fluctuations around a mean value throughout the majority of the outburst.
However, on day 250, the three colors show a noticeable redward trend, coincident with a bump in the optical light curves. This is followed by a blueward trend after day 300, as the outburst ends.

We have near-IR photometry only for a short period near the start of the plateau.
Aside from a trend toward redder colors as the source brightened (between days 76 and 82, corresponding to late part of the rising phase), the
near-IR colors were fairly stable, at $I-K=2.25 \pm 0.03$,
$J-K=1.24 \pm 0.02$, and $H-K=0.51 \pm 0.03$.

Using the optical and UV photometry, we created spectral energy distributions (SEDs) during different stages of the outburst (the rise -- from start to day 90; the plateau -- between days 90 and 300; and the decline -- after day 300; Figure~\ref{Fig:SED}). We fit blackbody functions to the SEDs to estimate the temperature of the emitting source. The best-fit blackbody models fit the data poorly, indicating complex emission, possibly from multiple sources (e.g, the nebula of the 2005 nova event, the white dwarf surface, the accretion disk, and ejected material/outflow). Therefore, caution is required when interpreting a blackbody fit. That being said, throughout the outburst, the temperatures derived from the best fit blackbody models are in the range of 17000\,K to 21000\,K.
These are somewhat high compared to the temperatures modeled
for DN outbursts in the outer part of the disk ($\sim$ 10000\,K \citealt{Godon_etal_2020}) but is in the right range for the temperatures in the inner part of the disk (between 30000 and 60000\,K \citealt{Long_1996,Godon_etal_2017}). In GK~Per outburst, \citet{Kim_etal_1992} estimated an accretion disk temperature of around 6000 to 10000\,K in the outer part of the disk based on their disk-instability-induced outburst model. The $(B-V)_0$ colors of V1047 Cen range between $-0.3$ and $+0.3$, with an average value of $\approx$ 0.0 throughout the outburst (Figure~\ref{Fig:color_evolution}), which translates to a blackbody temperature between 7000\,K and 16000\,K, with an average blackbody temperature of around 10000\,K. These temperatures are consistent with those expected
in a disk instability event.
We do not have simultaneous $BVRIJHK$ and UV data, but in one single instance (day $\approx$ 308), we have simultaneous optical, UV, and NEOWISE NIR photometry. We create the SED from these observations and we present the best fit blackbody model, which shows a similar outcome to the above conclusions (Figure~\ref{Fig:SED_extended}). However, what is most important about this SED is that it shows how poorly the blackbody model fits the data, implying that any blackbody temperatures derived from the SEDs or even color indices are not reliable.

\begin{figure}
\begin{center}
  \includegraphics[width=\columnwidth]{color_evolution.pdf}
\caption{The evolution of the extinction-corrected broadband colors $(B-V)_0$, $(V-R)_0$, and $(R-I)_0$, throughout the outburst.}
\label{Fig:color_evolution}
\end{center}
\end{figure}

\begin{figure}
\begin{center}
  \includegraphics[width=\columnwidth]{SED.pdf}
\caption{Extinction-corrected SED plots, showing the evolution of the SED throughout different phases of the outburst. The error bars are 1$\sigma$ uncertainties and they include contributions from the photometric and extinction uncertainties. The red-dashed curves represent the best fit blackbody models. The days of the observations are indicated in blue. The temperatures quoted are derived from the best fit blackbody model.}
\label{Fig:SED}
\end{center}
\end{figure}

\subsection{Timing analysis of the optical light curves}
The \textit{TESS} light curve of V1047~Cen, presented in Figure~\ref{Fig:TESS_LC}, shows the early rise of the outburst between days 17 and 44. We searched for periodicity in the \textit{TESS} data by computing the Lomb-Scargle periodogram \citep{Lomb_1976,Scargle_1982,2018ApJS..236...16V} after detrending the light curve.
The periodogram shows modulation at a period of
$P_{\rm \it TESS} = 0.361 \pm 0.030$\,days (Figure~\ref{Fig:TESS_LC}) .
The period uncertainty is conservatively estimated as
\begin{equation}
\label{eq:perror}
P_{\rm error} = P^2/(2 T),
\end{equation}values

corresponding to a period change that would cause the easily noticeable phase shift of 0.5 between the first and the last
points in the light curve. The \textit{TESS} light curve covers the time range of $T = 21.5$\,days after detrending and removing the outlier measurements.
The peak-to-peak amplitude of the periodic variations is about 0.04\,mag in the \textit{TESS} photometric band (Figure~\ref{Fig:folded_LC}).

To confirm the significance of the detected periodicity, we compute
the false alarm probability of $10^{-4}$ using eq.~(14) of \cite{Scargle_1982} and
estimating the number of independent frequencies
(208 for the trial period range 0.1--3\,days)
following the prescription in sec.~5.3 of \citet{2003ASPC..292..383S}.
We confirm the above false alarm probability estimate with $10^4$ iterations
of bootstrapping \citep[e.g., sec.~7.4.2.3. of][]{2018ApJS..236...16V}.
The periodic signal persists whether we use polynomial or piecewise linear detrending techniques.
The signal can also be recovered from different subsections of the \textit{TESS} light curve and is not associated with any known instrumental period signal.

\begin{figure*}
\begin{center}

  \includegraphics[width=\columnwidth]{TESS_LC.pdf}~~~~~~~~
  \includegraphics[width=\columnwidth]{test_ls_periodogram-eps-converted-to.pdf}
\caption{\textit{Left:} the \textit{TESS} light curve of the 2019 outburst of V1047~Cen. \textit{Right:} Lomb-Scargle Periodogram of the \textit{TESS} data showing modulation at a period of 0.361 $\pm$ 0.03\,days.}
\label{Fig:TESS_LC}
\end{center}
\end{figure*}

We also performed Lomb-Scargle periodogram analysis of the photometric data
collected by the AAVSO observers in $B$, $V$, $R$, $I$ and $CV$ (clear filter observations with $V$ zero-point) bands.
In the section of the light curve between $t_0+76$\,days (when the multicolor
photometric observations start) and the seasonal gap starting on
$t_0+181$\,days, we found a period of $P_{\rm AAVSO} = 8.36 \pm 0.33$\,days.
The period is derived from a light curve combining measurements in all filters, its uncertainty is estimated using Eq.~(\ref{eq:perror}).
The Lomb-Scargle false alarm probability is $>10^{-5}$ for each photometric band
and much lower for the combined light curve.
The amplitude of the periodic variations is $\sim 1.2$\,mag in $V$ band (Figure~\ref{Fig:folded_LC}).
With the low false alarm probability, detection in different photometric bands as well as the subsets of data
collected by individual observers (Josch Hambsch, Godron Meyer)
and 12 cycles covered before the seasonal gap, the detection of $P_{\rm AAVSO}$
can be regarded as highly significant.

Surprisingly, $P_{\rm AAVSO}$ completely disappears from the AAVSO light curve after the seasonal gap.

$P_{\rm AAVSO}$ is not found in the \textit{TESS} data, however detrending removes variations on the relevant timescales. The duration of the \textit{TESS} light curve is only $3.2 \times P_{\rm AAVSO}$,
so periodic variations on this timescale cannot be properly probed with these data.

While the \textit{TESS} data set spans only 27.5 days, the AAVSO
broadband data span several hundred days and some nights include time-resolved photometry.
Therefore, one should expect to find the period
$P_{\rm \it TESS}$ in the AAVSO data; however it was not found.
The variations with period $P_{\rm \it TESS}$ could have been missed due
to a combination of their low amplitude and difficulties in detrending (large night-to-night
variations while the duration of each night's run is $<P_{\rm \it TESS}$).

Due to the large difference between $P_{\rm \it TESS}$ and $P_{\rm AAVSO}$,
it is very unlikely that one of them is the orbital period while the other is the superhump period (see e.g., \citealt{Patterson_etal_2005,Pearson_etal_2006,Smith_etal_2007,Smak_2020} and references therein).
While based on the \citet{Stolz_schoembs_1984} relation, the orbital period-superhump period excess and the orbital period are correlated, for an orbital period of 0.361\,days, the superhump period should be around 15 hours \citep{Otulakowska-Hypka_etal_2016}. Based on the current data, it is not possible to confirm whether one of these periods is the orbital or superhump period. However, there is other evidence that suggests the longer 8.36 d period is indeed the orbital period of the binary, implying an evolved secondary (see Section 4 for additional discussion).

It is possible that the two transient periods may be associated with different structures within the binary system. The structure associated with the eruption and responsible for $P_{\rm AAVSO}$ (perhaps the expanded atmosphere of the accreting white dwarf) may outshine the structure responsible for $P_{\rm \it TESS}$ (perhaps the accretion disk) that was dominating the light of the system in quiescence and at the initial stages of the brightness rise. This could lead to the TESS variations becoming swamped and hence undetectable, even if present in the underlying system.

After the seasonal gap, the AAVSO lightcurve reveals irregular variations on a timescale of days with amplitudes reaching 1\,mag. The switch from
nearly-periodic (with $P_{\rm AAVSO}$) to irregular variations could reflect a change in physical conditions of the emitting region, such as pulsation from the expanded atmosphere of the white dwarf (e.g., \citealt{2002ASPC..259..580S}). Another possibility for the disappearance of the $P_{\rm AAVSO}$ after the seasonal gap is that the inner part of the binary system
is hidden by an optically thick outflow that started around day 230 after outburst (see Section~\ref{spec_features}).

\begin{figure}
\begin{center}
  \includegraphics[width=\columnwidth]{aavso_V_everyone_before_gap_ls_periodogram-eps-converted-to.pdf}
\caption{Lomb-Scargle Periodogram of the $V$-band data showing modulation at a period of 8.36 $\pm 0.33$\,days.}
\label{Fig:LSP_V}
\end{center}
\end{figure}

\subsection{Radio light curve and spectral indices}

The MeerKAT radio light curve is plotted in Figure~\ref{Fig:optical_UV}. The $L$-band (900 - 1670 MHz) flux shows variability between $\simeq$ 0.7 and 1.0\,mJy during the optical outburst. There is no obvious correlation between the optical and radio emission.
At day 700, ten months after the end of the optical outburst, the radio emission from the system is still bright at $\simeq$ 0.91\,mJy (Figure~\ref{Fig:spec_fitting}).

\begin{figure*}
\begin{center}
  \includegraphics[width= 0.48\textwidth]{radio_LC_bars.pdf}
  \includegraphics[width= 0.48\textwidth]{spectral_fitting_1.pdf}
  \includegraphics[width= 0.48\textwidth]{spectral_fitting_2.pdf}
  \includegraphics[width= 0.48\textwidth]{spectral_fitting_3.pdf}
\caption{The MeerKAT radio light curve measured at 1.28 GHz is plotted in the \textit{top left panel}, for comparison with three radio spectra measured on days 256, 294, and 700. Spectral indices are listed in each panel. The vertical dashed lines on the light curve plot signify the epochs when spectra are measured.}
\label{Fig:spec_fitting}
\end{center}
\end{figure*}

We derived the spectral indices of the MeerKAT observations
on days 258, 301, and 421 (Figure~\ref{Fig:spec_fitting}). To do that, we divided the data into eight frequency intervals centered at 859 MHz, 1.016 GHz, 1.123 GHz, 1.230 GHz, 1.337 GHz, 1.444 GHz, 1.551 GHz, and 1.658 GHz,
allowing us to measure the flux density of the source at each sub-band. Consequently we calculated the spectral index, $\alpha$, by fitting a
single power law to the radio spectra that is represented as:
 $ S \propto \nu^{\alpha}$.
We note that the sub-band calibration for MeerKAT is still an active area of exploration, and hence the systematic uncertainties on the sub-band flux densities may yet be underestimated. While the spectrum is mostly flat, given the large uncertainties
in the spectral indices, caution is required when interpreting these indices. Overall, it is hard to draw strong conclusions from these values, other than that the radio emission is not optically thick. We elaborate more on the origin of the radio emission in Section~\ref{Disc}.

\begin{figure*}
\begin{center}
  \includegraphics[width=0.9\textwidth]{spec_main_text_2.pdf}
\caption{The optical spectral evolution of V1047~Cen throughout different stages of its outburst. The numbers in brackets are days since $t_0$. We also quote the resolving power of the spectra on the plot. Note that the calibration pipeline introduced artifacts in the continuum of the spectrum of day 74, particularly below 5000\,$\mathrm{\AA}$.}
\label{Fig:main_spec}
\end{center}
\end{figure*}

\subsection{Optical spectral evolution}
In Figure~\ref{Fig:main_spec} we present a sample of the optical spectra, evolving throughout the outburst of V1047 Cen and covering three stages: the rise (day 74) , the plateau phase (days 113 and 303), and the post-outburst (day 407). The full spectral evolution is presented in Figures~\ref{Fig:spec_1} to~\ref{Fig:spec_5}. The first spectrum, obtained 74 days after the start of the outburst, shows relatively narrow emission lines of H (Balmer) and \eal{He}{I}, with Full Widths at Half Maximum (FWHMs) $<$ 300\,km\,s$^{-1}$ and Full Widths at Zero intensity (FWZIs) of around 500\,km\,s$^{-1}$, which are typical features of a DN outburst (e.g., \citealt{Morales-Rueda_Marsh_2002}). These features co-exist and are superimposed on top of broader emission lines of H, \feal{O}{III}, and \eal{C}{VI}, with FWZIs of around 2500\,km\,s$^{-1}$ for the Balmer lines and 1100\,km\,s$^{-1}$ for the \feal{O}{III} lines, which are characteristic of a classical nova nebula (likely arising in the remnant of the 2005 nova eruption). We also identify an even narrower feature in the Balmer lines, during the first 2 spectral epochs (days 74 and 81), characterized by FWHMs of $<$ 40\,km\,s$^{-1}$ (see Figures~\ref{Fig:Halpha} and~\ref{Fig:beta}). While we do not have a definite explanation for the origin of these narrow features, it might be originating in
a dense circumbinary medium related to the presence of an evolved secondary (see Section~\ref{sec:secondary}).

By day 113, the fluxes of the Balmer lines had increased significantly relative to the \feal{O}{III} lines
$F_{H\alpha}/F_{5007} =1.56$ on day 74, compared with $\sim$15 on day 113 and
(e.g., $F_{H\beta}/F_{5007}$ increased from $\sim$0.35 on day 74 to $\sim$3.1 on day 113; see Figure \ref{Fig:beta_OIII} for a direct comparison between the evolution of H$\beta$ and the \feal{O}{III}lines). The high-resolution spectra show that the Balmer and \feal{O}{III} line profiles are very different from one another, with the \feal{O}{III} lines (in the high-resolution spectra) having rectangular shapes and jagged tops, characteristic of nova nebular lines. In contrast, H$\beta$ shows complex profiles which vary throughout the outburst. The Balmer lines also had significantly broadened, with the FWZIs increasing by factors of $\sim$2 compared to day 74, reaching 4000\,--\,4500\,km\,s$^{-1}$ (Figures \ref{Fig:Halpha} and~\ref{Fig:beta_OIII}). Note that a broad base in H$\alpha$ can be observed as early as day 81 but it becomes prominent from day 113 onwards.

\begin{figure*}
\begin{center}
  \includegraphics[width=0.7\textwidth]{Halpha.pdf}
\caption{The line profile evolution of H$\alpha$ throughout the outburst of V1047~Cen. The numbers between brackets are days after outburst. The red dashed lines represents the rest velocity ($V_{\mathrm{vrad}}$ = 0\,km\,s$^{-1}$). A heliocentric correction is applied to the radial velocities.}
\label{Fig:Halpha}
\end{center}
\end{figure*}

\begin{figure*}
\begin{center}
  \includegraphics[width=0.77\textwidth]{Hbeta_OIII.pdf}
\caption{The evolution of the profiles of the \feal{O}{III} lines at 4959 and 5007\,$\mathrm{\AA}$ in comparison to H$\beta$, throughout the outburst of V1047~Cen. The numbers between brackets are days after outburst. The red dashed lines represents the rest velocity ($V_{\mathrm{vrad}}$ = 0\,km\,s$^{-1}$) of \feal{O}{III} 4959\,$\mathrm{\AA}$. A heliocentric correction is applied to the radial velocities.}
\label{Fig:beta_OIII}
\end{center}
\end{figure*}

From day 130, some of the \eal{He}{I} emission lines show P Cygni profiles with absorption troughs at blueshifted velocities of around 1400\,km\,s$^{-1}$ (Figures~\ref{Fig:spec_2} and~\ref{Fig:spec_3}).

Between days 262 and 303 \eal{O}{I} and \eal{N}{II} emission lines emerge, while the \feal{O}{III} nebular lines fade relative to the other spectral lines. At this stage, in addition to the \eal{He}{I} emission lines, Balmer, \eal{N}{II} and \eal{O}{I} emission lines also show P Cygni profiles with absorption troughs at velocities of around $-1200$ to $-1400$\,km\,s$^{-1}$ (Figure~\ref{Fig:PCygni}). During the same period, particularly on days 281 and 286, the optical spectra which extend above 8000\,$\mathrm{\AA}$ show broad double-peaked emission of \eal{O}{I} at 8446\,$\mathrm{\AA}$ with
FWZIs of around 3600\,km\,s$^{-1}$ (Figure~\ref{Fig:spec_6}). The \eal{O}{I} P Cygni profile at 7773\,$\mathrm{\AA}$ and the double-peaked \eal{O}{I} 8446\,$\mathrm{\AA}$ are not typical features of DN outbursts \citep{Morales-Rueda_Marsh_2002}.

\begin{figure}
\begin{center}
  \includegraphics[width=\columnwidth]{PCygni.pdf}
  \caption{The P Cygni line profiles of H$\beta$, \eal{He}{I} 7065$\mathrm{\AA}$, and \eal{O}{I} 7773\,$\mathrm{\AA}$ taken on day 281 and \eal{O}{I} 1.128 $\mu$m taken on day 322 (Numbers between brackets are days since $t_0$). The red and blue dashed lines represent $v_r$ = 0\,km\,$^{-1}$ and $v_r$ = $-1300$\,km\,$^{-1}$, respectively.}
\label{Fig:PCygni}
\end{center}
\end{figure}

After the end of the optical outburst, the spectrum obtained on day 407 shows substantial changes (Figure~\ref{Fig:main_spec}), with the \feal{O}{III} emission lines now dominating the spectrum relative to the Balmer lines. In Figure~\ref{Fig:line_ratio} we show the evolution of the line ratio between the \feal{O}{III} line at 5007\,$\mathrm{\AA}$ and H$\beta$ and the evolution of the equivalent width of H$\beta$, \feal{O}{III} 5007\,$\mathrm{\AA}$, and \feal{O}{III} 4995\,$\mathrm{\AA}$ emission lines. It is remarkable how the \feal{O}{III} lines were still relatively strong during the rise of the 2019 outburst, but fade throughout the outburst, before strengthening again relative to the Balmer lines by the end of the outburst.

Half a year later, our spectra on days 615 and 643 still show strong \feal{O}{III} lines, in addition to other lines, which likely originate in the 2005 classical nova nebula, such as \feal{O}{II} 7320\,$\mathrm{\AA}$, \feal{N}{II} 5755\,$\mathrm{\AA}$, and high ionization \feal{Fe}{VII} lines (Figure~\ref{Fig:spec_5}).

Our last spectrum taken on day 774, more than 2 years after the start of the 2019 outburst, shows in addition to the 2005 nova nebular lines, weak lines from high ionization transitions of He, N, O, and C, such as the \eal{He}{II} lines at 4686 and 5412\,$\mathrm{\AA}$, \eal{N}{V} 4603\,$\mathrm{\AA}$, \eal{O}{V} 5920\,$\mathrm{\AA}$, \eal{C}{IV} 5802\,$\mathrm{\AA}$, and \eal{O}{IV} 7713\,$\mathrm{\AA}$ or \eal{Ne}{IV} 7716\,$\mathrm{\AA}$ (Figure~\ref{Fig:spec_7}). Such lines have been observed in systems like V617~Sgr and V~Sge and are associated with nuclear shell burning (e.g., \citealt{Herbig_etal_1965,Cieslinski_etal_1999,Steiner_etal_1999}).

\begin{figure}
\begin{center}
  \includegraphics[width=\columnwidth]{line_ratio.pdf}
\caption{\textit{Top}: the optical light curve in the $I$-band. \textit{Middle}: The evolution of the flux ratio between H$\beta$ and \feal{O}{III} 5007\,$\mathrm{\AA}$. \textit{Bottom}: The evolution of the of the equivalent width (EW) of the H$\beta$, \feal{O}{III} 5007\,$\mathrm{\AA}$, and \feal{O}{III} 4995\,$\mathrm{\AA}$ emission lines. }
\label{Fig:line_ratio}
\end{center}
\end{figure}

\begin{figure}
\begin{center}
  \includegraphics[width=0.9\columnwidth]{Halpha_jets.pdf}
\caption{High-velocity ``satellite" features in the late H$\alpha$ line profiles, taken on days 615 and 643. The origin of these features is considered in the Discussion.}
\label{Fig:jet_lines}
\end{center}
\end{figure}

\begin{figure*}
\begin{center}
  \includegraphics[width=1.01\textwidth]{Gemini_spec.pdf}
\caption{The Gemini-S 0.9\,--\,2.5\,$\mu$m spectra obtained on days 323 and 330. We identify hydrogen recombination and other He and O lines. For the recombination lines, we show the transition of the longest wavelength member.}
\label{Fig:Gemini_spec}
\end{center}
\end{figure*}

Several months after the end of the optical outburst on days 615 and 643, H$\alpha$ shows satellite emission components at $\pm$ 2500\,km\,s$^{-1}$ (Figures~\ref{Fig:jet_lines} and \ref{Fig:spec_5}). We elaborate on the origin of these components in the discussion.

\subsection{IR spectral evolution}
The Gemini and SOFIA IR spectra are represented in Figures~\ref{Fig:Gemini_spec} and \ref{Fig:SOFIA_spec}, respectively. A detailed description and analysis of the IR spectral evolution
during the first 120 days of the outburst is presented in \citet{Geballe_etal_2019}. The Gemini $JHK$ spectra obtained on day 322 show emission lines of \eal{H}{I} (Paschen, Brackett, and Pfund series), \eal{He}{I}, \eal{He}{II}, and \eal{O}{I}. The FWZIs of the lines are around 4000\,km\,s$^{-1}$, similar to the ones measured from the optical spectral lines. The \eal{O}{I} line at  1.1289$\mu$m and some of the \eal{H}{I} and \eal{He}{I} lines show P Cygni profiles with absorption troughs at blue-shifted velocities of 1100 to 1800\,km\,s$^{-1}$, also comparable to the ones measured for optical lines (Figure~\ref{Fig:PCygni}).

Figure~\ref{Fig:SOFIA_spec} presents the SOFIA composite
spectra obtained on day 87, including contemporaneous, dereddened
(see \S~\ref{red_dist}) $BVRI$
photometry obtained from the AAVSO database, as well as NEOWISE
photometery obtained 18 days later. A blackbody fit to the 2019
SOFIA spectra yields T$_{\rm bb} = 425 \pm 12$~K, which we
interpret as thermal emission from circumstellar material heated
by the outburst event.  WISE photometery obtained
in 2010 (prior to the eruption) is also plotted and a blackbody
fit to this photometry yields a cooler temperature of
T$_{\rm bb} = 315 \pm 30$~K. Thus the circumstellar
material (dust, likely from the 2005 nova event) has been heated as a result of processes related to the 2019 outburst. The SOFIA spectra show no evidence for strong H or \eal{He}{I} emission lines on day 87. However the [\ion{O}{4}] 25.91~\micron{} fine structure line, frequently seen in
other novae \citep{Gehrz_etal_2015,Evans_etal_2012,Helton_etal_2012}
is marginally detected. A Gaussian fit gives a line flux of
$6.9 \pm 1.6~\rm{W}~\rm{m}^{-2}$. The upper level of this line
is collisionally de-excited at electron densities ($n_{e}$) in excess of
$9.9 \times 10^{3}$~cm$^{-3}$ (for an electron temperatures of  $10^{4}$~K);
the presence of the line therefore indicates that the electron density
in the region where the line is produced is less than this value.

\begin{figure}
\begin{center}
\includegraphics[width=0.51\textwidth, height = 7.4cm]{V1047Cen_210813_LogLin_erg_s_cm2_Trim_NO_Label2BBFit+SOFIA+WISE+FilterPhot.pdf}
\caption{The 2019 SOFIA mid-infrared spectra of V1047 Cen obtained on
day 87. The blue symbols are the SOFIA spectra (combination of two
grating settings G111 and G227) which includes all spectral points
where the atmospheric transmission at the flight altitudes are great
than 70\% (hence the gap near $\sim 9.5$~\micron{}). The wavelength
coverage of the two gratings is not continuous. Superposed is the
best-fit composite blackbody (red solid line), i.e., the sum
of a T$_{\rm bb} = 7070\pm950$~K from the 2019 photometry
data  plus a T$_{\rm bb} = 425\pm12$~K, the latter is interpreted as
thermal emission from heated circumstellar material (dust) from
the blast of the 2019 eruption. The inverted orange
triangles are contemporaneous (2019) dereddned $BVRI$ and NEOWISE
photometry, wile the filled-black circles are 2010 WISE
photometry showing that the infrared spectral energy distribution
peak was represented by a cooler blackbody (solid gray line)
of T$_{\rm bb} = 315\pm30$~K.}
\label{Fig:SOFIA_spec}
\end{center}
\end{figure}

\subsection{The 2005 eruption vs the 2019 outburst}

The 2005 nova eruption, reached magnitudes brighter than 8\,mag in the visual, compared to a peak $V$-band magnitude of 13.5 for the 2019 outburst (Figure~\ref{Fig:LC_comp} and~\ref{Fig:optical_UV}). These translate to an absolute visual magnitude of $\approx -8$
for the 2005 eruption and $\approx -2$
for the 2019 outburst, using $A_V = 3.03$ and a distance of 3.2\,kpc (see Section~\ref{red_dist}). This clearly implies that the 2005 event was much more energetic than the 2019 one. The morphology of the optical light curves are also distinctively different with the 2005 light curve typical of a fast nova, while the light curve of the 2019 event plateaued for around a year. The velocities measured from the spectra taken during the 2005 eruption range from 750 to 1800\,km\,s$^{-1}$. These velocities are slow to moderate in comparison to the velocities observed in
classical novae. During the rise of the 2019 outburst, the velocities measured from some of the lines were around a few hundreds km\,s$^{-1}$ (typical of CVs), but later in the outburst, the spectral lines showed velocities of $\lesssim$ 2000 km\,s$^{-1}$, raising questions about the nature of this event.

\section{Discussion}
\label{Disc}

\subsection{Reddening and distance}
\label{red_dist}
In order to derive the reddening towards the system, we use the equivalent widths of several absorption lines from diffuse interstellar bands, in combination with the empirical relations from \citet{Friedman_etal_2011}. We derive an average $E(B-V) \approx 1.0 \pm 0.1$ and $A_V = 3.0 \pm 0.1$ for $R_V = 3.1$, in good agreement with the results of \citet{Geballe_etal_2019}.
We avoid using the interstellar \eal{Na}{I}~D absorption doublet due to saturation.

The Gaia parallax measurement of V1047~Cen suffers from large uncertainties. The latest Gaia EDR3 parallax of the system is 0.338$\pm$0.249\,mas \citep{Gaia_EDR3}. With a flat prior (Galactic latitude-based priors are not appropriate for novae), this implies a distance of $2.7^{+3.9}_{-1.0}$ kpc, which is not a useful constraint.

Due to the high uncertainty in the Gaia parallax, we use the three dimensional Galactic reddening maps from \citet{Chen_etal_2019} and our measured reddening value to estimate the distance to V1047 Cen. The reddening map uses measurements from the Gaia DR2, 2MASS and WISE surveys. Therefore, we converted the previously derived $E(B-V)$ measurement to reddening values in the 2MASS $JHK$ filters and the Gaia DR2 $G$, $G_{Bp}$, and $G_{Rp}$ bands, using the extinction law from \citet{Wang_etal_2019} and \citet{Chen_etal_2019}.
We derive $E(G-K_s) = 2.2 \pm 0.1$, $E(G_{Bp}-G_{Rp}) = 1.3 \pm 0.1$, and $E(H-K_s) = 0.2 \pm 0.1$. Using these reddening values, we derive an average distance of $3.2 \pm 0.2$\,kpc, consistent with the Gaia parallax distance within its large uncertainties.
The above uncertainty for the reddening-based distance likely underestimates the systematic uncertainties in this calculation.

\subsection{Evolutionary stage of the secondary}
\label{sec:secondary}

If the 8.36\,day period found in the AAVSO data is the orbital period of the system, given that the donor is filling its Roche lobe, it implies that the star is evolved, either a subgiant or a lower luminosity red giant \citep{Webbink_etal_1983}. Two similar examples are (1) V392~Per (Nova Per 2018), which has an orbital period of 3.51 days, and has been demonstrated by \citet{Munari_etal_2020} to host a subgiant secondary; and (2) GK Per, which has an orbital period of 1.9 days and hosts a subgiant K secondary star \citep{Bianchini_etal_1981}. The 8.36 days period would be the longest for a CV that showed a nova eruption, but is on the shorter end for systems with luminous red giant secondaries, which are typically characterized by periods of the order of hundreds of days \citep{Whitelock_1987,Mikolajewska_2003}. This would make V1047 Cen an intermediate system between typical CVs and systems with luminous giant secondaries (symbiotic systems; \citealt{Kenyon_1986,Mikolajewska_etal_1988,Munari_2019}).

The average $V$- and $I$-band magnitudes as measured by OGLE long after the nova, between 2013 and 2018, are $I$ = 17.12 and $V$ = 17.5.
The post-nova color index is then $(V-I)$ = 0.38$\pm$0.07, implying $(V-I)_0$ = -1.19. In Figure~\ref{Fig:charts} we show the OGLE light curve during the six years prior to the 2019 outburst. Adopting a distance of 3.2\,kpc and $A_I$ = 1.46, we derive $M_I \approx 3.5$ during the post-nova period. While this value is consistent with the absolute magnitude of a subgiant star, we also expect contributions from the accretion disk, hot spot, and the 2005 nova nebula. Therefore, using the optical magnitudes/colors during quiescence to constrain the evolutionary stage of the companion star is not straightforward.
We note that the system is not listed in the USNO-B1.0 catalog \citep{Monet_etal_2003}, and is not apparent in visual inspection of blue or red plates of the Digitized Sky Survey (DSS; Figure~\ref{Fig:charts}). Since the completeness level of this survey is roughly $V \sim 20.5$--21 mag, this suggests the quiescent magnitude of the system is fainter than in the post-nova OGLE imaging and rules out anything more luminous than a marginally evolved secondary. In Figure~\ref{Fig:charts}, we show charts of the field of V1047~Cen from DSS (red plates taken in Feb. 1999), OGLE (taken in Feb 2014), and SOAR (taken in June 2019). Clearly, there is an excess in brightness of the system in 2014, nine years after the 2005 nova eruption, compared to pre-nova. This indicates that the system did not return to the pre-nova brightness, which could be due to ongoing nuclear burning on the surface of the white dwarf.

The OGLE $(V-I)$ colors between 2010 and 2019 indicate a blue source with extremely high temperatures, in excess of $10^5$\,K. These nonphysically blue colors (for an evolved or main sequence star) could be due to the contribution of emission lines to the spectra. Particularly, the forbidden \eal{O}{III} lines from the nova nebula, whose fluxes contribute to the $V$-band. These lines were relatively strong compared to the other lines during the early spectra of the 2019 outburst and after the end of the outburst (Figure~\ref{Fig:main_spec}). Unlike the $V$-band, no strong nebular lines contribute to the flux in the $I$-band (Figure~\ref{Fig:spec_1}). In addition, the blue colors could again indicate that ongoing nuclear burning on the surface of the WD. Some nova systems have shown continued thermonuclear burning on the surface of the WD several years after nova eruptions (e.g., \citealt{Zemko_etal_2015,Zemko_etal_2016}), but these authors' conclusions rely mostly on delayed X-ray emission. In systems where the nuclear burning lasts for several years after nova eruptions, the masses of the WDs are expected to be low ($M_{\mathrm{WD}} <1$\,M$_{\odot}$; \citealt{Wolf_etal_2013,Henze_etal_2014_Mar}).

Note that the \textit{Swift} non-detection of V1047 Cen on April 2008, less than 3 years after the 2005 nova, argues against residual nuclear burning on the white dwarf. If residual nuclear burning is ongoing on the white dwarf we expect to detect supersoft X-ray emission (e.g., \citealt{Page_etal_2020}) from the system, unless the nova ejecta were still optically thick to the supersoft emission from the white dwarf 3 years after the nova. Alternatively, nuclear burning could have resumed on the white dwarf at a low rate, a few years after the nova, due to enhanced mass-transfer rate from an irradiated, evolved companion.

\begin{figure*}
\begin{center}
  \includegraphics[width= 8.3cm, height = 7.3cm]{pre_outburst.pdf}
  \includegraphics[width= 8.3cm, height = 7.3cm]{DSS.pdf}
  \includegraphics[width= 8.3cm, height = 7.3cm]{OGLE.pdf}
  \includegraphics[width= 8.3cm, height = 7.3cm]{SOAR.pdf}
\caption{\textit{Top-left:} the OGLE light curve between 2013 and 2019. Days are relative to the 2019 outburst start. \textit{Top-right:} A red plate of the field of V1047~Cen from the Digitized Sky Survey (DSS) obtained in February 1999 (position of V1047~Cen indicated by a green circle). \textit{Bottom-left:} combined OGLE images of V1047~Cen (circled in green) taken between in February 2014. \textit{Bottom-right:} SOAR acquisition image of V1047~Cen (circled in green) taken during the in June 2019, during the 2019 outburst. In all the charts the field is 2 $\times$ 2 arcmin. V1047 Cen lies on the edge of the OGLE images, which is the reason why the chart is cropped on the west edge.}
\label{Fig:charts}
\end{center}
\end{figure*}

\subsection{The 2019 outburst of V1047~Cen - a record breaker}

Only a few classical novae are known to have shown DN outbursts after a nova eruption, such as GK Per (Nova Persei 1901; \citealt{Sabbadin_Bianchini_1983,Bianchini_etal_1986}), V1017~Sgr (Nova Sagittarii 1919; \citealt{Sekiguchi_1992,Salazar_etal_2017}) and V446 Her (Nova Herculis 1960; \citealt{Honeycutt_etal_1995,Honeycutt_etal_2011}). The first DN outbursts of these novae were recorded 47 years (GK Per), 54 years (V1017~Sgr), and 30 years (V446~Her) after the nova eruption. The 14 year gap between the classical nova eruption and the 2019 outburst of V1047~Cen is the shortest ever recorded, if it is indeed a DN outburst.
Note that the older novae such as GK~Per and V1017~Sgr might have had earlier DN outbursts that were missed due to monitoring gaps.

The rise to peak of V1047~Cen, its plateau phase, and its decline lasted 100, 210, and 100 days, respectively. This 400 day outburst is twice as long as the longest known DN outburst, previously recorded for V1017~Sgr and four times longer than the longest outburst recorded for GK Per (e.g., the 2006 outburst; \citealt{Evans_etal_2009}).

The peak absolute magnitude at $V$-band during the 2019 outburst reached $\approx -2$\,mag. This is higher than the typical absolute magnitude of dwarf novae during outburst, which ranges between 3.8 and 2.6 \citep{Ramsay_etal_2017}. Some systems with long periods (of the order of days) such as V1017~Sgr and V630~Cas, have more than usual luminous DN outbursts, with absolute magnitudes reaching $-0.3$ and 1.4, respectively. However, even these are still substantially less luminous compared to V1047~Cen. Therefore, if 2019 outburst of V1047~Cen is a DN, it would be the most luminous outburst of a DN observed to date. Note that \citet{Kawash_etal_2021a} showed that some DN outbursts could reach absolute magnitudes brighter than zero but accurate distances and extinction values were unavailable for these systems in Kawash et al.'s study.

DNe are known to show a diversity of light curve morphologies, even in one system. These outbursts could be narrow  (a rise and decline with only a short period of time spent at peak/plateau), or wide and long (characterized by an extended plateau phase; \citealt{Hellier_2001}).

The narrow but slow outbursts of V1017~Sgr show a short plateau of less than three weeks \citep{Salazar_etal_2017}, in contrast to the very long (more than 200 days) plateau duration of V1047~Cen. However, the morphology of the light curve of V1047~Cen is more similar to some of the outbursts of GK Per, particularly the 2006 one, which lasted for around 100 days and showed substantial variability during a plateau interval of more than 2 months \citep{Evans_etal_2009}. The plateau phase of a DN occurs when the entire accretion disk enters outburst and sustains it, with matter continuously falling onto the white dwarf. In some cases, enhanced mass transfer due to irradiation of the companion star helps sustain the plateau for an extended period of time \citep{Hellier_2001}.

The decline rate of DN outbursts is known to correlate with the orbital period of the system. This relation is known as the Bailey relation \citep{Bailey_1975}, and is particularly valid for systems above the period gap \citep{Otulakowska-Hypka_etal_2016}. The rise rate is associated with the direction of the propagation of the heating or instability wave, whether it is propagating outside-in (rapid rise) or inside-out (slow rise; see \citealt{Otulakowska-Hypka_etal_2016} and references therein). But the rise time should also be correlated with the size of the disk and therefore, the orbital period of the system \citep{Warner_2003}.
If the 2019 V1047~Cen is a DN outburst, the slow rise rate indicates an outburst propagating inside-out in a system with an orbital period of the order of a few days. We used a rise rate of $\approx$ 25\,d\,mag$^{-1}$ measured from the $OGLE$ light curve, and the empirical relations of both \citet{Warner_2003} and \citet{Otulakowska-Hypka_etal_2016}, to estimate an orbital period of $\approx$ 3.8 days and 29.1 days, respectively. While the relations suffer from large uncertainties, they both indicate that the system should have an orbital period of the order of days up to a few weeks.
We also used similar relations from both of the above papers
which relate the rate of decline to the orbital period. Using these relations and a decline rate of 28.75\,d\,mag$^{-1}$, we derive periods of $\approx$ 4.9 days and 2.9 days, from the relations of \citet{Warner_2003} and \citet{Otulakowska-Hypka_etal_2016}, respectively.

Both \citet{Smak_2000} and \citet{Otulakowska-Hypka_etal_2016} also found a correlation between the orbital period and duration of the outburst. Based on this relation and an outburst duration of 400 days, the orbital period of the system is also expected to be of the order of days. Note that the duration of the outburst of V1047~Cen is an outlier to the sample used in these studies. While all these relations suffer from large uncertainties, they all indicate that the system is very likely to have an orbital period of the order of days, which is consistent with the 8.36 days period found in the AAVSO data.

\subsection{Implications for the hibernation scenario}
The hibernation scenario, which was first suggested to explain the apparent lack of nova eruptions in CV systems showing DNe, suggests that all CVs undergo cycles of low and high states of mass transfer rate \citep{Shara_etal_1986}. After a classical nova eruption, the hibernation scenario predicts that the CV undergoes a phase of high-mass transfer rate due to irradiation of the secondary by the hot white dwarf; this can last for several centuries \citep{Hillman_etal_2020}. However, due to mass loss in the nova eruption, the binary now has a wider separation, and as irradiation shuts off and the secondary relaxes into its new Roche lobe, the mass transfer rate will drop and perhaps even turn off entirely. During this hibernation phase, standard angular momentum loss mechanisms like magnetic braking will act, gradually bringing the binary into contact again. Mass transfer will then resume onto the white dwarf surface, eventually accumulating enough accreted material  to trigger the next classical nova eruption.

During the state of high mass transfer after a classical nova eruption (\.M$ \approx 10^{-8}$\,M$_{\odot}$\,yr$^{-1}$),
the system is not expected to undergo DN outbursts, which require lower mass transfer rates of the order of $\sim$ 10$^{-10}$\,M$_{\odot}$\,yr$^{-1}$ \citep{Osaki_1974,Cannizzo_etal_1988,Warner_1995}. Therefore, the 14 year gap between the 2005 nova eruption and 2019 outburst of V1047~Cen challenges the hibernation scenario on one hand and raises question about the nature of this outburst on the other hand.

However, for systems with orbital periods of a few days, DN outbursts could still take place a few decades after a nova eruption, during a phase of high mass transfer. This is because, in systems characterized by long orbital periods, the accretion disk is large enough that the critical mass accretion rate, below which DN outbursts are possible, is also large. This critical mass accretion rate scales with the radius of the accretion disk as $R^{2.6}$  \citep{Kim_etal_1992,Cannizzo_1993,Warner_1995,Geballe_etal_2019}.
Therefore, in the cases of GK~Per and V1017~Sgr, with orbital periods around 2 and 5.7 days respectively, the fact that they show DN outbursts soon after nova eruptions is not surprising \citep{Kim_etal_1992}.
If the period of 8.36 days found in the broadband AAVSO photometry is the orbital period of V1047 Cen, its accretion disk---and the critical mass transfer rate below which outbursts are possible---is  even larger than in GK Per and V1017~Sgr.

If the period of 0.361 days derived from the \textit{TESS} data is the orbital period of the system, then the size of V1047~Cen's accretion disk is still within the typical range for CVs. In this case it is more challenging to explain the 14 year gap between the nova eruption and the 2019 outburst. However, V446~Her, which was observed to show regular DN outbursts three decades after its classical nova eruption in 1960, is characterized by a period of 0.207 days \citep{Honeycutt_etal_1995,Honeycutt_etal_2011}. This could mean that under certain conditions, a DN outburst is still possible during post-nova high mass transfer rate conditions, even if the system is harboring an average size accretion disk. It is therefore essential to accurately constrain the period of the system spectroscopically, which will be the aim of future work.

\subsection{Origin of the spectral features}
\label{spec_features}
During the rise phase (days 0 to 90), the spectrum of V1047 Cen resembled that of a DN superimposed on
a spectrum of an old classical nova shell. However, the subsequent evolution, particularly of the Balmer lines, differed considerably from that of DNe. The widths of the Balmer lines reached velocities much larger than those of the nova's nebular lines, with FWZIs of around 4000\,km\,s$^{-1}$ (compared to FWZIs $\approx$ 1100\,km\,s$^{-1}$ for the latter). These velocities are also much larger than those measured for the Balmer lines during GK Per's recent outburst (FWHM $\sim$ 500-650\,km\,s$^{-1}$; e.g., \citealt{ATel_7217}).
The Balmer lines also showed substantial increases in their line fluxes relative to the nova nebular lines more than 100 days into the 2019 outburst (Figure~\ref{Fig:beta_OIII}). The cause of this dramatic flux increase and broadening is not clear, but it could be due to a fast, low density outflow. All this indicates that, in addition to the potential DN outburst, there are other mechanisms shaping the electromagnetic signatures of V1047~Cen.

Between days 262 and 310, most of the lines in the spectrum (e.g., Balmer, \eal{He}{I}, and \eal{O}{I}) showed P Cygni-like absorption features at velocities of around $-1400$\,km\,s$^{-1}$. Such features are not characteristic of a DN outburst. They probably originate in an optically thick wind/outflow. The onset of these features coincides with the start of the brightness increase between days $\approx$ 260 and 310 (Figure~\ref{Fig:optical_UV}), supporting the possibility of a wind or shell ejection during this phase of the outburst.

The satellite emission components at $\pm$ 2500\,km\,s$^{-1}$ in H$\alpha$ around 7 months after the end of the outburst, are reminiscent of the features observed in Z And-like classical symbiotic outbursts.
They are associated with collimated bipolar flows (see e.g., \citealt{Burmeister_etal_2007,Skopal_etal_2013,Tomov_Tomova_2013}).

\subsection{The origin of the bright radio emission}
\label{radio_disc}
In the past decade, CVs have been established as radio emitters during both quiescent and outburst phases (e.g., \citealt{Cordova_etal_1983,Coppejans_etal_2015,Coppejans_etal_2016,Barrett_etal_2017}). Their radio properties are diverse, and the emission mechanisms powering CV radio luminosity are still being established.

\citet{Coppejans_etal_2015} observed several nova-like CVs and found that they are significant radio emitters. Nova-like CVs are non-magnetic CVs characterized by a
sufficiently high mass transfer rate
to maintain the accretion disk in a constant hot state, unlike CV systems which undergo DN outbursts. For the individual systems in their sample, \citet{Coppejans_etal_2015} found that the emission in these nova-like CVs is consistent with optically thick synchrotron, gyro-synchrotron, or cyclotron maser emission. Their 6 GHz spectral radio luminosities are $\sim10^{15}-10^{17}$ erg\,s$^{-1}$\,Hz$^{-1}$ \citep{Coppejans_etal_2015}.

Non-magnetic CVs which undergo DN outbursts are also known to be radio sources during outburst \citep{Coppejans_etal_2016}. The most famous of these systems is SS~Cyg, which shows a radio flare during the early days of its optical outbursts, peaking at $\sim$ 1\,mJy before fading gradually, dropping below radio detectability by the end of the optical outburst. This flaring radio emission resembles that of X-ray binaries (XRBs) and is suggested to be the result of synchrotron emission from a transient jet (e.g., \citealt{Kording_etal_2008,Miller-Jones_etal_2011,Russell_etal_2016,Coppejans_etal_2016,Coppejans_etal_2020}). \citet{Coppejans_etal_2016} found that the specific radio luminosities of DNe in outburst range between 10$^{14}$ and 10$^{16}$\,erg\,s$^{-1}$\,Hz$^{-1}$ at 10\,GHz.

Based on a radio survey targeting a large sample of magnetic CVs, \citet{Barrett_etal_2017} suggest that they are also radio emitters, dominated by weakly polarized gyrosynchrotron emission or highly polarized electron-cyclotron maser emission. Most of these are nearby (less than a kpc) sources and are characterized by flux densities, $\sim$20--400 $\mu$Jy; the implication is radio spectral luminosities in the range $\sim10^{14}-10^{17}$ erg\,s$^{-1}$\,Hz$^{-1}$ \citep{Barrett_etal_2020}. Only one of them, AE~Aqr, shows substantially higher flux density ($\sim$5\,mJy), but it is located remarkably nearby, at a distance of $\approx$ 90\,pc \citep{Ramsay_etal_2017}.

Similar to the other observational features of V1047~Cen across the spectrum, the radio emission from this system is record breaking and puzzling.  Assuming a distance of 3.2\,kpc, the spectral luminosity of V1047~Cen at 2 GHz is $\sim$ 10$^{19}$\,erg\,s$^{-1}$\,Hz$^{-1}$. This means that V1047~Cen is at least two orders of magnitude brighter at radio wavelengths than other cataclysmic variables.

If the radio emission from a DN outburst is optically thick thermal emission from ionized gas at a typical brightness temperature of 10$^4$\,--\,10$^5$\,K, the emitting regions should have radii of $\sim$ 10$^{14}$\,--\,10$^{15}$\,cm (assuming circular sources as projected on the sky) to produce the observed flux densities in V1047~Cen. During a DN outburst, it is reasonable to assume that the emitting gas has a size equal to the size of the binary. CVs with orbital periods of the order of a few hours have orbital radii of $\sim$ 10$^{11}$\,cm. Even for CVs with orbital periods of around 9 days, the orbital radius is $\sim$ 10$^{12}$\,cm, still at least two orders of magnitude smaller in size than the optically thick emitting gas responsible for the observed fluxes. This implies that if the outburst is a DN, any emission of order the size of the orbit must be non-thermal (brightness temperature $> 10^5$ K).

One possibility is that the radio emission is non-thermal synchrotron emission from a transient radio jet, similar to the case of other DN outbursts. While the flux density of V1047~Cen equals that of SS~Cyg during its short-lasting flare, the luminosity of V1047~Cen is several orders of magnitude brighter than that of SS Cyg (the distance to SS Cyg is $\approx$ 114\,pc).
In Figure~\ref{Fig:LxLr} we plot V1047~Cen on the $L_{\mathrm{radio}}$ vs $L_X$ diagram in comparison with jet emission in other CVs and other astrophysical events such as XRBs, accreting millisecond X-ray pulsars, and transitional millisecond pulsars (we multiply the spectral luminosity by 5 GHz to estimate the radio luminosity at 5 GHz assuming a flat spectrum).
V1047 Cen's 5\,GHz luminosity, $\sim$ 1 $\times 10^{29}$\,erg\,s$^{-1}$, is much higher than that observed in CVs and is comparable to (or even brighter in some cases than) XRBs and pulsars with compact primaries like neutron stars and black holes. Such an energetic jet should also be a bright X-ray source. We used WebPIMMS \footnote{\url{https://heasarc.gsfc.nasa.gov/cgi-bin/Tools/w3pimms/w3pimms.pl}} to derive the X-ray unabsorbed flux using the stacked \textit{Swift} X-ray detection and assuming a 5\,keV thermal bremsstrahlung model and $N(H)$ = 8.4 $\times$ 10$^{21}$\,cm$^{-2}$ (we use $N(H)$ = 2.81 $\times$ 10$^{21}$ $A_V$; \citealt{Bahramian_etal_2015}). This translates to an X-ray luminosity of around 3.7$\times$10$^{31}$\,erg\,s$^{-1}$ at a distance of 3.2\,kpc, which is comparable to the X-ray luminosity of other CVs in outburst.

\begin{figure*}
\begin{center}
  \includegraphics[width=0.8\textwidth]{lrlx_plot.pdf}
\caption{The plane of 5 GHz radio luminosity ($L_{\mathrm{radio}}$) vs. 1--10 keV X-ray luminosity ($L_X$), with points representing X-ray binaries (accreting neutron stars, black holes), accreting and transitional millisecond pulsars, and CVs. The radio luminosity of V1047 Cen at 5 GHz is plotted with dashed blue horizontal line and is calculated assuming a flat spectrum.
Adopted from \citet{Bahramian_etal_2018_1252036} \url{https://zenodo.org/record/1252036/export/hx}.}
\label{Fig:LxLr}
\end{center}
\end{figure*}

It is not possible to determine if V1047~Cen showed flaring radio emission during the early days
of the outburst like SS~Cyg and was an even brighter radio source at that time, or if the $\sim$ 1\,mJy emission remained more or less constant throughout the optical outburst. Nevertheless, what is really puzzling is that at the end of the optical outburst, the radio flux did not drop. This is in contrast to the trends shown by SS Cyg and other DN outbursts, where the radio flux drops below the detection limit by the end of the optical outburst \citep{Russell_etal_2016}. This disparity between the radio behavior and energetics of V1047~Cen
and other DN outbursts raises additional questions about the origin of the radio emission in V1047~Cen and if a transient synchrotron jet is an appropriate explanation.

\citet{Bode_etal_1987} conducted a radio survey of 26 classical novae observed with the VLA less than 10 years after eruption and detected only two in the sample. Both novae were detected at 5\,GHz, 2 years (V4077~Sgr) and 8 years (NQ~Vul) after the nova eruptions. The emission from these two novae was suggested to be thermal and had radio spectral luminosities $\sim10^{18}-10^{19}$ erg\,s$^{-1}$\,Hz$^{-1}$. The higher-end radio spectral luminosity of V4077~Sgr is due to the fact that the system was observed only 2 years after the nova eruption; radio emission from classical novae often peaks on timescales $\sim$1--3 yr after eruption \citep{2021arXiv210706251C}. In addition, given that radio emission from decade old or older nova ejecta originates in the nova's extended thermal remnant, it can vary only slowly, not on the weeks--months timescales observed in V1047~Cen \citep{2021arXiv210706251C}. Therefore, it is very unlikely that the radio emission from V1047~Cen originates in the ejecta of the 2005 classical nova eruption.

The long-lasting, luminous radio emission in V1047~Cen is remarkably similar to the radio emission from Z And-like classical symbiotic, accretion-powered, outbursts (e.g., \citealt{Crocker_2001,Mikolajewska_2002,Brocksopp_2004,Sokoloski_etal_2006}). These systems are luminous radio sources ($L \approx 10^{18}-10^{19}$ erg\,s$^{-1}$\,Hz$^{-1}$; \citealt{Ogley_etal_2002, Brocksopp_2004}), where the emission is usually suggested to be thermal, originating from bipolar collimated jets/outflows. Classical symbiotic systems consist of giant stars, which in most cases do not fill their Roche Lobes, transferring material onto WDs;
accretion can proceed via a disk or not, depending on the relative velocity of the red giant wind and the accreting WD \citep{Livio_1992, Mohamed_Podsiadlowski_2007}. Even out of outburst, symbiotic stars can be luminous radio emitters ($L \approx 10^{17}-10^{20}$ erg\,s$^{-1}$\,Hz$^{-1}$), as the hot accreting white dwarf ionizes the substantial circumstellar material expelled by the giant companion \citep{Seaquist_etal_1984, Seaquist_Taylor_1990}.

We showed earlier that if the radio emission is optically thick thermal emission from an ionized gas at a typical brightness temperature of 10$^4$\,--\,10$^5$\,K, the emitting regions should have a radius of $\sim$ 10$^{14}$\,--\,10$^{15}$\,cm to produce the observed flux densities.
If the radio emission observed from V1047~Cen proves to be transient and associated with the 2019 outburst, it might be attributed to outflows moving at velocities $\gtrsim$2000\,km\,s$^{-1}$ (based on the optical line profiles), which will cover a radius of more than 10$^{15}$\,cm in less than 2 months. The variability of the radio source can then be explained as variability in the fast outflow.
If V1047~Cen remains a luminous radio source out of outburst, the radio emission may instead be attributable to circumstellar material, expelled by the potentially evolved companion (assuming an 8.36 day orbital period) and ionized by the hot WD.

\subsection{The nature of the outburst}

The spectroscopic follow up of V1047~Cen, during the rise to peak (first 90 days) is fairly consistent with a DN spectrum; however, the 400 day long outburst, the 14 year gap between the 2005 nova eruption and the 2019 outburst, the luminous radio and optical emission, the dramatic spectral changes and high velocity spectral components all raise questions about the nature of this event.

If the 0.361 day period found in the \textit{TESS} data is the orbital period of the system, it would
challenging to explain the observational features of the system as a DN outburst. However, if the 8.36 day period found in the AAVSO data is the orbital period, the system could potentially host a large accretion disk (maybe one of the largest in currently known CV systems), which has the potential to explain some of the record-breaking observational features, such as the
short gap between the nova eruption and DN outburst and the 400 day outburst. This could potentially place V1047~Cen along with GK~Per and V1017~Sgr in a new class of CVs, which show long-lasting DN outbursts (a few decades after a nova eruption) and are characterized by orbital periods of the order of a few days with mildly evolved secondaries and massive accretion disks. However, some of the observational features of V1047~Cen outburst, such as the long-lasting luminous radio emission and the optical spectral features are challenging to explain in the context of a DN outburst.

Many of the features observed in V1047~Cen resemble those of ``combination novae'' in Z And-like classical nova outbursts \citep{Sokoloski_etal_2006}. This type of event starts with a disk instability outburst, leading to the accretion of a massive accretion disk onto the white dwarf. The accretion burst is large enough to trigger enhanced nuclear burning on the surface of the white dwarf and the ejection of an optically thick outflow/shell (see also \citealt{Munari_2019}).

While classical symbiotic systems are characterized by orbital periods of the order of hundreds of days (\citealt{Mikolajewska_2003}; much longer than the potential 8.36 days orbital period in V1047~Cen), the observational features of classical symbiotic outbursts have many similarities with those of V1047~Cen (the long-lasting optical outburst, the long-lasting and luminous radio emission, the P Cygni spectral features with velocities $\approx$ 1500\,km\,s$^{-1}$, and the satellite spectral features around the Balmer lines).

Therefore, we suggest that the 2019 outburst of V1047~Cen could have indeed started as a DN disk instability type outburst (Based on the early spectroscopy), leading to the accretion of a massive accretion disk onto the white dwarf. The accretion of the large disk might have led to enhanced nuclear shell burning on the surface of the white dwarf, which launched a radiation driven wind/outflow,
as quoted by \citet{Sokoloski_etal_2006} ``in a milder version of the phenomenon that powers classical novae''. The accretion could have been ongoing for several years at a low level prior to the 2019 outburst (see below). This scenario could explain the 400 day outburst duration in the optical (e.g., the 2000's outburst of Z~And lasted for 2 years) and the dramatic brightening and broadening of the Balmer lines.
The light curve bump or re-brightening that occurred $\sim$260 days after the outburst start could mark the stage at which an optically thick outflow started. This is consistent with the development of P Cygni-like absorption features in the optical spectral lines around the same period (Figure~\ref{Fig:PCygni}). The absence of the P Cygni profiles in the spectra obtained on day 407 after the end of the optical outburst (Figure~\ref{Fig:main_spec}; note that the P Cygni profiles were present in the previous epoch on day 303) indicates that the optically thick wind/outflow
dissipated rapidly and became optically thin, which is very similar to the behavior observed in the combination nova of
Z~And \citep{Sokoloski_etal_2006}. The presence of high ionization lines of He, C, O, and N in the late spectra of V1047~Cen (Figure~\ref{Fig:spec_7}), obtained after the end of the optical outburst could indicate ongoing nuclear shell burning, which would strengthen the case for a combination nova.
The same high ionization lines of \eal{N}{V} and \eal{O}{VI} are typically observed in V~Sge like stars---the Milky-Way counterparts of the supersoft sources in the Magellanic Clouds---typically indicate nuclear shell burning (e.g., \citealt{Herbig_etal_1965,Cieslinski_etal_1999,Steiner_etal_1999}).

The combination nova scenario could also explain the long-lasting, luminous radio emission. We showed in Section~\ref{radio_disc} that the flux densities observed in V1047~Cen could be explained as thermal emission from an outflow of ionized gas. Such radio emitting outflows/jets are characteristics of combination nova events (e.g, \citealt{Brocksopp_2004} and references therein). In addition, there is strong evidence for the presence of such collimated outflows or jets in V1047~Cen based on the satellite emission components observed in H$\alpha$ (Figure~\ref{Fig:jet_lines}), which are similar to those observed in Z And-like classical symbiotic outbursts (e.g., \citealt{Munari_etal_2005,Burmeister_etal_2007,Skopal_etal_2013,Tomov_Tomova_2013}).

If strong internal shocks are not present within the outflow, we do not expect bright X-ray emission, because the X-ray emission associated with the nuclear-burning WD would be hidden by the puffed up photosphere and outflow during the combination nova event. This could explain the lack of X-ray detection in the \textit{Swift} individual observations.

The 8.36 day period is shorter than the typical periods of classical symbiotic systems. Therefore, the V1047~Cen system would be intermediate
between typical CVs and classical symbiotics. The latter are also characterized by low mass WDs \citep{Sokoloski_etal_2006}. We do not have insight into the mass of the WD in V1047~Cen. The 2005 nova eruption of V1047~Cen did show a rapid decline in the optical light curve, which is indicative of an ejection of a low-mass envelope; low mass ejecta in turn tend to be associated with eruptions occurring on massive WDs \citep{Yaron_etal_2005}. However, the ejecta velocities measured from the optical spectral lines during the 2005 classical nova are moderate \citep{Liller_2005}, which argues against a massive WD. \citet{Hachisu_Kato_2007} derived a WD mass of around 0.7\,M$_{\odot}$ for V1047~Cen, consistent with a system with a low mass WD.  However, the values derived by \citet{Hachisu_Kato_2007} are based on poor ``eye-fitting'' of the visual light curve to their models, and therefore are uncertain.

\citet{Sokoloski_etal_2006} suggest that pre-existing quasi-steady shell burning on the WD surface
is necessary to trigger a combination nova. Given the scarcity of observations between the 2005 nova eruption and 2019 outburst, it is not possible to definitively know if residual nuclear burning on the surface of the WD has been ongoing after the 2005 eruption. However, OGLE photometry between 2013 and 2019 showed that the system has not
returned to the pre-2005-nova brightness and the colors of the system during this period suggest  emission from a hot source (see Section~\ref{sec:secondary}). Both of these indicate that nuclear burning on the WD surface could have been ongoing between 2005 and 2019.

The absence of the system from the DSS images before the 2005 nova eruption, places an upper limit of $V \approx$ 20.5 on the system during quiescence. This implies that the absolute magnitude of the system before 2005 is dimmer than 5\,mag, leaving a small chance for the companion star to be an evolved sub-giant star, unless both the reddening and distance to the system are not well constrained. This challenges the 8.36 days period and the potential combination nova scenario, and raises more questions about the nature of the outburst.

That being said, a DN outburst seems unlikely to explain all the observational features of the 2019 outburst of V1047~Cen and a phenomenon similar to a combination nova seems to agree better with the observations. If so, this would be the first combination nova observed in a system with an intermediate orbital period between typical CVs and classical symbiotics.

One other possible explanation of the 2019 outburst is that the system experienced a non-ejection nova event.
The extended grid of nova models by \citet{Yaron_etal_2005} shows that some combination of parameters might lead to a thermonuclear runaway without ejecting material
(see also \citealt{Fujimoto_1982,Shara_etal_1977}). \citet{Yaron_etal_2005} suggest that such outbursts cause only a slow increase in luminosity,
followed by a slow decay, particularly in the case of low mass WDs where the timescales of the rise and decay could be thousands of days. The supersoft X-ray transient ASASSN-16oh has been suggested to be a thermonuclear runaway event without mass ejection \citep{Hillman_etal_2019}. The system showed a slowly rising (1585 days) and declining (268 days) optical light curve, with an amplitude of less than 4 magnitude. While there are similarities between the optical light curves of ASASSN-16oh and V1047~Cen, the former has been detected as a supersoft X-ray source at the distance of the Small Magellanic Cloud. This is very different from the case of V1047~Cen, which was not detected in individual epochs by \textit{Swift} during the 2019 outburst despite our extensive monitoring.
In case of a thermonuclear runaway, the luminosity of the supersoft source is expected to be of the order of 10$^{38}$\,erg\,s$^{-1}$, which would have been easily observed at a distance of 3.2\,kpc by \textit{Swift}. The supersoft emission of ASASSN-16oh was characterized by a luminosity of around 10$^{37}$\,erg\,s$^{-1}$, which is lower than that expected
from the models of non-ejecting thermonuclear runaways. This low luminosity has been attributed to an optically thick accretion disk hiding most of the WD surface. While a massive accretion disk could also be blocking some of the emission from the WD surface in case of V1047~Cen, it is a less likely possibility given the \textit{Swift} non-individual-detections over several months.

\section{Conclusions}
\label{Conc}
We have presented multi-wavelength observations of the 2019 outburst in the 2005 classical nova V1047~Cen. The outburst amplitude reached around 6 magnitudes in the optical and lasted for more than 400 days. We derive a distance of 3.2$\pm$0.2\,kpc to the system and an peak absolute magnitude of the outburst $M_V = -2$. The first spectra we obtained of the system around 74 days after the start of the outburst were consistent with a disk instability dwarf nova event in a classical nova system. If V1047 Cen is a DN, the event would be a record breaker, making it the longest DN outburst on record, the shortest gap between a nova and DN in a CV system, and the most luminous optical and radio DN outburst. We find a potential orbital period of 8.36 days in the optical photometry, which could explain the long-lasting outburst and the 14 year gap with the nova eruption, however, the observational features of the outburst across the electromagnetic spectrum point towards a phenomenon more exotic than just a DN outburst.

We therefore suggest that the event is an outburst resembling those observed in classical symbiotic systems which are known as combination novae. In this scenario, the outburst might have started with a DN-like instability in a massive accretion disk, which then triggered enhanced nuclear shell burning on the surface of the white dwarf and eventually led to an optically thick, radiation driven wind/outflow. This scenario fits well the 400 day outburst duration, the dramatic changes in the optical line profiles and the $>$ 2000\,km\,s$^{-1}$ velocities inferred from these profiles, the P Cygni line profiles which appeared several months after the outburst start, the high peak optical brightness, and most importantly the long-lasting, superluminous radio emission, which likely originates from ionized gas (thermal emission) in collimated bipolar flows, which are characteristic of combination nova events. Strong evidence for such outflows can be seen in the late optical spectral line profiles. Mid-infrared observations also indicate that pre-existing dust --- likely formed during the 2005 classical nova eruption --- has been heated by the radiation from the 2019 outburst.

Therefore, the 2019 outburst of V1047~Cen could either be a record breaking DN outburst or more likely a combination-nova-like event -- a phenomenon observed for the first time in an intermediate system between a CV and a symbiotic system, which has undergone a classical nova eruption.

\section*{Acknowledgments}

We thank B. Schaefer for useful discussion. We thank the AAVSO observers from around the world who contributed their magnitude measurements to the AAVSO International Database used in this work.

EA, LC, and KVS acknowledge NSF award AST-1751874, NASA award 11-Fermi 80NSSC18K1746, NASA award 16-Swift 80NSSC21K0173, and a Cottrell fellowship of the Research Corporation. JS was supported by the Packard Foundation. DAHB gratefully acknowledges the receipt of research grants from the National Research Foundation (NRF) of South Africa. PAW kindly acknowledges the National Research Foundation and the University of Cape Town. KLP acknowledges funding from the UK Space Agency. Nova research at Stony Brook University has been made possible by NSF award AST-1611443. MG is supported by the EU Horizon 2020 research and innovation programme under grant agreement No 101004719.
A part of this work is based on observations made with the Southern African Large Telescope (SALT), with the Large Science Programme on transients 2018-2-LSP-001 (PI: DAHB). Polish participation in SALT is funded by grant no. MNiSW DIR/WK/2016/07. This paper was partially based on observations obtained at the Southern Astrophysical Research (SOAR) telescope, which is a joint project of the Minist\'{e}rio da Ci\^{e}ncia, Tecnologia e Inova\c{c}\~{o}es (MCTI/LNA) do Brasil, the US National Science Foundation's NOIRLab, the University of North Carolina at Chapel Hill (UNC), and Michigan State University (MSU).
The OGLE project has received funding from the National Science Centre, Poland, grant MAESTRO 2014/14/A/ST9/00121 to AU.
The CHIRON and $Andicam$ instruments are managed by the Todd Henry and the SMARTS Consortium. The MeerKAT telescope is operated by the South African Radio Astronomy Observatory, which is a facility
of the National Research Foundation, an agency of the Department of Science and Innovation. We acknowledge the use of the ilifu cloud computing facility - \url{www.ilifu.ac.za}, a partnership between the University of Cape Town, the University of the Western Cape, the University of Stellenbosch, Sol Plaatje University, the Cape Peninsula University of Technology and the South African Radio Astronomy Observatory. The ilifu facility is supported by contributions from the Inter-University Institute for Data Intensive Astronomy (IDIA - a partnership between the University of Cape Town, the University of Pretoria and the University of the Western Cape), the Computational Biology division at UCT and the Data Intensive Research Initiative of South Africa (DIRISA).
This research is based in part on observations obtained at the international Gemini Observatory, a program of NSF's NOIRLab, which is managed by the Association of Universities for Research in Astronomy (AURA) under a cooperative agreement with the National Science Foundation. on behalf of the Gemini Observatory partnership: the National Science Foundation (United States), National Research Council (Canada), Agencia Nacional de Investigaci\'{o}n y Desarrollo (Chile), Ministerio de Ciencia, Tecnolog\'{i}a e Innovaci\'{o}n (Argentina), Minist\'{e}rio da Ci\^{e}ncia, Tecnologia, Inova\c{c}\~{o}es e Comunica\c{c}\~{o}es (Brazil), and Korea Astronomy and Space Science Institute (Republic of Korea). DPKB is supported by a CSIR Emeritus Scientist grant-in-aid, which is being hosted by the Physical Research Laboratory, Ahmedabad.
This work is based in part on observations made with the
NASA/DLR Stratospheric Observatory for Infrared Astronomy (SOFIA). SOFIA is jointly operated by the Universities Space Re- search Association, Inc. (USRA), under NASA contract NNA17BF53C, and the Deutsches SOFIA Institut (DSI) under DLR contract 50 OK 0901 to the University of Stuttgart. Financial support for CEW/RDG
related to this work was provided by NASA through award SOF07-0005 issued by USRA to the University of Minnesota. This publication makes use of data products from the Near-Earth Object Wide-field Infrared Survey Explorer (NEOWISE), which is a joint project of the Jet Propulsion Laboratory/California Institute of Technology and the University of Arizona. NEOWISE is funded by the National Aeronautics and Space Administration. VARMR acknowledges financial support from the Funda\c{c}\~{a}o para a Ci\^encia e a Tecnologia (FCT) in the form of an exploratory project of reference IF/00498/2015/CP1302/CT0001, and from the Minist\'erio da Ci\^encia, Tecnologia e Ensino Superior (MCTES) through national funds and when applicable co-funded EU funds under the project UIDB/EEA/50008/2020, and supported by Enabling Green E-science for the Square Kilometre Array Research Infrastructure (ENGAGE-SKA), POCI-01-0145- FEDER-022217, and PHOBOS, POCI-01-0145-FEDER029932, funded by Programa Operacional Competitividade e Internacionaliza\c{c}\~ao (COMPETE 2020) and FCT, Portugal.

\subsubsection*{#1}}

\pagestyle{headings}
\markright{Reference sheet: \texttt{natbib}}

\usepackage{shortvrb}
\MakeShortVerb{\|}

\begin{document}
\thispagestyle{plain}

\newcommand{\textsc{Bib}\TeX}{\textsc{Bib}\TeX}
\newcommand{\texttt{#1}\def\filedate{#2}\def\fileversion{#3}}}{\texttt{#1}\def\filedate{#2}\def\fileversion{#3}}}
\begin{center}{\bfseries\Large
    Reference sheet for \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}\ usage}\\
    \large(Describing version \fileversion\ from \filedate)
\end{center}

\begin{quote}\slshape
For a more detailed description of the \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}\ package, \LaTeX\ the
source file \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}\texttt{.dtx}.
\end{quote}
\head{Overview}
The \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}\ package is a reimplementation of the \LaTeX\ |\cite| command,
to work with both author--year and numerical citations. It is compatible with
the standard bibliographic style files, such as \texttt{plain.bst}, as well as
with those for \texttt{harvard}, \texttt{apalike}, \texttt{chicago},
\texttt{astron}, \texttt{authordate}, and of course \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}.

\head{Loading}
Load with |\usepackage[|\emph{options}|]{|\texttt{#1}\def\filedate{#2}\def\fileversion{#3}}|}|. See list of
\emph{options} at the end.

\head{Replacement bibliography styles}
I provide three new \texttt{.bst} files to replace the standard \LaTeX\
numerical ones:
\begin{quote}\ttfamily
 plainnat.bst \qquad abbrvnat.bst \qquad unsrtnat.bst
\end{quote}
\head{Basic commands}
The \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}\ package has two basic citation commands, |\citet| and
|\citep| for \emph{textual} and \emph{parenthetical} citations, respectively.
There also exist the starred versions |\citet*| and |\citep*| that print
the full author list, and not just the abbreviated one.
All of these may take one or two optional arguments to add some text before
and after the citation.
\begin{quote}
\begin{tabular}{l@{\quad$\Rightarrow$\quad}l}
  |\citet{jon90}| & Jones et al. (1990)\\
  |\citet[chap.~2]{jon90}| & Jones et al. (1990, chap.~2)\\[0.5ex]
  |\citep{jon90}| & (Jones et al., 1990)\\
  |\citep[chap.~2]{jon90}| & (Jones et al., 1990, chap.~2)\\
  |\citep[see][]{jon90}| & (see Jones et al., 1990)\\
  |\citep[see][chap.~2]{jon90}| & (see Jones et al., 1990, chap.~2)\\[0.5ex]
  |\citet*{jon90}| & Jones, Baker, and Williams (1990)\\
  |\citep*{jon90}| & (Jones, Baker, and Williams, 1990)
\end{tabular}
\end{quote}
\head{Multiple citations}
Multiple citations may be made by including more than one
citation key in the |\cite| command argument.
\begin{quote}
\begin{tabular}{l@{\quad$\Rightarrow$\quad}l}
  |\citet{jon90,jam91}| & Jones et al. (1990); James et al. (1991)\\
  |\citep{jon90,jam91}| & (Jones et al., 1990; James et al. 1991)\\
  |\citep{jon90,jon91}| & (Jones et al., 1990, 1991)\\
  |\citep{jon90a,jon90b}| & (Jones et al., 1990a,b)
\end{tabular}
\end{quote}

\head{Numerical mode}
These examples are for author--year citation mode. In numerical mode, the
results are different.
\begin{quote}
\begin{tabular}{l@{\quad$\Rightarrow$\quad}l}
  |\citet{jon90}| & Jones et al. [21]\\
  |\citet[chap.~2]{jon90}| & Jones et al. [21, chap.~2]\\[0.5ex]
  |\citep{jon90}| & [21]\\
  |\citep[chap.~2]{jon90}| & [21, chap.~2]\\
  |\citep[see][]{jon90}| & [see 21]\\
  |\citep[see][chap.~2]{jon90}| & [see 21, chap.~2]\\[0.5ex]
  |\citep{jon90a,jon90b}| & [21, 32]
\end{tabular}
\end{quote}
\head{Suppressed parentheses}
As an alternative form of citation, |\citealt| is the same as |\citet| but
\emph{without parentheses}. Similarly, |\citealp| is |\citep| without
parentheses. Multiple references, notes, and the starred variants
also exist.
\begin{quote}
\begin{tabular}{l@{\quad$\Rightarrow$\quad}l}
  |\citealt{jon90}| & Jones et al.\ 1990\\
  |\citealt*{jon90}| & Jones, Baker, and Williams 1990\\
  |\citealp{jon90}| & Jones et al., 1990\\
  |\citealp*{jon90}| & Jones, Baker, and Williams, 1990\\
  |\citealp{jon90,jam91}| & Jones et al., 1990; James et al., 1991\\
  |\citealp[pg.~32]{jon90}| & Jones et al., 1990, pg.~32\\
  |\citetext{priv.\ comm.}| & (priv.\ comm.)
\end{tabular}
\end{quote}
The |\citetext| command
allows arbitrary text to be placed in the current citation parentheses.
This may be used in combination with |\citealp|.
\head{Partial citations}
In author--year schemes, it is sometimes desirable to be able to refer to
the authors without the year, or vice versa. This is provided with the
extra commands
\begin{quote}
\begin{tabular}{l@{\quad$\Rightarrow$\quad}l}
  |\citeauthor{jon90}| & Jones et al.\\
  |\citeauthor*{jon90}| & Jones, Baker, and Williams\\
  |\citeyear{jon90}|   & 1990\\
  |\citeyearpar{jon90}| & (1990)
\end{tabular}
\end{quote}
\head{Forcing upper cased names}
If the first author's name contains a \textsl{von} part, such as ``della
Robbia'', then |\citet{dRob98}| produces ``della Robbia (1998)'', even at the
beginning of a sentence. One can force the first letter to be in upper case
with the command |\Citet| instead. Other upper case commands also exist.
\begin{quote}
\begin{tabular}{rl@{\quad$\Rightarrow$\quad}l}
  when & |\citet{dRob98}| & della Robbia (1998) \\
  then & |\Citet{dRob98}| & Della Robbia (1998) \\
     &   |\Citep{dRob98}| & (Della Robbia, 1998) \\
     &   |\Citealt{dRob98}| & Della Robbia 1998 \\
     &   |\Citealp{dRob98}| & Della Robbia, 1998 \\
     &   |\Citeauthor{dRob98}| & Della Robbia
\end{tabular}
\end{quote}
These commands also exist in starred versions for full author names.

\head{Citation aliasing}
Sometimes one wants to refer to a reference with a special designation,
rather than by the authors, i.e. as Paper~I, Paper~II. Such aliases can be
defined and used, textual and/or parenthetical with:
\begin{quote}
\begin{tabular}{lcl}
  |\defcitealias{jon90}{Paper~I}|\\
  |\citetalias{jon90}| & $\Rightarrow$ & Paper~I\\
  |\citepalias{jon90}| & $\Rightarrow$ & (Paper~I)
\end{tabular}
\end{quote}
These citation commands function much like |\citet| and |\citep|: they may
take multiple keys in the argument, may contain notes, and are marked as
hyperlinks.
\head{Selecting citation style and punctuation}
Use the command |\bibpunct| with one optional and 6 mandatory arguments:
\begin{enumerate}
\item the opening bracket symbol, default = (
\item the closing bracket symbol, default = )
\item the punctuation between multiple citations, default = ;
\item the letter `n' for numerical style, or `s' for numerical superscript
      style, any other letter for
      author--year, default = author--year;
\item the punctuation that comes between the author names and the year
\item the punctuation that comes between years or numbers when common author
      lists are suppressed (default = ,);
\end{enumerate}

The optional argument is the character preceding a post-note, default is a
comma plus space. In redefining this character, one must include a space if
one is wanted.

Example~1, |\bibpunct{[}{]}{,}{a}{}{;}| changes the output of
\begin{quote}
  |\citep{jon90,jon91,jam92}|
\end{quote}
into [Jones et al. 1990; 1991, James et al. 1992].

Example~2, |\bibpunct[; ]{(}{)}{,}{a}{}{;}| changes the output of
\begin{quote}
  |\citep[and references therein]{jon90}|
\end{quote}
into (Jones et al. 1990; and references therein).

\head{Other formatting options}
Redefine |\bibsection| to the desired sectioning command for introducing
the list of references. This is normally |\section*| or |\chapter*|.

Define |\bibpreamble| to be any text that is to be printed after the heading but
before the actual list of references.

Define |\bibfont| to be a font declaration, e.g.\ |\small| to apply to
the list of references.

Define |\citenumfont| to be a font declaration or command like |\itshape|
or |\textit|.

Redefine |\bibnumfmt| as a command with an argument to format the numbers in
the list of references. The default definition is |[#1]|.

The indentation after the first line of each reference is given by
|\bibhang|; change this with the |\setlength| command.

The vertical spacing between references is set by |\bibsep|; change this with
the |\setlength| command.

\head{Automatic indexing of citations}
If one wishes to have the citations entered in the \texttt{.idx} indexing
file, it is only necessary to issue |\citeindextrue| at any point in the
document. All following |\cite| commands, of all variations, then insert
the corresponding entry to that file. With |\citeindexfalse|, these
entries will no longer be made.

\head{Use with \texttt{chapterbib} package}

The \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}\ package is compatible with the \texttt{chapterbib} package
which makes it possible to have several bibliographies in one document.

The package makes use of the |\include| command, and each |\include|d file
has its own bibliography.

The order in which the \texttt{chapterbib} and \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}\ packages are loaded
is unimportant.

The \texttt{chapterbib} package provides an option \texttt{sectionbib}
that puts the bibliography in a |\section*| instead of |\chapter*|,
something that makes sense if there is a bibliography in each chapter.
This option will not work when \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}\ is also loaded; instead, add
the option to \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}.

Every |\include|d file must contain its own
|\bibliography| command where the bibliography is to appear. The database
files listed as arguments to this command can be different in each file,
of course. However, what is not so obvious, is that each file must also
contain a |\bibliographystyle| command, \emph{preferably with the same
style  argument}.
\head{Sorting and compressing citations}
Do not use the \texttt{cite} package with \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}; rather use one of the
options \texttt{sort} or \texttt{sort\&compress}.

These also work with author--year citations, making multiple citations appear
in their order in the reference list.

\head{Long author list on first citation}
Use option \texttt{longnamesfirst} to have first citation automatically give
the full list of authors.

Suppress this for certain citations with |\shortcites{|\emph{key-list}|}|,
given before the first citation.

\head{Local configuration}
Any local recoding or definitions can be put in \texttt{#1}\def\filedate{#2}\def\fileversion{#3}}\texttt{.cfg} which
is read in after the main package file.

\head{Options that can be added to \texttt{\char`\\ usepackage}}
\begin{description}
\item[\ttfamily round] (default) for round parentheses;
\item[\ttfamily square] for square brackets;
\item[\ttfamily curly] for curly braces;
\item[\ttfamily angle] for angle brackets;
\item[\ttfamily colon] (default) to separate multiple citations with
     colons;
\item[\ttfamily comma] to use commas as separaters;
\item[\ttfamily authoryear] (default) for author--year citations;
\item[\ttfamily numbers] for numerical citations;
\item[\ttfamily super] for superscripted numerical citations, as in
     \textsl{Nature};
\item[\ttfamily sort] orders multiple citations into the sequence in
     which they appear in the list of references;
\item[\ttfamily sort\&compress] as \texttt{sort} but in addition multiple
     numerical citations are compressed if possible (as 3--6, 15);
\item[\ttfamily longnamesfirst] makes the first citation of any reference
     the equivalent of the starred variant (full author list) and subsequent
     citations normal (abbreviated list);
\item[\ttfamily sectionbib] redefines |\thebibliography| to issue
     |\section*| instead of |\chapter*|; valid only for classes with a
     |\chapter| command; to be used with the \texttt{chapterbib} package;
\item[\ttfamily nonamebreak] keeps all the authors' names in a citation on
     one line; causes overfull hboxes but helps with some \texttt{hyperref}
     problems.
\end{description}
\end{document}

redpajama_set_name: RedPajamaArXiv