Multiwavelength Investigation of Pulsar Wind Nebula DA 495 with HAWC, VERITAS, and NuSTAR

DA 495 (G65.7+1.2) was discovered by the Dominion Astrophysical (DA) survey (Galt & Kennedy 1968). It was identified as a point source in the DA survey due to the coarse angular resolution but there was insufficient signal to determine any spectral information at the time. Follow-up as part of a SNR search, performed with the National Radio Astronomy Observatory and the Vermilion River observatory, found that DA 495 had an extended structure and a nonthermal spectrum suggestive of a Crab-like SNR (Willis 1973, and erratum). Further confirmation of the center-filled SNR hypothesis came with additional Dominion Radio Astrophysical Observatory observations (Landecker & Caswell 1983). We now know these center-filled SNRs to be PWNe, and DA 495 has been studied as such since then; however, an associated pulsation has not been found. More recent results from observations in radio and other wavelengths are detailed in the next sections.

The most recent radio analysis of DA 495 is reported in Kothes et al. (2008), hereafter referred to as K08. The maps generated by K08 from observations done with the Canadian Galactic Plane Survey (CGPS) in 408 and 1420 MHz (Taylor et al. 2003) and the Effelsberg Radio Telescope in 4850 and 10,550 MHz show an approximately circular diffuse source of about 25' in diameter. That size is fairly consistent across the radio spectrum reported in K08 and is consistent with earlier observations, but the nebula is not present in IRAS 60 μm observations. Fractional polarization is about 25% at higher frequencies (2695, 4850, and 10,550 MHz) and is ordered in a way that indicates a central dipole magnetic field with a superimposed toroidal component. The magnetic field of the nebula, measured from a synchrotron cooling break at 1.3 GHz, was determined to be extremely high at B = 1.3 mG. Using H i absorption and kinematic measurements relating to galactic rotation, K08 also estimated a distance to the source of d = 1.0 ± 0.4 kpc. We adopt the radio fluxes reported in K08, which have compact sources and the foreground H ii region removed for use with our combined spectrum described in Section 5. A map of radio emission at 1.4 GHz from the CGPS is shown as contours in Figure 1.

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DA 495 was first detected in the X-ray band in 2004 March using archival ROSAT and ASCA data (Arzoumanian et al. 2004). ROSAT source 1WGA J1952.2+2925, a faint, compact X-ray source, was assumed to be associated with the surrounding radio nebula. The X-ray flux was found to be nonvariable, nonthermal, and extended, leading to its identification as emission from a wind nebula. DA 495 was followed up with Chandra in 2007 (Arzoumanian et al. 2008), which was able to resolve a central point source inside the ∼40'' diameter (0.2 parsec at a distance of 1 kpc) X-ray nebula. Chandra spectral analysis of the extended emission resulted in a photon index of Γ = 1.6 ± 0.3, allowing for confirmation of the extended region as a wind nebula. The central point source was found to have a purely thermal spectrum, a result confirmed by Karpova et al. (2015; hereafter K15). K15 jointly fit the archival XMM-Newton and Chandra data with an absorbed power-law-plus-blackbody model and an absorbed power-law-plus-neutron-star atmosphere model (NSMAX) (Mori & Ho 2007; Ho et al. 2008) to characterize the point source emission. The blackbody fit, likely modeling emission from a polar cap, resulted in a temperature of T ≈ 0.22 keV, a radius of R ≈ 0.6 km, and a neutral hydrogen column density of N_H ≈ 2.6 × 10²¹ cm⁻². The NSMAX fit, modeling emission from the entire neutron star surface, resulted in a temperature of T ≈ 0.08 keV, a radius of R ≈ 10 km, and a neutral hydrogen column density of N_H ≈ 3.5 × 10²¹ cm⁻². Neither model was ruled out. No pulsations were detected in the Chandra data, but the pure thermal spectrum of the point source implies its likely identification as the neutron star powering the wind nebula. K15 further constrained the pulsation non-detection through timing analysis of the XMM-Newton data, setting the upper limit for a pulsed fraction of 40 percent in a range of ≥12.5 ms. K15 estimated the distance to the putative neutron star to be ∼2.4–3.3 kpc (depending on which spectral models were applied to the point source spectrum) using the N_H–D relation.

The first TeV gamma-ray detection in the region of DA 495 was reported by the High Altitude Water Cerenkov Telescope (HAWC), which detected 2HWC J1953+294 as a point source within 02 of the radio center reported by K08 (Abeysekara et al. 2017a). Given the 25' diameter radio size of DA 495 and HAWC's localization uncertainty of 01 it is certainly plausible that DA 495 is associated with 2HWC J1953+294 by position alone. Additional evidence that 2HWC J1953+294 is associated with DA 495 is the photon index at 7 TeV of Γ = 2.78 ± 0.15, which is consistent with other PWNe in this energy range (such as the Crab). The TeV gamma-ray detection of DA 495 was confirmed by follow-up observations with the Very Energetic Radiation Imaging Telescope Array System (VERITAS), which detected VER J1952+293 (Abeysekara et al. 2018). VER J1952+293 is centered within 005 of the radio center and has a extension defined by a 2D Gaussian with σ = 014 ± 002, which is consistent with the size of the radio nebula. The extent and position of the emission detected by VERITAS provides compelling evidence that the TeV and radio nebula are associated. The photon index measured by VERITAS at 1 TeV, Γ = 2.65 ± 0.49, also agrees well with the TeV PWN interpretation.

There is significant disagreement between the VERITAS and HAWC flux measurements of their respective sources. At 1 TeV, HAWC has a flux that is about 7× higher than that of VERITAS (Abeysekara et al. 2018) and is probably attributed to a nearby extended source and/or diffuse emission that contaminates the flux reported for 2HWC J1953+294. For this study, we adopt the VERITAS flux points and accept the HAWC measurement as an upper limit.

Observations of DA 495 were performed by NuSTAR on 2017 June 9 (ObsID 30362003002) for a total exposure of 60 ks. NuSTAR pointed at the X-ray centroid of the PWN at R.A. (J2000) = 19^h52^m1704, decl. (J2000) = 29°25'525 (Arzoumanian et al. 2008). The NuSTAR data was processed with nupipeline 0.4.6. DA 495 was observed by Chandra/ACIS-I on 2002 December 9 (ObsID 3900) for a total of 25 ks. Observations also took place with XMM-Newton EPIC/MOS (Full Frame Mode, medium filter setting) on 2007 April 21 for a total of about 50 ks. XMM-Newton data were processed using XMM SAS 1.2.

The NuSTAR, Chandra, and XMM-Newton data sets were used for both image analysis and spectroscopy. All three data sets underwent spectral extraction and joint fitting in XSPEC (v12.9.0), providing a more complete picture of the DA 495 X-ray emission.

4.1.1. NuSTAR Data Reduction

The NuSTAR data were processed and analyzed using the HEASOFT V6.21 software package, including NUSTARDAS 06December16 V1.7.1, with NuSTAR Calibration Database (CALDB) files from 2017 June 14.

NuSTAR background-subtracted images were obtained using the nuskybgd software (Wik et al. 2014). Nuskybgd maps and generates images for the entire background, allowing for specific background spectrum extraction given a source region. It takes into account stray light leaking through the aperture stop, focused cosmic X-ray background, instrumental background, and soft environmental neutrons from cosmic rays. When modeling the background with nuskybgd, three source-free regions were selected from each module, each region on a different detector chip. Although stray light contamination mainly manifested below 3 keV, the affected region, the top third of module A, was avoided when choosing background regions for modeling. Near-source background light curve analysis revealed no periods of high background or flaring.

After background subtraction three X-ray sources were visible in the NuSTAR field of view (Figure 2(a)). CIAO wavdetect detected the two point sources, S1 and S2, in the NuSTAR data from 3 to 10 keV, but did not detect them in the 10–20 keV band. DA 495 appears up to 20 keV. The archival Chandra data revealed counterpart point sources to S1 and S2: CXO J195205.6+292808 (S1) and CXO J195222.6+293005 (S2). These nonvariable low-energy point sources are extremely unlikely counterparts to the TeV gamma-ray emission and are excluded from further analysis.

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We extracted spectra from both the NuSTAR data and archival Chandra and XMM-Newton data and jointly fit the low- and high-energy spectra.

4.2.1. NuSTAR

4.2.2. Chandra

4.2.3. XMM-Newton

4.2.4. Fitting Results

When fitting, the NuSTAR spectrum was cut off above 20 keV where the background began to dominate. Thus we used a range of 3–20 keV for the NuSTAR spectrum, a range of 0.5–10 keV for Chandra and a range of 0.2–10 keV for XMM-Newton EPIC PN, MOS 1, and MOS 2. All error bars were calculated to 1σ significance. We fit the NuSTAR, Chandra, and XMM-Newton data jointly with an absorbed power-law model plus a blackbody component (tbabs∗(bbodyrad+powerlaw)) to account for the central pulsar (abundances from Wilms et al. 2000). This fit (Figure 3) produced a photon index of Γ = 2.0 ± 0.1 (typical for a PWN), a neutral hydrogen absorption of ${N}_{{\rm{H}}}={3.3}_{-0.8}^{+1.2}\times {10}^{21}$ cm⁻², a blackbody temperature of $T={0.18}_{-0.04}^{+0.05}$ keV, and a blackbody radius of $R={0.5}_{-0.3}^{+0.8}$ km (Table 1). All results are consistent with Karpova et al. (2015), who theorized that this blackbody radius and temperature model emission come from a hot polar cap. Absorbed X-ray flux was 2.4 ± 0.1 × 10⁻¹³ erg s⁻¹ cm⁻² in the 2–20 keV band and ${2.4}_{-0.2}^{+0}\times {10}^{-13}$ erg s⁻¹ cm⁻² in the 0.5–8 keV band, with ${L}_{2-20}={2.9}_{-0.3}^{+0.2}\times {10}^{31}$ erg s⁻¹ (unabsorbed) at a distance of 1 kpc. This fit resulted in a reduced chi-squared of 1.0 for 482 degrees of freedom, confirming that the spectrum is nonthermal and fits well to a single power-law model. There is no evidence of a spectral cutoff or break up to 20 keV.

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Table 1.  X-Ray Joint Fit Spectral Parameters

Model	NH, 1021 cm−2	T, eV	R, km	Γ	Power-law Norm, 10−5	${\chi }_{v}^{2}$	dof
BB	${3.3}_{-0.8}^{+1.2}$	${180}_{-40}^{+50}$	${0.5}_{+0.8}^{-0.3}$	2.0 ± 0.1	${6.7}_{-1.3}^{+1.4}$	1.0	482
NSMAX (B = 1012 G)	${4.0}_{-0.8}^{+0.2}$	28+6	${10}_{-9}^{+3}$	1.9 ± 0.1	${5.8}_{-0.7}^{+0.5}$	1.0	482

Note. All errors are given with 1σ confidence.

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The spectra were also fit with a neutron star atmosphere model as a replacement for the blackbody model (tbabs∗(NSMAX+powerlaw)), with B = 10¹² G to model thermal emission from the entire neutron star surface. This fit resulted in a photon index of Γ = 1.9 ± 0.1, a neutral hydrogen absorption of ${N}_{{\rm{H}}}={4.0}_{-0.8}^{+0.2}\times {10}^{21}$ cm⁻², an atmosphere temperature of T = 28⁺⁶ eV, and a radius of $R={10}_{-9}^{+3}$ km (Table 1). No lower limit was derived for the temperature, as the temperature fit to the lower limit value of the model. The gravitational redshift was frozen to a standard value of 0.3. The NSMAX fit is not shown in Figure 3, as it resulted in similar residuals and goodness of fit as the blackbody model.

For spectral energy distribution (SED) fitting, an annulus of r = 2'' centered on the central point source was subtracted from the Chandra data to avoid thermal emission from the putative neutron star, and photons below 2 keV (where the absorption component cut off) were omitted to account for neutral hydrogen absorption at low energies. Both the NS thermal emission and ISM absorption are negligible in the NuSTAR band (3–79 keV). Thus, the X-ray data (Chandra and NuSTAR only) fit in the SEDs contains only the power-law contribution.

4.2.5. Spatially Resolved Spectral Analysis with Chandra

4.2.6. X-Ray Flux Upper Limit

An X-ray flux upper limit for SED fitting was extracted by analyzing the extended emission within the radio nebula region r < 125 (Figure 5(b)). In a single electron population leptonic emission scenario one would expect diffuse X-ray emission coincident with the large radio nebula. XMM-Newton was the only X-ray telescope with a large enough field of view to be used to attempt a diffuse nebula upper limit. EPIC-MOS 2 was the sole camera used for analysis, as two CCD chips in MOS 1 were compromised, and EPIC PN was operated in the Small Window mode. XMM SAS 17.0.0 ESAS procedures were used for source and background spectrum generation. ESAS background subtraction procedures account for the quiescent particle background, but do not account for CXB emission, solar wind charge exchange (SWCX) background, soft proton background, or instrumental lines. Because instrumental lines and SWCX effects manifest below 2 keV, all energies below 2 keV were ignored. The light curve of the region was examined and no flares were found. The CXB was accounted for by fitting a second power law to the spectrum using the known photon index of Γ = 1.46 (Snowden et al. 2004). It was not possible to directly subtract a local background spectrum, as the radio nebula covers the entire XMM field of view and therefore no region was guaranteed to be source-free. All point-like sources, including DA 495, were masked using the XMM ESAS cheese command, and the extended emission spectrum was fit separately from the compact nebula. Because of the sharp (factor of ∼6) drop in flux between the compact X-ray nebula and surrounding background (Figure 5), it is unlikely that much of the measured X-ray flux from the radio emission region is from DA 495.

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When fitting the extended region, N_H was frozen to the joint fit value of 3.3 × 10²¹ cm⁻². Subtracting the relevant background components and ignoring the photons below 2 keV resulted in a flux (2–10 keV) of ${3.6}_{-0.5}^{+0.7}\times {10}^{-12}$ erg s⁻¹ cm⁻² when the DA 495 photon index was allowed to fit freely (very soft at Γ ≈ 5.6), and a flux of ≈4.7 × 10⁻¹² erg s⁻¹ cm⁻² when the photon index was frozen to the DA 495 best fit value of 2.0 (with the normalization of the CXB component frozen to the value found in the previous fit). The compact X-ray nebula upper limit flux was measured to be ≈1.6 × 10⁻¹³ erg s⁻¹ cm⁻² in the 2–10 keV band. The extended region flux (with all point-like sources masked) was added to the DA 495 compact flux value to gain a diffuse X-ray nebula upper limit of (3.3–4.9) × 10⁻¹² erg s⁻¹ cm⁻².

In a pure leptonic model, the radio to TeV emission generally consists of two components, namely, a low-energy component from synchrotron emission from primary electrons and a high-energy component from inverse Compton scattering of CMB photons by the same electrons. Given the large extents of the radio and TeV nebulae, the primary electron synchrotron photon density is much lower than the CMB. Thus, we do not expect a significant synchrotron–self-Compton contribution. The compact X-ray nebula originates from the neighborhood of the central pulsar, where the higher magnetic field and freshly accelerated electrons trigger an additional synchrotron component.

We can estimate the average magnetic field strength within the radio and TeV nebulae based on the spectral shapes. In a pure leptonic model, the radio and TeV emission should come from the same electron population. Because the radio spectrum is rising in a power-law shape while the TeV spectrum appears like a cutoff, we expect that the underlying electron spectrum has a spectral break at some maximal energy. To upscatter CMB photons to TeV energies, the primary electron cutoff should be ∼20 TeV. Generally speaking, this spectral break can have two origins, either the synchrotron cooling break or the intrinsic spectral cutoff. If the electron cutoff at ∼20 TeV results from synchrotron cooling, the synchrotron cooling timescale should be comparable to the source age,

$$\begin{eqnarray}\begin{array}{rcl}{t}_{\mathrm{cool}} & = & \displaystyle \frac{{m}_{e}{c}^{2}}{4/3c{\sigma }_{T}\gamma ({u}_{B}+{u}_{\mathrm{CMB}})}\\ & = & 0.976\,\mathrm{yr}\,{\gamma }^{-1}{\left({u}_{{\rm{B}}}+{u}_{\mathrm{CMB}}\right)}^{-1}\sim {t}_{\mathrm{age}}=20\,\mathrm{kyr}.\end{array}\end{eqnarray} \tag{ 1 }$$

Therefore, we can find that B ∼ 8 μG, consistent with the ISM magnetic field. This implies that the synchrotron component should peak at

$$\begin{eqnarray}&&h{\nu }_{c}=h\displaystyle \frac{3{eB}}{4\pi {m}_{e}c}{\gamma }^{2}\sim 1.74\times {10}^{-8}\,\mathrm{eV}\,B{\gamma }^{2}\sim 50\,\mathrm{eV},\end{eqnarray} \tag{ 2 }$$

which is in the ultraviolet.

The extent of the TeV nebula is likely determined by the diffusion of TeV electrons. The typical interstellar diffusion coefficient of ∼10 TeV electrons is ∼10²⁹ cm² s⁻¹. Considering the radiative cooling of these TeV electrons in the ISM, the diffusion radius is given by

$$\begin{eqnarray}&&{d}_{\mathrm{dif}}=2\sqrt{{{Dt}}_{\mathrm{cool}}}\sim 100\,\mathrm{pc}.\end{eqnarray} \tag{ 3 }$$

Apparently, the extent should be much larger than the observed angular extent assuming a distance of $1\,\mathrm{kpc}$. To trap the high-energy electrons within the observed ∼2 pc region, the diffusion coefficient should be D ∼ 4 × 10²⁵ cm² s⁻¹, much smaller than the ISM diffusion coefficient. Comparing to the Bohm diffusion coefficient, which is given by

$$\begin{eqnarray}&&{D}_{\mathrm{Bohm}}=\displaystyle \frac{1}{3}\displaystyle \frac{\gamma {{mc}}^{2}}{{eB}}c\sim 6.67\times {10}^{25}\,{\mathrm{cm}}^{2}\,{{\rm{s}}}^{-1}\,{\left(\displaystyle \frac{B}{5\mu {\rm{G}}}\right)}^{-1},\end{eqnarray} \tag{ 4 }$$

the observed slow diffusion then sets a lower limit on the average magnetic field, B ≳ 10 μG. Recent HAWC observation of the Geminga PWN suggested a similarly small diffusion coefficient within the Geminga nebular diffusion comparable to the local Bohm diffusion coefficient (Abeysekara et al. 2017b; López-Coto et al. 2018), which is consistent with our DA 495 fitting parameters.

Because the same electrons produce the radio emission through synchrotron and the TeV emission through inverse Compton scattering, their flux levels can be used to derive the ratio between the magnetic energy density and the target photon energy density, namely, u_B/u_CMB. Figure 6 shows our leptonic fitting results. We find that the magnetic field cannot be too high, or else there should exist a very bright X-ray nebula of similar size to the radio and TeV nebulae, which cannot be confirmed in our observations. Notice that our spectral fitting considers cooling effects, thus we observe a cooling break in the B = 15 μG case.

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The sharp cutoff at the edge of the central compact X-ray nebula implies a stronger magnetic field and newly accelerated electrons. Figure 6 shows a sample fitting with 50 μG magnetic field, although with merely X-ray data we cannot constrain both the magnetic field strength and the electron spectrum. The higher magnetic field strength should infer brighter radio emission within the compact X-ray nebula. Instead, K08 suggests an apparent radio hole in that region. This casts doubt on simple leptonic models.

In a pure hadronic model, the broadband emission of the extended nebula originates from proton–proton (pp) collisions, where the neutral pion decay makes the high-energy spectral component while the charged pion decay results in secondary electron–positron pairs, which then give rise to the low-energy spectral component through synchrotron radiation. The compact X-ray nebula is likely the primary electron synchrotron radiation, and its small extent is due to the fast cooling time of these electrons in a high magnetic field.

The observed radio and TeV spectra can put strong constraints on the underlying proton spectral distribution and magnetic field strength. The soft TeV spectral shape indicates the neutral pion decay spectral cutoff. This corresponds to a proton cutoff energy at (Kelner et al. 2006)

$$\begin{eqnarray}&&{E}_{p,\mathrm{cut}}=10\times {E}_{\gamma ,\mathrm{cut}}\sim 10\,\mathrm{TeV}.\end{eqnarray} \tag{ 5 }$$

Meanwhile, the rising radio spectrum marks the low-energy cutoff of the secondary electrons. These electrons, which come from the charged pion decay, should have a low-energy cutoff at ∼500 MeV (Kelner et al. 2006). This is because the cross section of the pp collisions cuts off when the nonthermal protons become nonrelativistic. Therefore, we can quickly derive from Equation (2) that the magnetic field strength should be

$$\begin{eqnarray}&&B\sim 0.5\,\mathrm{mG}\,\displaystyle \frac{{E}_{\mathrm{radio}}}{{10}^{-5}\,\mathrm{eV}}{\left(\displaystyle \frac{\gamma }{{10}^{3}}\right)}^{-2}.\end{eqnarray} \tag{ 6 }$$

Since the cross section for charged pions is about half of the neutral pion cross section, and part of the charged pion energy goes into neutrinos, the total electron power from pp collisions is typically about one-third of the total γ-ray power. We can see from the DA 495 spectrum that the TeV flux is ∼5 times higher than the radio flux, indicating that the secondary synchrotron emission is very efficient (Figure 7). Indeed, we can quickly estimate from Equation (1) that the cooling time of the secondary electrons in a ∼0.5 mG magnetic field should be ∼90 kyr for the low-energy cutoff electrons at ∼500 MeV, on the same order of the source age.

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Given the extent of the TeV nebula, we can estimate the diffusion coefficient inside. Because the protons that produce the TeV emission do not cool in the nebula, the diffusion coefficient is

$$\begin{eqnarray}&&D={\left(\displaystyle \frac{R}{2}\right)}^{2}/{t}_{\mathrm{age}}\sim 1.5\times {10}^{25}\,{\mathrm{cm}}^{2}\,{{\rm{s}}}^{-1},\end{eqnarray} \tag{ 7 }$$

which is a little smaller than that in the leptonic scenario. But here the magnetic field is two orders of magnitude higher than the typical ISM value, thus the diffusion coefficient is well above the Bohm limit.

While the radio and TeV nebulae are produced by pp collisions, the compact X-ray nebula near the central pulsar comes from the primary electrons coaccelerated with the protons. Due to the high magnetic field, these electrons will cool in a short time, which naturally explains the small spatial extent of the X-ray nebula. From Equation (2), the X-ray spectral shape implies that the primary electrons should be accelerated to ∼20 TeV, consistent with the maximal proton energy.

Certainly, without a detected pulsar the age estimation becomes more challenging and the displacement of the neutron star from the radio nebula's center, presuming a reasonable kick, is consistent with the currently estimated ages (Arzoumanian et al. 2004). There are several relationships for PWN that can place rough age constraints on the undetected pulsar based on observational parameters alone.

A study on the X-ray and γ-ray luminosity of PWN (Mattana et al. 2009) describes a relationship between ${L}_{1\mbox{--}30\mathrm{TeV}}/{L}_{2\mbox{--}10\mathrm{keV}}$ that can be used to find τ_c of the pulsar and $\dot{E}$. This is a characteristic relationship due to the different electron populations that generate the X-ray synchrotron and TeV inverse Compton emissions. Using the 1 kpc luminosity from the VERITAS spectral fit and the joint X-ray fit in this work: ${L}_{1\mbox{--}30\mathrm{TeV}}=(1.3\pm 0.2)\,\times {10}^{31}\,\mathrm{erg}\ {{\rm{s}}}^{-1}$ and ${L}_{2\mbox{--}10\mathrm{keV}}=(2.6\pm 0.2)\times {10}^{31}\,\mathrm{erg}\ {{\rm{s}}}^{-1}$ for a ratio of ${L}_{1\mbox{--}30\mathrm{TeV}}/{L}_{2\mbox{--}10\mathrm{keV}}=0.50\pm 0.05$. This results in a derived $\dot{E}=(2.6\pm 0.2)\times {10}^{37}\,\mathrm{erg}\ {{\rm{s}}}^{-1}$ and τ_c = 3.5 ± 0.5 kyr—dramatically different than the K08 estimated values of $\dot{E}\,\sim {10}^{35}\,\mathrm{erg}\ {{\rm{s}}}^{-1}$ and τ_c ∼ 20 kyr. The authors in Mattana et al. (2009) note that DA 495 is expected to be an exception with respect to their fits due to the strong magnetic field inferred from radio observations. Now that we have included the TeV emission, we can confirm their suspicions that DA 495 does not line up with the rest of the PWN population from an observational perspective. If the younger age estimate is to be believed, then we would expect a shell of angular size θ_SNR ≈ 31' (using Equation (1) from Romani et al. 2005). Without clear evidence for a radio shell at this size, we do not find this estimated age to be more reliable than the synchrotron cooling age that was estimated by K08.

Another population study of PWNe was performed by HESS (H.E.S.S. Collaboration et al. 2018), where we get yet another story about the age. For this study, TeV luminosity (${L}_{1\mbox{--}10\mathrm{TeV}}$) was used to find estimate the age, spindown luminosity, and size. Using the K08 age for DA 495, the expected ${L}_{1\mbox{--}10\mathrm{TeV}}\ (@20\ \mathrm{kyr}\ \mathrm{and}\ 1\ \mathrm{kpc})\ \mathrm{is}\sim 7\times {10}^{33}\,\mathrm{erg}\ {{\rm{s}}}^{-1}$. The measured ${L}_{1\mbox{--}10\mathrm{TeV}}\ (@1\ \mathrm{kpc})=(1.3\pm 0.2)\times {10}^{31}\,\mathrm{erg}\ {{\rm{s}}}^{-1}$ would put the age at >1 Myr when using the HESS flux. So, while the enhanced TeV emission is causing issues with understanding the SED models with a strong magnetic field, by these estimates the TeV emission is very low for the expected age, as the TeV luminosity is expected to decrease with the PWN age. There is some suggestion in K08 that DA 495 is at a larger distance (D ∼ 5 kpc). The measured luminosity, ${L}_{1\mbox{--}10\mathrm{TeV}}\ (@1\ \mathrm{kpc})=(3.3\pm 0.4)\times {10}^{31}\,\mathrm{erg}\ {{\rm{s}}}^{-1}$, still has significant disagreement with any assumed age. Without any other age estimates, the luminosity at 5 kpc would still put the age of DA 495 at >1 Myr. We are not suggesting that DA 495 is a remnant that old, but that it is an outlier from an observational standpoint and that standard PWN evolutionary scenarios do not apply.

These conflicting results on the age of the object reinforce our conclusions that DA 495 has not followed a typical evolutionary path. Time-dependent modeling of atypical evolutionary scenarios of the PWN including SNR evolution would be the most likely way to help resolve some of the questions regarding possible age.

The magnetic field strength estimated from the radio observations and the SED may be reconciled using a model including hadronic emission. This requires hadronic target material and the presence of a SNR shell. This is contrary to the K08 claim about the lack of a shell, though earlier interpretation of the radio emission from Velusamy et al. (1989) argued that DA 495 is a composite remnant with a thick shell due to slow supernova ejecta interacting with the reverse shock. The thick shell claim may still be considered a possibility if there is a superimposed toroidal magnetic field as K08 suggests. The magnetic field strength of the hadronic scenario may be possible with there only being a shock interacting with ejecta in a ring around the source. This, in part, could also explain the very low surface brightness of the nebula, because the radio and TeV emission would be coming only from the equatorial ring. However, the physical reason for why the toroidal magnetic field was entrapped is not clear. This interpretation still has issues within the context of other observations, as there is not an obvious, cooler, dense, target region as expected with most shock interactions.

A way to further investigate the existence of a shell-type remnant would be to examine the TeV morphology more closely. The radio hole is about 2' in diameter, and if the radio and TeV emission both share this feature, then this could indicate that the emission is from a shocked region where the magnetic field was compressed. The putative hadronic targets in this possible scenario would be slower ejecta from the supernova. If the TeV emission is smooth through the middle of the region, then that would indicate that the electrons from the central pulsar are likely responsible for the TeV emission. This argues against the hadronic scenario, as the radio and TeV emissions are no longer spatially linked. Currently, none of the gamma-ray instruments are sensitive enough to definitively make this distinction, but the Cerenkov Telescope Array, due to come online in the next few years, may have the enhanced angular resolution necessary (The CTA Consortium et al. 2018).

While both leptonic and hadronic scenarios can reasonably reproduce the multiwavelength observation (Figures 7 and 6), we can quickly note several key predictions from the spectral modeling. For a pure leptonic model, we expect a straight power law in the low-energy spectral component, from radio to optical. K08, however, suggested a spectral break in the radio component. Future observations at higher radio frequencies to the far-infrared may help to diagnose the significance of this spectral break. Additionally, the pure leptonic model predicts rather high optical to ultraviolet flux. Most interestingly, it implies that there exists a diffusive X-ray nebula that is of similar size as the radio and TeV nebulae. The total flux is comparable to the central compact X-ray nebula, but because it spreads out ∼2 pc, its surface brightness can be much lower than the central compact X-ray nebula, which is very hard to detect. The leptonic model requires an average magnetic field that is slightly higher than the typical ISM value, probably due to the interaction between pulsar wind and the ISM, or maybe the remnant magnetic field from the supernova. The total energy budget for the leptonic model is ∼10⁴⁸ erg, typical for PWNe.

The hadronic models lead to very different predictions. In a pure hadronic model, the radio spectral break is a natural result of the secondary pair synchrotron. Both the infrared to optical and Fermi γ-ray spectra should appear flat, assuming an underlying proton energy spectral index of ∼2.0. The pure hadronic model also suggests that there is no diffusive X-ray nebula, which is consistent with our observations. However, the hadronic model requires a very strong magnetic field pervading the radio and TeV nebulae, which is likely powered by the central pulsar, consistent with K08. Nonetheless, protons are generally unlikely to be accelerated at typical particle acceleration sites near the pulsar, such as the polar cap. Instead, they should be accelerated due to the interaction between the pulsar wind and the ISM, or magnetic reconnection in the highly magnetized nebula. Given the fitting parameters in Table 2, we can see that the total magnetic energy within the PWN is about 2.5 × 10⁴⁹ erg, indicating that the nebula is considerably magnetized. These features require that DA 495 be a very unusual PWN, one that has extremely high power and can extend its magnetic field to ∼2 pc into the ISM. Interestingly, the latter is supported by the radio polarization map by K08, where they find a dipole shape magnetic field morphology in the radio nebula.