CONSTRAINTS ON THE PROGENITOR SYSTEM AND THE ENVIRONS OF SN 2014J FROM DEEP RADIO OBSERVATIONS

Type Ia supernovae (SNe) are the end-products of white dwarfs with a mass approaching the Chandrasekhar limit, which results in a thermonuclear explosion of the star. In addition to their use as cosmological distance indicators (e.g., Riess et al. 1998; Perlmutter et al. 1999), Type Ia SNe (henceforth SNe Ia) are a major contributor to the chemical evolution of galaxies. It is therefore unfortunate that we do not yet know what makes an SN Ia. This lack of knowledge makes it difficult to gain a physical understanding of the explosions so that we can model possible evolution, which compromises their use as distance indicators. It also means we do not fully understand the timescale over which SNe Ia turn on, adding a large uncertainty to our understanding of the chemical evolution of galaxies.

Unveiling the progenitor scenario for SNe Ia is difficult because white dwarfs (WDs) can, theoretically, reach their fatal Chandrasekhar mass in many ways, and disentangling which is the correct one (if there is just one), is challenging from an observational point of view. Nonetheless, there are two basic families of models leading to an SN Ia, the single-degenerate model (SD) and the double-degenerate model (DD). In the SD scenario, a WD accretes mass from a hydrogen-rich companion star before reaching a mass close to the Chandrasekhar mass and going off as an SN. In the DD scenario, two WDs merge, with the more-massive WD being thought to tidally disrupt and accrete the lower-mass WD (see, e.g., Maoz et al. 2014, and references therein).

Observations can potentially discriminate between the progenitor models of SNe Ia. For example, in all scenarios with mass transfer from a companion, a significant amount of circumstellar gas is expected (see, e.g., Branch et al. 1995), and therefore a shock is bound to form when the SN ejecta are expelled. The situation would then be very similar to circumstellar interaction in core-collapse SNe, where the interaction of the blast wave from the SN with its circumstellar medium results in strong radio and X-ray emission (Chevalier 1982b). On the other hand, the DD scenario will not give rise to any circumstellar medium close to the progenitor system, and hence essentially no radio emission is expected.

Radio and X-ray observations of SN 2011fe have provided the most sensitive constraints on possible circumstellar material (Chomiuk et al. 2012; Margutti et al. 2012) around a normal SN Ia. The claimed limits on mass-loss rate from the progenitor system are $\dot{M}= 6\times 10^{-10}$ M_☉ yr⁻¹ and $\dot{M}= 2\times 10^{-9}$ M_☉ yr⁻¹ from radio (Chomiuk et al. 2012) and X-rays (Margutti et al. 2012), respectively, assuming a wind velocity of 100 km s⁻¹. Radio (e.g., Panagia et al. 2006; Hancock et al. 2011) and X-ray (e.g., Hughes et al. 2007; Russell & Immler 2012) observations of other, more distant SNe Ia, have resulted in less constraining upper limits on wind density. The non-detections of radio and X-ray emission from SNe Ia have added to a growing consensus that a large fraction of SNe Ia may not be the result of SD scenarios (e.g., Maoz et al. 2014).

Despite the non-detection of radio and X-ray emission, there is evidence of possible circumstellar material in the form of time-varying absorption features in the optical Na i D line for a few SNe Ia (Patat et al. 2007; Simon et al. 2009; Dilday et al. 2012), supposed to arise in circumstellar shells. The exact location of the absorbing gas is still debated (e.g., Chugai 2008; Soker 2014), and probably varies from case to case. The number of SNe Ia showing indications of circumstellar shells could be significant, although the uncertainty is still large ((18 ± 11)%; Sternberg et al. 2014). Just as with the radio and X-rays, no optical circumstellar emission lines from normal SNe Ia have yet been detected (e.g., Lundqvist et al. 2013), although there are a few cases with strong emission (see, e.g., Maoz et al. 2014, for an overview). Those SNe Ia with strong circumstellar interaction constitute a very small fraction of all SNe Ia, probably only ∼1% (Chugai et al. 2004).

Recently, Fossey et al. (2014) serendipitously discovered SN 2014J in the nearby galaxy M82 (D = 3.5 Mpc). Cao et al. (2014) classified SN 2014J as an SN Ia, which makes it the closest SN Ia since SN 1986G in Cen A, almost three decades ago. The SN exploded between UT 2014 January 14.56 and 15.57 according to the imaging obtained by Itagaki et al. (2014), ¹⁵ and its J2000.0 coordinates are R.A. = 09:55:42.121, decl. = +69:40:25.88 (Smith et al. 2014). For further discussion on the discovery and early rise of the optical/IR emission, we refer to Goobar et al. (2014) and Zheng et al. (2014). The vicinity of SN 2014J makes it a unique case for probing its prompt radio emission, and thus constrain its progenitor system.

We observed SN 2014J with the electronic Multi Element Radio Interferometric Network (eMERLIN) at 1.55 and 6.17 GHz, and with the electronic European Very Long Baseline Interferometry Network (EVN) at a frequency of 1.66 GHz. We show in Table 1 the summary for our observations, along with the radio data obtained by others, and in Figure 1 the main results from our eMERLIN and EVN observations.

Table 1. Log of Radio Observations

Starting	T	tint	Array	ν	Sν	Lν, 23	$\dot{M}_{-9}$
(UT)	(day)	(hr)		(GHz)	(μJy)
Jan 23.2	8.2	...	JVLA	5.50	12.0	1.77	0.70 (4.2)
Jan 24.4	9.4	...	JVLA	22.0	24.0	3.51	3.7 (22)
Jan 28.8	13.8	13.6	eMERLIN	1.55	37.2	5.46	1.15 (7.0)
Jan 29.5	14.5	14.0	eMERLIN	6.17	40.8	5.97	3.6 (22)
Feb 4.0	20.0	11.0	eEVN	1.66	32.4	4.74	1.69 (10)
Feb 19.1	35.0	10.0	eEVN	1.66	28.5	4.17	2.9 (16)

Notes. The columns starting from left to right are as follows: starting observing time, UT; mean observing epoch (in days since explosion, assumed to be on Jan 15.0); integration time, in hr; facility; central frequency in GHz; 3σ flux density upper limits, in μJy; the corresponding 3σ spectral luminosity, assuming a distance of 3.5 Mpc to M82, in units of 10²³ erg s⁻¹ Hz⁻¹; inferred 3σ upper limit to the mass-loss rate in units of 10⁻⁹ M_☉ yr⁻¹, for an assumed wind velocity of 100 km s⁻¹. (The values for $\dot{M}_{-9}$ are for _B = 0.1 and, in parenthesis, for _B = 0.01, and have been calculated using our model described in Section 4.) The Jansky Very Large Array (JVLA) data are taken from Chandler & Marvil (2014), while the rest of the data are from this work.

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2.1. eMERLIN Observations

2.2. eEVN Observations

The radio and X-ray non-detections of SNe Ia, in conjunction with indications of circumstellar shells around some SNe Ia (see Section 1), is a conundrum that yet has to find a solution. The nearby northern hemisphere SNe 2011fe and 2014J offer a possibility of using the most sensitive radio facilities present to probe circumstellar emission. In particular, we now interpret the upper limits on radio emission from SN 2014J in Section 2 within the framework of circumstellar interaction. Indeed, when the SN shockwave ploughs through the circumstellar gas, a high-energy-density shell forms. Within this shell, electrons are accelerated to relativistic speeds and significant magnetic fields are generated, especially if the circumstellar gas is pre-ionized. For the low wind densities discussed in this paper, pre-ionization is likely to occur (Cumming et al. 1996). The relativistic electrons radiate synchrotron (radio) emission (e.g., Chevalier 1982b).

A proper modeling of the radio emission from SNe requires, in principle, taking into account Coulomb, synchrotron, and (inverse) Compton losses of the relativistic electrons. However, since we only have upper limits for the radio emission from SN 2014J, we will discuss the radio emission from SNe Ia within a scenario of Type Ib/c SNe (see, e.g., Chevalier & Fransson 2006), neglecting energy losses for the relativistic electrons (see Fransson & Björnsson 1998; Martí-Vidal et al. 2011, for a more general treatment). The spectrum of the radio emission from those SNe follows the "synchrotron self-absorption" (SSA) form, i.e., a rising power law with ν^5/2 (low-frequency, optically thick regime) and a declining power law, ν^α (high frequency, optically thin regime), where α is assumed to be constant. For most well studied SNe, α ≈ −1 (Chevalier & Fransson 2006). We assume that electrons are accelerated to relativistic energies, with a power law distribution, dN/dE = N₀ E^−p ; where E = γm_ec² is the energy of the electrons and γ is the Lorentz factor. For synchrotron emission, α = (p − 1)/2, which indicates that p ≈ 3 should be used.

Here, we study both the case of a circumstellar structure created by a wind, as well as the case with constant density circumstellar gas. For the wind case, we make the standard assumption that the SN progenitor has been losing matter at a constant rate, $\dot{M}$, so that the circumstellar density has a radial profile: $\rho (r) = n_{\rm CSM}(r)\mu = \dot{M}/(4\pi r^2 v_w)$, where v_w is the wind velocity, r is the radial distance from the star, n(r) is the particle density, and μ is the mean atomic weight of the circumstellar matter.

To calculate the shock expansion, we use the thin-shell approximation (Chevalier 1982b), with the extensions of Truelove & McKee (1999). We assume that the innermost ejecta has a density slope of ρ_{ej, inner}∝r^−δ, which at some velocity of the ejecta rolls over to a steeper density profile, ρ_{ej, outer}∝r⁻ⁿ (n > δ). We assume δ = −2, motivated by the explosion models of Fink et al. (2014), and use n = 10.2, which is a good approximation to the outer density profile of an SN that stems from a radiative star (Matzner & McKee 1999). Assuming an ejecta mass of 1.4 M_☉ and a kinetic energy of the explosion of 10⁵¹ erg, the break in the power-law index in our model occurs at ≈1.25 × 10⁴ km s⁻¹, which agrees with the angle-averaged results of Fink et al. (2014). The SN expansion can be well approximated by a power law, r_s ∝t^m , so that the shock speed, v_s = m  r_s /t (Chevalier 1982b). Here, m = (n − 3)/(n − s), and s is the density slope of the circumstellar gas, which for the steady wind case is s = 2. (In Section 4.1.2 we also discuss the case s = 0.) The shock speed at 10 days in this model is v_s ≈ 8.3 × 10⁴ km s⁻¹ for $\dot{M}= 1 \times 10^{-9}$ M_☉ yr⁻¹ and v_w = 100 km s⁻¹; v_s and r_s both scale as $(\dot{M}/v_w)^{-1/(n-s)}$.

For any sensible pre-SN wind speed, the SN shock is strong. Assuming a polytropic gas with γ = 5/3, the compression of the gas across the shock is η = 4, and the post-shock thermal energy density is $u_{\rm th} = ({9}/{8})\rho v_s^2$, where ρ is the pre-shock density. Following Chevalier & Fransson (2006), we denote _B = u_B/u_th, where u_B = B²/(8π) is the (post-shock) magnetic energy density; and _rel = u_rel/u_th, where u_rel is the energy density of the relativistic particles, assumed for simplicity to be electrons.

We assume that the power-law index of the relativistic electron population stays constant with time at p = 3, although we have also studied cases with p = 2.5 (see Section 4.2).

The most uncertain parameters refer to the microphysics of the shocked gas, namely _rel and, to a greater extent, _B. Indeed, it seems that _rel ∼ 0.1 with some small dispersion around this value (Chevalier & Fransson 2006), whereas _B appears to vary more among SNe, and is hence largely unknown. Therefore, we fix _rel = 0.1, and take _B as a free parameter. We can easily find N₀ by integrating the relativistic electron distribution between E_min = γ_min  m_ec² and infinity, which yields $N_0 = (p-2) \epsilon _{\rm rel} u_{\rm th} E_{\rm min}^{p-2}$.

We estimate the minimum Lorentz factor of the relativistic electrons, γ_min, assuming that all post-shock electrons go into the power-law distribution with energy index p (cf. Chevalier & Fransson 2006). This means that _rel u_th ≈ ηn_eE_min[(p − 1)/(p − 2)] (see also Chomiuk et al. 2012). Here, n_e is the electron density of the pre-shocked gas when it is fully ionized. We assume a mix of H and He with an abundance ratio 10:1, which together with _rel = 0.1 and η = 4 means that γ_min ≈ 1.64[v_s /(70, 000 km s⁻¹)]². Following Chevalier (1998), we add the constraint that γ_min ⩾ 1.

To calculate the synchrotron spectrum, we follow the method by Björnsson & Lundqvist (2014), i.e., we use the observational evidence that the brightness temperature, T_bright, is expected to be somewhat below 10¹¹ K (cf. Readhead 1994; Björnsson & Lundqvist 2014). While we defer a more complete discussion about this to a future paper (C.-I. Björnsson, in preparation), we have chosen a likely value of T_bright = 5 × 10¹⁰ K, which should be correct to within a factor of ∼2. The intensity at the frequency of the peak of the synchrotron spectrum, ν_peak, is then defined as $I_{{\nu }_{\rm peak}} \equiv 2kT_{\rm bright} ({\nu }_{\rm peak}/c)^2$, whereas the intensity at any frequency is I_ν = S_ν[1 − exp(− τ_ν)]. Here S_ν∝ν^5/2 is the source function and τ_ν the synchrotron optical depth. The latter is just τ_ν = κ_νΔs, where Δs is the path length through the emitting region along the line of sight, and κ_ν = ϰ(p)N₀ B^{(p + 2)/2}ν^{−(p + 4)/2}. Like Chevalier (1998), we make the simplification that B sin(θ) ≈ B, where θ is the particle pitch angle. The constant ϰ(p) can be found in, e.g., Rybicki & Lightman (1979). The path length Δs depends on the thickness of the synchrotron emitting region, Δr. At the center, ξ_h ≡ Δs(h)/(2Δr) = 1 (assuming the SN ejecta to be transparent to radio emission), but can become significantly larger than unity toward the limb. h is the normalized impact parameter, so that 0 ⩽ h ⩽ 1. We assume constant properties of the plasma within Δr. For s = 2, we have assumed a thickness of Δr/r_s = 0.2, which corresponds to that of the shocked circumstellar gas for n ∼ 10 and s = 2 in the similarity solutions of Chevalier (1982a), namely, Δr/r_s ≃ 0.19 ± 0.01 for 9 ⩽ n ⩽ 12. For s = 0, the similarity solutions give Δr/r_s ≃ 0.116 ± 0.008 for the same range in n, and we have chosen Δr/r_s ≈ 0.12 for this s-value. ξ_h is therefore just due to the geometrical increase of the path length as h increases.

For convenience, we also introduce, in addition to ν_peak, the frequency ν_abs, defined as $\tau _{\nu _{\rm abs}} = 1$. In general, τ_ν = (ν/ν_abs)^{−(p + 4)/2}. For h = 0, we denote ν_abs = ν_{abs, 0}, $\tau _{\nu _{\rm abs}} = \tau _{\nu _{{\rm abs},0}}$ and $\tau _{\nu } = \tau _{\nu _{0}}$. We can then derive the intensity for any impact parameter as

$$\begin{equation} I_{\nu }(h) = \frac{2kT_{\rm bright}}{c^2 f\left(\frac{{\nu }_{\rm peak}}{{\nu }_{\rm abs}}\right)} \frac{\nu ^{5/2}}{\nu _{{\rm abs},0}^{1/2}} [ 1-{\rm exp} (-\xi _h \tau _{\nu _{0}}) ], \end{equation} \tag{ 1 }$$

where f(x) depends on p such that

$$\begin{equation} f (x) = x^{1/2} [1 - {\rm exp} (- x^{-(p+4)/2})] \end{equation} \tag{ 2 }$$

(see also Björnsson & Lundqvist 2014). For p = 3(2.5), ν_peak/ν_abs ≈ 1.137(1.235) and f(x) ≈ 0.503(0.440). To obtain the luminosity, one integrates over h, so that $L_{\nu } = 8 \pi ^2 r_{s}^{2} \int _0^1 I_{\nu }(h) h dh$. The longer path length toward the limb makes L_ν larger for the optically thin part of the spectrum than just assuming $L_{\nu } = L_{{\nu },0} = 4 \pi ^2 r_{s}^{2} I_{\nu }(0)$. For p = 3, the factor ϑ_ν ≡ L_ν/L_{ν, 0} in the optically thin part is a weak function of Δr, being ≈1.81(1.63) for Δr/r_s = 0.1(0.2). For the optically thick part, ϑ_ν = 1. This makes the observed spectrum peak at a somewhat higher frequency than ν_peak.

4.1. Modeling the Data for SN 2014J

The radio emission from the SN is subject to both SSA and possible external free–free absorption in the ambient medium. While we include SSA in our model, we do not include free–free absorption since, as we show below, it is negligible. In previous analyses of SNe Ia (Panagia et al. 2006; Hancock et al. 2011), free–free absorption was assumed to be the most important factor when deriving wind densities. For SNe 2011fe and 2014J, X-ray non-detections (Margutti et al. 2012, 2014) have put limits on $\dot{M}/v_w$ of order $10^{-9} {\,{M_{\odot } \,\rm yr^{-1}}}$ for v_w = 100 km s⁻¹. From Equation (6) in Lundqvist & Fransson (1988) it follows that the free–free optical depth, τ_ff, for a fully ionized wind at 10⁴ K and moving at v_w = 100 km s⁻¹, is $\tau _{\rm ff} \sim 10^{-8} \lambda ^{2}\ (\dot{M}/10^{-9} {\,{M_{\odot } \,\rm yr^{-1}}})^2 (r_s/10^{15}\ {\rm cm})^{-3}$, where λ is in cm. For such a low wind density, the shock radius is ∼10¹⁵ cm already at 2 days, which means that $\tau _{\rm ff} \sim 3\times 10^{-7}\ (\dot{M}/10^{-9} {\,{M_{\odot } \,\rm yr^{-1}}})^2$ at 5.5 GHz at such an early epoch. Free-free absorption is thus insignificant and can be dismissed from our analysis. We also note that Horesh et al. (2011) used a similar argument to dismiss free–free absorption in their analysis of radio emission from SN 2011fe. In what follows, we therefore only consider SSA.

4.1.1. The Wind Case, s = 2

We now compare the radio data for SN 2014J in Section 2 with the predictions of the model presented in Section 3. If the SN happens in an SD scenario, the accreting WD is expected to have lost some of the accreted material from the donor star through a wind. This sets up a ρ∝r⁻² circumstellar structure (see Section 3).

As we show below, for the epochs of the radio observations, SN 2014J was clearly in its optically thin phase. This simplifies the expressions above, so that the luminosity becomes

$$\begin{equation} L_{\nu, {\rm thin}}=\frac{8 \pi ^{2} kT_{\rm bright} \vartheta _{\nu } r_s^2}{c^2 f\left(\frac{{\nu }_{\rm peak}}{{\nu }_{\rm abs}}\right)} \nu _{{\rm abs},0}^{(p+3)/2} \nu ^{-(p-1)/2}, \end{equation} \tag{ 3 }$$

where

$$\begin{equation} \nu _{{\rm abs},0}= (2 \Delta r \varkappa (p) N_0 B^{(p+2)/2})^{2/(p+4)}. \end{equation} \tag{ 4 }$$

From this, together with expressions in Section 3 and assuming p = 3, one gets

$$\begin{equation} L_{\nu, {\rm thin}} \propto T_{\rm bright} \epsilon _{\rm rel}^{1.71} \epsilon _{\rm B}^{1.07} (\dot{M} / v_w)^{1.37} t^{-1.55}, \end{equation} \tag{ 5 }$$

if γ_min is not fixed. If it is fixed

$$\begin{equation} L_{\nu, {\rm thin}} \propto T_{\rm bright} \epsilon _{\rm rel}^{0.86} \epsilon _{\rm B}^{1.07} (\dot{M} / v_w)^{1.27} t^{-1.35}. \end{equation} \tag{ 6 }$$

At early epochs, when the shock velocity is high, γ_min is always larger than unity, and decreases as time goes on. Therefore, Equation (5) applies. At later epochs, when the shock velocity is such that it would formally imply γ_min ⩽ 1, our constraint on γ_min takes effect, and Equation (6) applies. As stated in Section 3, we fixed _rel to 0.1 and allowed _B to vary. Equations (5) and (6) can be used to scale _rel even if only _B is varied.

In Figure 2 we show models for _B = 0.01 and 0.1. An almost perfect overlap between the modeled light curves occurs for the combination $\dot{M} = 7.0\times 10^{-10}{\,{M_{\odot } \,\rm yr^{-1}}}$ and _B = 0.1, and $\dot{M} = 4.2\times 10^{-9}{\,{M_{\odot } \,\rm yr^{-1}}}$ and _B = 0.01. Only at very early epochs (≲ 3 days after explosion) does SSA play a role for the lowest frequencies. The overlap is not surprising, since Equation (5) shows that $\dot{M}/v_{w} \propto \epsilon _{\rm B}^{-0.78}$ for fixed luminosity at early epochs in the optically thin part of our model. We note that Chomiuk et al. (2012) obtain a slightly different power-law index, −0.7, in their model. For all models in Figure 2, γ_min > 1 for the time span shown. This means that Equation (5) describes all light curves well, except for the lowest frequencies at t ≲ 3 days.

The values of $\dot{M}/v_{w}$ for SN 2014J in Figure 2 are chosen so that the 5.50 GHz light curves go through the JVLA 3σ upper limit on day 8.2. The light curves for other frequencies lie below their corresponding upper limits. The second most constraining limit is from our 1.55 GHz eMERLIN observation on day 13.8, yielding $\dot{M} \lesssim 1.15 (7.0) \times 10^{-9}{\,{M_{\odot } \,\rm yr^{-1}}}$ for _B = 0.1(0.01) and v_w = 100 km s⁻¹. We show in Table 1 the upper limits for all data points.

Figure 2 also includes the most constraining upper limit for SN 2011fe (Chomiuk et al. 2012), together with a 5.9 GHz light curve using $\dot{M} = 5.0 \times 10^{-10}{\,{M_{\odot } \,\rm yr^{-1}}}$, v_w = 100 km s⁻¹ and _B = 0.1. The limit on the mass-loss rate is somewhat below that of Chomiuk et al. (2012), who obtained $6.0 \times 10^{-10}{\,{M_{\odot } \,\rm yr^{-1}}}\ (v_w / 100 {\,km $^{-1}})$. The difference in those values probably stems from the difference in shell thickness of the emitting region, where we have adopted Δr/r_s = 0.2 versus Δr/r_s = 0.1 (Chomiuk et al. 2012), and our fixed T_bright. In any case, the difference in the upper limit on $\dot{M}/v_{w}$ is much smaller than that due to the uncertainty in _B.

In principle, radio non-detections could also be due to SSA during the observed epochs. In this case, the observed frequency ν_obs < ν_abs and from Equation (1) we find that the observed flux at ν_obs is $\propto r_s^2 (v_{\rm obs}^5/v_{\rm abs})^{1/2}$. The combination of $r_s^2 v_{\rm abs}^{-1/2}$ is a weak function of $\dot{M}/v_{w}$, and to make SSA important for the observations discussed here would require values of $\dot{M}/v_{w}$ much larger than those at which free–free absorption becomes important. We can therefore fully dismiss SSA as a cause for the radio non-detections of SN 2014J.

4.1.2. The Constant Density Case, s = 0

If the progenitor of SN 2014J followed the double-degenerate channel, then the exploding WD is expected to be surrounded by the interstellar medium (ISM), which has a constant density. Chomiuk et al. (2012) discussed this scenario for SN 2011fe, and obtained a limit for the density of ≲ 6  cm⁻³ (_B = 0.1).

The general behavior of the radio light curves for the constant density ISM case is different from the n_CSM∝r⁻² wind case in Section 4.1.1. For the constant density case and p = 3, the radio luminosity increases with time (see also Chomiuk et al. 2012) according to

$$\begin{equation} L_{\nu, {\rm thin}} \propto T_{\rm bright} \epsilon _{\rm rel}^{1.71} \epsilon _{\rm B}^{1.07} n_{\rm ISM}^{1.10} t^{0.38}, \end{equation} \tag{ 7 }$$

if γ_min is not fixed. If it is fixed

$$\begin{equation} L_{\nu, {\rm thin}} \propto T_{\rm bright} \epsilon _{\rm rel}^{0.86} \epsilon _{\rm B}^{1.07} n_{\rm ISM}^{1.27} t^{0.88}. \end{equation} \tag{ 8 }$$

Here, we substituted n_CSM with n_ISM to highlight the likely origin of the gas in the s = 0 case. Figure 3 shows models with densities n_ISM = 1.33  cm⁻³ (_B = 0.1) and n_ISM = 12  cm⁻³ (_B = 0.01). Scaling according to $n_{\rm ISM} \propto \epsilon _{\rm B}^{-0.97}$ (cf. Equation (7)) makes the light curves for these models overlap fully, except for t ≳ 25 days, when the condition γ_min ⩾ 1 becomes important for the n_ISM = 12  cm⁻³ model. The most constraining data are our eEVN 1.66 GHz data, stacked together, and the model parameters were chosen for the modeled radio luminosity to match those data. However, due to the γ_min ⩾ 1 constraint, the 35 day data alone are almost as constraining as the stacked data.

In Figure 3, we also show the stacked 5.9 GHz data for SN 2011fe (Chomiuk et al. 2012), together with a model characterized by n_ISM = 7.0  cm⁻³ and _B = 0.1. We are thus close to Chomiuk et al. (2012) regarding the limit on n_ISM for SN 2011fe. For s = 0, we used Δr/r_s = 0.12, which is close to the value 0.1 used by Chomiuk et al. (2012). The limit on n_ISM we find for SN 2014J is ≈5.3 times lower than for SN 2011fe, and is therefore clearly the lowest limit on density for the constant density case in any SN Ia.

As for the s = 2 case, SSA is unimportant for the s = 0 case. To be efficient enough to mute the radio emission to be consistent with the observed upper limit, n_ISM would have to be ≳ 10¹⁴  cm⁻³ (for _B ⩽ 0.1), which is fully ruled out from X-ray limits, as well as the normal optical behavior of the SN.

4.2. Sensitivity of Results to Parameters

5.1. The Possible Progenitors of SN 2014J

5.2. Broader Picture

5.3. Future Outlook for Radio Observations of SNe Ia

We report deep eEVN and eMERLIN radio observations of the Type Ia SN 2014J in the nearby galaxy M82, along with a detailed modeling of its radio emission. Our observations result in non-detections of the radio emission from SN 2014J. Yet, radio data and modeling allow us to place a tight constraint on the mass-loss rate from the progenitor system of SN 2014J. Namely, if the exploding WD was surrounded by a wind with a density profile ρ∝r⁻², as expected for an SD scenario, then our upper limit to the mass-loss rate is $\dot{M} \lesssim 7.0\times 10^{-10} {\,{M_{\odot } \,\rm yr^{-1}}}$, for a wind speed of 100 km s⁻¹.

If, on the contrary, the circumstellar gas has a constant density, as expected to be the case for the DD scenario (but also in a small region of the parameter space of SD scenarios), then our modeling yields an upper limit on the gas density, such that n_ISM ≲ 1.3  cm⁻³.

Our stringent upper limits to the circumstellar density around SN 2014J allow us to exclude completely symbiotic systems and the majority of the parameter space associated with stable nuclear burning WDs, as viable progenitor systems for SN 2014J. For the case of recurrent novae with main sequence or subgiant donors, we cannot rule them out completely, yet most of their parameter space is also excluded by our observations for the standard assumption of _B = 0.1, where _B is the ratio of magnetic energy density to post-shock thermal energy density.

We have also reassessed the radio limits on wind density for SN 2011fe, and for _B = 0.1 we obtain $\dot{M} \lesssim 5.0\times 10^{-10} {\,{M_{\odot } \,\rm yr^{-1}}}$ (for a wind speed of 100 km s⁻¹) and n_ISM ≲ 7.0  cm⁻³. These limits are close to those calculated by Chomiuk et al. (2012). Our limit on $\dot{M}/v_w$ for SN 2014J is thus similar to that for SN 2011fe, whereas for the constant density case we obtain a much lower limit than for SN 2011fe, and hence the lowest limit for a constant density ambient around an SN Ia.

The combined radio limits on circumstellar gas around SNe 2011fe and 2014J add to evidence, mainly from non-detections of X-rays from SN 2011fe and 2014J (Margutti et al. 2012, 2014) and no detection of Hα in the nebular phase of SN 2011fe (Shappee 2013b), that SNe Ia are very likely to stem from the DD scenario, rather than SD scenarios.

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Finally, we highlight future observations with the SKA. When fully completed, the sensitivity of the SKA will yield limits on circumstellar gas for future SNe Ia significantly better than the limits reported here for the nearby SNe 2011fe and 2014J, and up to distances beyond the Virgo cluster. With such upper limits (or detection), SKA is likely to be fully conclusive regarding the origin of the progenitor systems of SNe Ia.

We are grateful to Carles Badenes for useful comments on the manuscript, to Rubén Herrero-Illana for Pythonic advice to produce Figure 5, and to Vicent Peris and Oscar Brevià from the Observatorio Astronómico de Aras (Valencia, Spain) for the optical image of M82 used in Figure 1. We acknowledge the eMERLIN and EVN programme committees, and the directors of the EVN stations, for supporting the radio observations of SN 2014J. The European VLBI Network (EVN) is a joint facility of European, Chinese, South African, and other radio astronomy institutes funded by their national research councils. The electronic Multi-Element Radio Linked Interferometer Network (eMERLIN) is the UK's facility for high resolution radio astronomy observations, operated by The University of Manchester for the Science and Technology Facilities Council (STFC). The research leading to these results has received funding from the European Commission Seventh Framework Programme (FP/2007-2013) under grant agreement No. 283393 (RadioNet3). A.A., J.C.G., J.M.M., M.A.P.T., E.R., and I.M.V. acknowledge support from the Spanish MICINN through grants AYA2012-38491-C02-01 and AYA2012-38491-C02-02. P.L. acknowledges support from the Swedish Research Council. The research leading to these results has received funding from the European Commission Seventh Framework Programme (FP/2007-2013) under grant agreement No. 283393 (RadioNet3).

Facilities: EVN - European VLBI Network, MERLIN - Multi-Element Radio-Linked Interferometer Network Array