Near-infrared Studies of Nova V1674 Herculis: A Shocking Record Breaker

Nova V1674 Her (TCP J18573095+1653396) was discovered by Seidji Ueda on 2021 June 12.5484 UT, at an apparent visual magnitude of 8.4; it attained a peak brightness of V ∼ 6 mag (Munari et al. 2021; Quimby et al. 2021). V1674 Her has been intensely observed since discovery, with optical spectroscopy at high cadence (Woodward et al. 2021c, and references therein), near-infrared (NIR) spectroscopy (Woodward et al. 2021a, and references therein), and ultraviolet photometry and spectroscopy (Kuin et al. 2021). First detected in X-rays on June 14.41 UT, V1674 Her began to show supersoft (SS) X-ray emission near July 1 (Page et al. 2021). Drake et al. (2021) give a detailed discussion of its early X-ray evolution. Radio observations reveal nonthermal emission (Paragi et al. 2021; Sokolovsky et al. 2021). V1674 Her is also a γ-ray source (Maccarone et al. 2021; Li 2021b, and references therein). NICER and Chandra X-ray observations (Pei et al. 2021), as well as pre-outburst archival r-band photometry (Mroz et al. 2021), reveal a ∼501 s oscillation period that is interpreted to be the spin period of a WD in an intermediate polar system.

V1674 Her is distinctive in many ways. Classification-based optical spectra obtained within days of outburst (Woodward et al. 2021c, and references therein) indicate that V1674 Her is of the Fe ii class (Williams 1992; 2012), with several Fe ii lines appearing, most prominent of which was the Fe ii(42) multiplet (however, the possibility of a hybrid Fe iib classification based on a Fe ii to He/N transition around ∼day 5.5 may be noted; Woodward et al. 2021b). At this early epoch the peak of the Hα profile was corrugated, comprising around 9–10 distinct subpeaks separated from each other in velocity by amounts ranging from 180 to 770 km s⁻¹, while exhibiting an overall FWHM ≃ 5850 km s⁻¹. The OI 8446 Å line, dominantly excited by Lyman-β fluorescence, was also prominent, having a FWHM of 5650 km s⁻¹. V1674 Her also showed coronal line emission remarkably early. The detection of coronal lines at 11 days past maximum light is possibly the earliest reported detection of such lines in a nova (see Benjamin & Dinerstein 1990).

On day 28, Wagner et al. (2021) report the detection of broad [Ne v] 3426 Å, as well as [Ne iv] at 4714 Å, and 4724 Å arguing that the presence of these strong neon emission lines is likely attributable to overabundances of neon. If substantiated, V1674 Her would be a member of the “neon nova” class that includes QU Vul, V838 Her, and V1974 Cyg among others.

Here we discuss the exceptionally rapid near-infrared (NIR) spectral evolution of V1674 Her during the first ∼70 days post-eruption, focusing our analyses on the very strong coronal emission that provides direct measures of gas temperatures, ejecta abundances, and insight into the TNR process in novae.

From the AAVSO database (Kafka 2021), peak brightness was reached on MJD 59,377.9605 (UTC 2021 June 12.96) at V = 6.14, in agreement with Munari et al. (2021), who find V = 6.18 on June 12.90 UT. We adopt June 12.96 UT (MJD 59,377.96) as our time origin, t₀, and V = 6.14 at maximum. Based on this t₀, and the AAVSO light curve in Figure 1, t₂, t₃, and t₆ are 1 day, 2.2 days, and 14 days, respectively.

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When the t₂ and t₃ times are considered together, V1674 Her is again remarkable. It is the fastest Galactic nova known to date, of all the compiled, detailed light curves (Strope et al. 2010) surpassing V838 Her (t₂ = 1.0, t₃ = 3.0) and the recurrent nova U Sco (t₂ = 1.0, t₃ = 4.0). Optical (Wagner et al. 2021; Woodward et al. 2021c) and NIR (Woodward et al. 2021a) spectroscopy bear this out, with emission line widths and P Cygni profiles indicating ejection velocities ∼6000 km s⁻¹. Such extreme eruptions are characteristic of a TNR that occurred on a WD with a mass very close to the Chandrasekhar limit (Starrfield et al. 2020).

The reddening, E_B−V , is in the range 0.52–0.58 based on the 6613 Å diffuse interstellar band feature, clearly seen in the high-resolution (R = 220K) optical spectra (Woodward et al. 2021c). From these high-resolution spectra, the interstellar resonance K i line at 7699 Å is measured to have an equivalent width of 0.1349 Å, which gives a reddening E_B−V = 0.52 using the calibration by Munari & Zwitter (1997). Munari et al. (2021), from the K i 7699 Å line, measure E_B−V = 0.55. The observed color of B–V = 0.72 and 0.55 at t₀ and t₂, respectively, when compared with the intrinsic colors of novae expected at these early evolutionary times (van den Bergh & Younger 1987), yields E_B−V values in the range 0.49–0.57. We adopt E_B−V = 0.55.

Adopting the method described in Banerjee et al. (2018, and references therein), an independent estimate of the reddening, together with a distance estimate, can simultaneously be obtained from the intersection of the extinction versus distance curve in the direction to the nova (Marshall et al. 2006), and the curve defined by

where m_v = 6.14 at maximum and M_V = −9.05 ± 0.5 (for t₂ = 1 d derived from the MMRD relation of Della Valle & Livio 1995). The intersection of these two curves yields a distance of kpc and visual extinction A_V = 1.82 ± 0.33, or E_B−V = 0.59±0.11, assuming that the total-to-selective extinction is 3.10. As a cross-check, applying this method to a similarly fast nova V838 Her (t₂ = 2 d), gives distance and A_V estimates of 3.83 kpc and 1.40 magnitudes, respectively, consistent with other compiled measurements (Harrison & Stringfellow 1994). Gaia estimates for the distance to V1674 Her (Bailer-Jones et al. 2021, 2018) have large uncertainties (the parallax measurements are at the limit of usefulness), but the central values are consistent with the latter kpc estimate.

The spectral energy distribution (SED) at maximum, constructed using BV RI values (Munari et al. 2021) dereddened by E_B−V = 0.55, yield a blackbody fit with T = 9900 K. The outburst luminosity, assuming d = 4.7 kpc, is L_out = 3.7 × 10⁵ L_⊙, which is ∼9 times the Eddington luminosity of a 1.3 M_⊙ WD.

Using the dereddened fluxes of Pa-β and Br-γ, we ascertain that their relative strengths follow Case B predictions on day 11.51, within the parameter space expected to represent the ejecta, viz. T = 5000–20,000 K, and electron density n_e ∼ 10⁶–10⁹ cm⁻³. Other H lines could not be used because of blending with neighboring lines, leading to inaccurate line flux estimates. We assume the ejecta density on day 11.51 to be n_e = 10^6.6 cm⁻³ since the [Al vi] 3.66 μm line, which has a critical density of n_e = 10^6.6 cm⁻³ at T = 10,000 K (Greenhouse et al. 1993), is detected on this day.

The ejecta mass is given by

where f is the dereddened flux in an H line in W m⁻², is the emissivity in the line in erg s⁻¹ cm³ (Storey & Hummer 1995), and n_e is the electron density per cm³. Using the Pa-β and Br-γ lines, and assuming kpc and a pure H composition, the ejecta mass is M_⊙, wherein the upper limit on the mass is obtained by choosing n_e = 10^6.6 cm⁻³.

The profiles of four nebular lines (two H i and two He i) on day 5.64 are shown in Figure 3(a) to illustrate the complexities of the line profiles. The similarities of the details in each of these profiles are remarkable. In particular, there are identifiable features at radial velocities –2480, –1820, –60, 730, 1140, 1530, and 2310 km s⁻¹, each with an uncertainty of ±100 km s⁻¹. The velocity structure seen at an epoch (day 50.43) when the coronal lines are strong, Figure 3b, is similar.

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The coronal line fluxes are listed in Table 2. Uncertainties in the flux determinations are estimated to be ≤ ±20% for the stronger lines, and as much as ±40% for the weaker. The table also lists the ionization potential (IP) of the lower ion (thus, the entry for [S ix] gives the IP of S viii). The parameter kT_* denotes the temperature T_* at which half the photons emitted by the corresponding blackbody can ionize the lower ion; we take this to imply the temperature of the ionizing source that is capable of producing the requisite ion.

Table 2. V1674 Her Coronal Line Dereddended Fluxes

Ion	λ		Transition		IP	kT*	Ψ	Line Fluxes ( × 10−17 W m−2)
	(μm)		u − ℓ		(eV)	(eV)		11.51d	24.38d	33.37d	50.43d	58.39d	68.32d
[S ix] a	1.2523		3P1 − 3P2		329	137	—	1736	98	49	9.2	5.0	1.7
[Si x]	1.4305		2PP		351	146	—	<2.3	1878	3437	1307	774	82
[Si ix]	1.5599		3P2 − 3P0		304	127	—	11.7	32	13	2.3	0.7	<
[P viii] b	1.7361		3P1 − 3P2		264	110	—	231	32	15	2.1	1.4	<
[Si xi]	1.9320		3PPo 1		401	167	—	<8.5			<1.8	< 0.2
[Si vi]	1.9650		2PP		167	70	0.423	1190	2274	927	185	133	95
[Al ix] c	2.0400		2PP		285	119	—	<4	1035	1004	296	216	72
[Ca viii]	2.3211		2PP		127	53	<5.88	5.1	1.25		<0.6	< 0.1
[Si vii]	2.4833		3P1 − 3P2		205	86	0.695	218	1990	1107	283	200	121
[Al v]	2.9052		2PP		120	50	0.456	51	287	—		—	—
[Mg viii]	3.0280		2PP		225	94	—	<5	288	—	103	—	—
[Ca iv] d	3.2067		2PP		51	21	—			—		—	—
[Al vi]	3.6597		3P1 − 3P2		154	64	1.365	47		—		—	—
[Al viii]	3.6900		3P2 − 3P1		242	101	—			—		—	—
[Si ix]	3.9357		3P1 − 3P0		304	127	—	<1.5	614	—	468	—	—

Notes. Upper limits are 1σ. IP is the ionization potential of the lower ionization state. See the text for explanation of kT_*. Table entry with “—” indicates that the observation did not cover the wavelength range.

^a Heavily blended with He i 1.253 μm. ^b Blended with H i 10–4 1.737 μm. ^c Blended with He i 2.058 μm. ^d Not strictly coronal according to the definition of Greenhouse et al. (1990). The effective collision strengths (Ψ) are the collisional strengths Ω averaged over a thermal electron distribution and are taken from the IRON Project (Hummer et al. 1993; Badnell et al. 2006) online database (http://cdsweb.u-strasbg.fr/tipbase/home.html) interpolated to a temperature of 10^5.57±0.05.

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The time dependence of the silicon coronal lines, among the strongest and best recorded coronal lines in the spectrum, is shown in Figure 1. It is clear from Table 2 and Figure 1 that coronal line emission was first unambiguously detected, via the [Si vi] 1.96 μm, [Si vii] 2.48 μm, and the [Al vi] 3.66 μm lines, on June 24 (day 11.51), well before the SS phase began (Page et al. 2021; Drake et al. 2021). In this respect V1674 Her is again a record breaker. Both V1500 Cyg and V838 Her, two of the fastest novae with t₂ times of 2 and 1 day(s) respectively, were first detected to show coronal emission at 29 and 17 days, respectively, past maximum light (Chandrasekhar et al. 1993; Benjamin & Dinerstein 1990).

The presence of coronal lines so early in the evolution implies that, despite the presence of a hot central WD remnant, the initial coronal emission could not have been caused by photoionization. For example, [Si vii] 2.48 μm is present on day 11.51 even though the ionization of Si vi requires 205 eV, far exceeding the temperature of the ionizing source (kT_* = 86 eV) that would emit sufficient photons, and is well in excess of the source temperature as measured by The Neil Gehrels Swift Observatory (see Figure 1).

We therefore conclude that the coronal emission arose from collisional ionization by particles that gained sufficient energy from shock heating. This contradicts the commonly assumed view that, when the central source is sufficiently hot, photoionization is the principal mechanism for coronal line generation (see Benjamin & Dinerstein 1990). Plausible support for shocks comes from the multipeaked line profiles shown in Figure 3, which were even more clearly resolved into a large number of substructures in very high resolution (R = 220,000) spectra (Woodward et al. 2021c). This points to parcels of gas in the ejecta having widely different velocities. Collisions between such parcels would not only lead to shock heating, but might also account for the gamma-ray emission seen in the early stages of V1674 Her (Li 2021a).

To make estimates, we suppose that the ejecta initially are fully ionized and have a temperature ∼10⁴ K, typical for CNe. For a gas of solar composition, the rms ion speed is V_rms ∼ 20 km s⁻¹. Figure 3 suggests that there must exist in the ejecta parcels of gas with relative speed V_rel ∼ 1000 km s⁻¹. Thus, V_rel/V_rms ∼ 50, implying that the gas is shocked, and that the shock is strong.

Under these circumstances, the temperature of the shocked gas could be as high as

which is ∼14 × 10⁶ K, or kT_shock ∼ 1200 eV. This is more than sufficient to account for the early presence of coronal lines.

However, photoionization may still have a role in the coronal line emission. As the SS phase commences, an additional hard ionizing source of radiation becomes available (in addition to shock heating). The impact of this source is considerable, especially in the case of V1674 Her since it has been shown to have the hottest galactic SS source of all well-documented SS sources (see Page & Osborne 2014). As the bottom panel of Figure 1 shows, the degree of ionization for the Si ion slowly increases with time as the SS source becomes hotter. The strength of the [Si x] 1.43 μm line (IP = 351 eV), the strongest coronal line in the spectrum, likely reflects the extreme high temperature of the central WD surface. The overall picture that emerges is that both collisional ionization and photoionization are operational, but the key result from this work is that collisional ionization preceded photoionization, since coronal lines were seen about a week before the SS phase was detected. Thus, a problem that is hard to resolve, whether collisional ionization or photoionization is responsible for coronal emission in CNe, can be resolved as has been done here by correlating the temporal sequence in which the X-ray SS phase and the NIR coronal emission appear.

The coronal lines were first detected on day 11.51, and likely the coronal emission was entirely the result of collisions. We can, therefore, estimate the gas temperature using the fluxes of the various silicon lines at this time. We estimate the relative proportions of Si ions as a function of temperature using the prescription of Greenhouse et al. (1990), the ionization equilibrium data in Arnaud & Rothenflug (1985), and the effective collisional strength values Ψ listed in Table 2. We find the gas temperature T_gas ≃ 10^5.57±0.05 K, (kT_gas ∼ 27 keV). This value is consistent with the upper limit on the flux in the [Si x] 1.43 μm line, which requires that T_gas be less than 10^5.7 K. This value of T_gas is similar to that found in the shock-excited gas around the recurrent nova RS Oph, in which the ejected material runs into, and shocks, the secondary wind (Evans et al. 2007; Banerjee et al. 2009). We assume T_gas ≃ 10^5.5 K in what follows.

Greenhouse et al. (1990, 1993) describe a method for estimating the abundance ratios for selected species. For the present work we use dereddened fluxes from Table 2 for lines where there are no blends. Using the [Al v] 2.90 μm line, we find that n(Al)/n(Si) = 0.17 using the [Si vii] line, and = 0.14 derived from the [Si vi] line. A value of n(Al)/n(Si) = 4.5 × 10⁻² is derived from the [Al vi] and [Si vi] lines. Hence the mean n(Al)/n(Si) is ≃0.11. Using the [Ca viii] line, we get a more consistent result, viz. n(Ca)/n(Si) = 2.6 × 10⁻³ from the [Si vi] line, and = 2.2 × 10⁻³ from the [Si vii] line. The latter values represent lower limits as collision strengths were not explicitly computed at the assumed temperature for the [Ca viii] line. The solar values are 0.081 (Al/Si) and 0.062 (Ca/Si), respectively (Asplund et al. 2021). We tentatively conclude that aluminum is near solar abundance, and calcium is underabundant relative to silicon, with respect to solar values. A more detailed abundance determination awaits future modeling efforts.

Synoptic multiwavelength studies of this remarkable nova continue as the system returns to quiescence.