抽象
本文研究了两颗贫氢超光超新星（SLSNe），iPTF15esb和iPTF13dcc，其光曲线（LCs）显示出与平滑上升和下降的显着偏差。iPTF15esb的液相色谱表现出两个峰和一个峰后平台，此外，iPTF15esb的晚期光谱显示出强而宽的Hα发射线。iPTF13dcc 的早期 LC 显示长时间凸起，然后是第二个峰。本文提出了一种涉及多个壳层/风的抛射-星周介质相互作用模型，并用它来解释iPTF15esb和iPTF13dcc的LCs。我们发现该模型再现的理论LC可以匹配iPTF15esb和iPTF13dcc的观测结果。根据我们的结果，我们推断祖先在SN爆炸之前经历了多次剧烈的质量损失过程。此外，我们发现我们推断的壳层密度的变化趋势与SN爆炸前的恒星质量损失历史所预测的一致。对其他凹凸不平的SLSNe/SNe的进一步研究将揭示它们的性质，并为其祖先的质量损失历史提供探测。
1. 引言
在过去的十年中，快速发展的非目标超新星（SN）勘测计划发现了一类新的不寻常的SNe，其峰值绝对星等为M_{峰}亮度大于−21等。这些非常明亮的SNe被称为超亮超新星（SLSNe）;Quimby 等人，2011 年;Gal-Yam 2012 年）。
看来，SLSNe可以简单地分为两个亚类，即I型和II型。SLSNe I 在缺乏氢吸收线的峰周围有光谱，它们的光曲线 （LC） 可能由对不稳定性 SNe 解释（Rakavy & Shaviv 1967;Heger & Woosley 2002;Heger et al. 2003），^{56}镍动力模型，磁星动力模型（Kasen & Bildsten 2010;伍斯利 2010;Inserra 等人，2013 年;Wang et al. 2015a， 2015b， 2016a;Dai 等人，2016 年;Liu 等人，2017 年;Yu et al. 2017），或喷射物-星周介质（CSM）相互作用模型（Chevalier & Irwin 2011;Chatzopoulos 等人，2012 年，2013 年;金茨堡和巴尔伯格 2012;Nicholl 等人，2014 年;Chen 等人，2015 年）。
另一方面，SLSNe II峰周围的光谱显示出较强的氢发射特征，并且几乎所有峰都显示出窄和中等宽度的Balmer发射线，与正常的SNe IIn相似。以前的研究（Smith & McCray 2007;Moriya 等人，2011 年，2013 年;Chatzopoulos et al. 2012， 2013）认为，SLSNe IIn的LCs可能由SN喷射物与致密、富氢和光学厚的CSM之间的相互作用提供动力。
然而，一些SLSNe I（例如，iPTF13ehe、iPTF15esb和iPTF16bad）的晚期光谱显示Hα发射线（Yan等人，2015,2017）使分类方案复杂化。 Yan等人（2015）估计，15%的SLSNe I可能具有这些光谱特征。在这些具有晚期Hα发射谱线的SLSNe I中，iPTF15esb以0.224的红移z爆炸，是最引人注目的。其后期光谱显示强而宽的Hα发射线，表明SN喷射物与SN祖细胞周围的富氢CSM壳层之间的相互作用。此外，其LC有两个峰，其光度彼此大致相等（L_{峰}≈ 4 × 10^{43}ERG S^{−1}）和持续约40天的平台期。其后期 LC 衰减为 L_{波尔}∝ 吨^{−2.5}.另一个有趣的案例是 iPTF13dcc（Vreeswijk 等人，2017 年），其 LC 显示初始缓慢下降，持续时间为 ∼30 天，然后重新变亮并达到第二个峰值。
iPTF15esb 和 iPTF13dcc 的 LC 都挑战了上述所有模型。由^{56}完全捕获γ射线的Ni级联衰变为每天0.0098等，由磁星驱动的晚期LC（完全捕获γ射线）可以用L来描述_{inp，mag}∝ 吨^{−2}.磁星模型，以及^{56}镍级联衰变伴γ射线泄漏（Clocchiatti & Wheeler 1997;Chen 等人，2015 年;Wang等人，2015a）能够解释iPTF15esb的晚期LC，但该模型无法产生显示两个明亮峰和一个平台的LC。似乎需要一个涉及多次能量注入的能量源模型来解释iPTF15esb的奇异LC。Wang等人（2016b）提出了一个三重能源模型（^{56}Ni加磁星加相互作用），并用它来解释iPTF13ehe的LC。然而，该模型仅涉及喷射物和 CSM 壳之间的一次碰撞，并且不能产生波动的 LC。Vreeswijk等人（2017）采用的包含冷却发射和磁星自旋或喷射-CSM相互作用的模型来拟合iPTF13dcc的双峰LC，还不能重现iPTF15esb的LC中看到的波动。
然而，正如Yan等人（2017）所指出的那样，光谱和LC似乎有利于SN喷射物与不同半径的多个CSM壳或CSM团块之间的相互作用。在这里，我们提出了一个涉及SN喷射与多个壳层和恒星风之间相互作用的喷射-CSM相互作用模型，并使用该模型拟合iPTF15esb的LC。双重相互作用模型也有望解释iPTF13dcc的LC。
本文的结构如下。在第 2 节中，我们详细介绍了该模型，并将其应用于第 3 节中 iPTF15esb 和 iPTF13dcc 的 LC。最后，我们讨论了我们的结果，并在第4节中提出了我们的结论。
2. 多重喷射-CSM相互作用模型
在本节中，我们将正常的喷射物-CSM相互作用模型推广为涉及多个CSM壳和风的更复杂的模型。该模型的基本物理情况如下所述。喷射物与预先存在的CSM的相互作用导致形成两种冲击波：通过CSM传播的前向激波（FS）和扫过SN喷射物的反向激波（RS）（Chevalier 1982;Chevalier&Fransson，1994年）。这种相互作用通过将动能转化为辐射来提供强大的能量来源。
基于SN爆炸的数值模拟，可以采用SN喷射密度的破幂律分布（Matzner&McKee 1999）。外部喷射物的密度分布为
其中 n 是外侧喷射物的斜率，取决于 SN 祖星，g_{n}是密度分布缩放参数，由（Chevalier & Fransson 1994;Chatzopoulos 等人，2012 年)
where δ is the inner density profile slope. Here, E_{SN} is the total SN energy, and M_{ej} is the total mass of the SN ejecta. The relation between E_{SN} and M_{ej} can be written as (Chatzopoulos et al. 2012)
where x_{0} denotes the dimensionless radius of a break in the supernova ejecta density profile from the inner component to the outer component.
Before the SN explosion, the mass loss of a massive star could erupt several gas shells surrounding the progenitor. We assume that the density of a circumstellar shell or wind is
where q_{i} is a scaling constant, and s_{i} is the power-law index for a CSM density profile and therefore s_{i} = 2 indicates stellar winds, while s_{i} = 0 indicates uniform density shells. The subscript "i" denotes the ith collision between the ejecta and the CSM shell. For a steady wind (s_{i} = 2) with a constant pre-explosion mass loss rate and wind velocity v_{w}, we have .
The shocked CSM and shocked ejecta are separated by a contact discontinuity. The radius of the contact discontinuity R_{cd} can be described by a self-similar solution (Chevalier 1982)
where A_{i} is a constant. The radii of the FS and RS are given by
and
where R_{in,i} is the initial radius of the ith interaction (which is equal to the inner radius of the CSM density profile), β_{FS} and β_{RS} are constants representing the ratio of the shock radii to the contact-discontinuity radius R_{cd}. The values of β_{FS} and β_{RS} are determined by the values of n and s_{i}. They are given in Table 1 of Chevalier (1982). For n = 7 and s_{i} = 2, we can obtain β_{FS} = 1.299, β_{RS} = 0.970, and A = 0.27; for n = 7 and s_{i} = 0, we have β_{FS} = 1.181, β_{RS} = 0.935, and A = 1.2.
The interaction radii that are equal to the inner radii of the CSM density profile are given by
where t_{tr,i} is the trigger time of the ith interaction relative to the time zero-point. Here, we set the first interaction between ejecta and the CSM as the time zero-point. The kinetic energy of the ith interaction is
where E_{rad} is the energy loss due to radiation. The ejecta mass of the ith interaction is
The interaction between the ejecta and the CSM would convert the kinetic energy to radiation. The luminosity input function of the FS is (Chatzopoulos et al. 2012)
while the RS's input luminosity is (Wang et al. 2017)^{6}
where and represent the Heaviside step function that controls the end times of FS and RS, respectively, and t_{int,i} ≈ R_{in,i}/v_{SN,i} is the time when the ejecta-CSM interaction begins. The temporal index is . Here, we fix n = 7. Consequently, we have α_{i} = −0.143 for the shell (s_{i} = 0), and α_{i} = −0.6 for the steady wind (s_{i} = 2).
The RS termination timescale t_{RS,*} is the time once the RS sweeps up all available ejecta (Chatzopoulos et al. 2012, 2013)
Under the same assumption, the FS terminates when the optically thick part of the CSM is swept up. The termination timescale of the FS, being approximately equal to the time of FS breakout t_{FS,BO}, is given by (Chatzopoulos et al. 2012, 2013)
where M_{CSM,th,i} is the mass of the optically thick CSM
Here, R_{ph,i} denotes the photospheric radius of the ith CSM shell, located at the optical depth τ = 2/3 under Eddington's approximation. R_{ph,i} is given by
where κ is the optical opacity of the CSM and R_{out,i} is the radius of the outer boundary of the CSM. R_{out,i} can be determined by
Both the FS and the RS heat the interacting material. The total luminosity input from the FS and RS can be written as
where _{i} is the conversion efficiency from the kinetic energy. Chatzopoulos et al. (2012) assumed that = 100%, which is unrealistic in the actual situation, especially in the M_{CSM} ≪ M_{ej} case. Due to poor knowledge of the process of converting the kinetic energy to radiation, for simplicity, we set _{i} as a free parameter.
Because the expansion velocity of the CSM is much lower than the typical velocity of the SN ejecta, Chatzopoulos et al. (2012) assumed a fixed photosphere inside the CSM. Under this assumption, the output bolometric LC can be written as
where t_{diff,i} is the diffusion timescale in the optically thick CSM. The diffusion timescales of the ith interaction can be written as
where β = 4π^{3}/9 ≃ 13.8 is a constant (Arnett 1982), and c is the speed of light.
In this multiple-interaction model, the theoretical bolometric LC of N times interactions can be described as
We assume that the bolometric luminosity comes from the blackbody emission from the photosphere whose radius is R_{ph}, and therefore the temperature in our model can be estimated by
where σ_{SB} is the Stefan–Boltzmann constant. By assuming a stationary photosphere, we have .
3. Modeling the Light Curves of SLSNe with Multiple Peaks
In this section, we use the model described above to fit the bolometric LCs of iPTF15esb and iPTF13dcc. In order to derive the best-fitting parameters and determine the ranges of relevant parameters, we develop a Markov Chain Monte Carlo method that can minimize the values of χ^{2} divided by the number of degrees of freedom (χ^{2}/dof) for the multiple ejecta-CSM interaction model, and adopt this model to fit the LCs of iPTF15esb and iPTF13dcc.
To reduce the number of free parameters in our calculations, we fix several parameters. We adopt the power-law index of the outer density profile n = 7 as an approximation for Type I SNe (Chevalier 1982), and the inner density slope δ = 0.
The free parameters in our model are the opacities of the CSM shells and winds κ, the mass of the SN ejecta M_{ej}, the total mass of the CSM M_{CSM}, the density of the CSM at the interaction radius ρ_{CSM,in}, the interaction radius (the inner radius of CSM) R_{in}, the conversion efficiency from the kinetic energy to radiation , the time of the collision between the SN ejecta and the CSM shells t_{tr}.
The opacities of the CSM shells and winds κ are related to their composition and temperatures. For hydrogen-poor matter, the dominant source of opacity is electron scattering, κ = 0.06–0.2 cm^{2} g^{−1} (see the references listed in Wang et al. 2015b). For hydrogen-rich matter, κ = 0.33 cm^{2} g^{−1}, which is the Thomson electron scattering opacity for fully ionized material with the solar metallicity (Moriya et al. 2011; Chatzopoulos et al. 2012).
3.1. iPTF15esb
It is reasonable to assume that there are at least three collisions between the SN ejecta and the CSM shells, since the LC of iPTF15esb shows two prominent peaks and a plateau. In this scenario, the interaction between the SN ejecta and the stellar wind (i.e., s_{1} = 2) powers the first peak of the LC of iPTF15esb, while the second peak and the plateau are powered by the interactions between the SN ejecta and CSM shells (s_{2} = s_{3} = 0) at different radii.
The composition of the first and second CSM shells cannot be well constrained, since no hydrogen emission lines in the early-time spectra of iPTF15esb have been detected. On the other hand, the strong, broad Hα emission at ∼70 days after the first LC peak might be prompted by the interaction between the ejecta and a hydrogen-rich shell. Therefore, the values of the opacity of the first and second CSM shells κ_{1} and κ_{2} can be fixed to be 0.06–0.2 or 0.33 cm^{2} g^{−1}, and the value of the opacity of the third CSM shell κ_{3} is supposed to be 0.33 cm^{2} g^{−1}. By analyzing the Fe ii 5169 Å line, Yan et al. (2017) found that the photospheric velocity around the first peak of iPTF15esb is v_{ph} ≈ 17,800 km s^{−1} which can be set to be the value of the characteristic velocity v_{SN} of the SN ejecta.
The theoretical LC of iPTF15esb is shown in Figure 1. The best-fitting parameters and the corresponding confidence contour corners of iPTF15esb are shown in Table 1 and Figure 2, respectively. We find that the multiple-interaction model can explain the bumpy LC of iPTF15esb well (χ^{2}/dof = 1.71) and the parameters are reasonable. The derived physical parameters of the CSM shells and the wind are listed in Table 2. The masses of the optically thick part of the CSM shells M_{CSM,th}, which are close to their total mass, can be calculated. The termination timescales of the FS and the RS can also be determined. The optical depth of CSM τ_{CSM} > 1, indicating that these shells are opaque.
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Standard image High-resolution imageTable 1. The Fitting Parameters for iPTF15esb and iPTF13dcc
ith | s | κ | M_{ej} | M_{CSM} | ρ_{CSM,in}^{a} | ^{b} | t_{tr} | R_{in} | |
---|---|---|---|---|---|---|---|---|---|
Interaction | |||||||||
iPTF15esb | |||||||||
1 | 2 | 0.2 | −8.5 | ||||||
2 | 0 | 0.2 | 4.41^{c} | 5.2 | 2.11^{d} | ||||
3 | 0 | 0.33 | 5.87^{c} | 24.2 | 5.15^{d} | ||||
iPTF13dcc | |||||||||
1 | 0 | 0.2 | −73.79 | ||||||
2 | 0 | 0.2 | 21.31^{c} | −24.69 | 6.85^{d} |
Notes.
^{ a}ρ_{CSM,in} is the density of the CSM at radius . ^{ b} is the conversion efficiency from the kinetic energy to radiation. ^{ c}M_{ej,2} and M_{ej,3} are not fitting parameters, but were calculated by Equation (10). ^{ d}R_{in,2} and R_{in,3} are not fitting parameters, but were calculated by Equation (8).Download table as: ASCIITypeset image
Table 2. The Derived Physical Parameters
ith Interaction | M_{CSM,th}^{a} | t_{FS,BO} | t_{RS,*} | t_{diff} | τ_{CSM}^{b} | R_{out} | t_{erupt}^{c} | ||
---|---|---|---|---|---|---|---|---|---|
iPTF15esb | |||||||||
1 | 0.47 | 9.53 | 178.7 | 10.0 | 78.8 | 0.97 | ... | 0.19–1.9 | |
2 | 1.35 | 16.8 | 39.7 | 9.0 | 10.1 | 2.25 | 6.69 | ...^{d} | |
3 | 1.77 | 39.6 | 120.0 | 6.2 | 4.1 | 5.5 | 16.34 | ...^{d} | |
iPTF13dcc | |||||||||
1 | 7.04 | 12.7 | 29.6 | 38.9 | 225.5 | 1.42 | 1.55 | ...^{d} | |
2 | 17.23 | 25.5 | 44.67 | 27.6 | 12.14 | 6.94 | 21.72 | ...^{d} |
Notes.
^{ a}M_{CSM,th} is the mass of the optically thick CSM. ^{ b}τ_{CSM} is the optical depth of the CSM. ^{ c}t_{erupt} is the time of the progenitor star erupting the CSM shells before explosion. Here, we assume the velocity of the progenitor wind v_{w} = 100–1,000 km s^{−1} and the shell expansion velocities v_{shell} = 100 km s^{−1}. ^{ d}The shells were promptly expelled by some instability process and their mass-loss rates cannot be calculated.Download table as: ASCIITypeset image
Provided that the velocity of the shells v_{shell} is ∼100 km s^{−1} and using t_{shell} ≈ R_{in}/v_{shell}, we can obtain the time when the progenitor expelled the CSM shells before explosion. We infer that the progenitor of iPTF15esb has undergone at least two violent shell eruption processes at 6.69 and 16.34 years before the SN explosion, respectively, then experienced a wind-like mass loss whose mass-loss rate () is 0.19–1.9 M_{⊙} yr^{−1}.^{7}
3.2. iPTF13dcc
Vreeswijk et al. (2017) used a model combining the cooling emission from a shock-heated extended envelope (Piro 2015) and energy injection from a magnetar or an ejecta-CSM interaction to fit the LC of iPTF13dcc. While this model is plausible, we suggest that the double-collision model is also a possible model that can account for the LC of iPTF13dcc. In our scenario, the early-time bump of iPTF13dcc might be powered by the first collision between the ejecta and the CSM shell, while the second collision at a larger radius powers the late-time rebrightening of iPTF13dcc.
The early bump and the late-time rebrightening of iPTF13dcc are powered by the interactions between the SN ejecta and CSM shells at different radii, in which s_{1} = s_{2} = 0 is adopted. In our fitting, we adopt the expansion velocity v_{SN} = 10,000 km s^{−1}, which is the same as Vreeswijk et al. (2017). We assume the opacities of the first and second CSM shells κ_{1} = κ_{2} = 0.2 cm^{2} g^{−1}.
The theoretical LC of iPTF13dcc is shown in Figure 3. The best-fitting parameters and the corresponding confidence contour corners of iPTF13dcc are shown in Table 1 and Figure 4, respectively. The multiple-collision model can match the unusual light curve of iPTF13dcc (χ^{2}/dof = 2.25). The ejecta mass is M_{ej} = 14.2M_{⊙}, and the masses of CSM shells associated with the first and second collisions are 7.1M_{⊙} and 18.3 M_{⊙}, respectively. Adopting t_{shell} ≈ R_{in}/v_{shell} and v_{shell} ∼ 100 km s^{−1}, we can infer that the two shells had been expelled at 1.55 and 21.72 years before the explosion, respectively.
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Standard image High-resolution imageDue to the lack of data before the maximum brightness of the early bump of iPTF13dcc, it is difficult to determine the rise time of the first bump. Based on Table 2 and Figure 3, we find that the rise time of the early bump depends on the forward-shock termination timescale of the first interaction t_{FS,BO} = 17.8 days, while the decline rate of the late-time LC is determined by the RS of the second collision.
3.3. The Factors Influencing the LC Features
The observed properties of a core collapse SN are determined by several physical parameters, including the mass of ejecta M_{ej}, the kinetic energy of the SN ejecta E_{K}, the composition of the ejecta, and the structure of the envelope of the progenitor at the time of explosion. These properties result in different types of observed SNe. The LCs of some multi-peaked SNe show evidence of multiple interactions between the SN ejecta and the pre-existing CSM shells. In the interaction model, several parameters related to the CSM, e.g., the mass of CSM M_{CSM}, the density profile of the CSM, and CSM composition, must be taken into account. Different parameters would lead to different observational features.
The CSM properties are directly reflected by the LC shape. As shown in Figures 1 and 3, the luminosities provided by the FSs are usually larger than those provided by the RSs. The peak times of bumps depend on the timescales of the FSs t_{FS,BO}, while the RS affects the final decline rate. A less massive CSM tends to power narrower LCs and the larger the CSM density, the slower the decline. The peak luminosity is sensitive to the ejecta mass M_{ej}, so the explosion of a more massive star yields a brighter peak.
4. Discussion and Conclusions
Massive stars could be unstable and experience mass losses in the form of eruptions in the final stage of their lives (see, Smith 2014 and the references therein). Ofek et al. (2014) pointed out that more than 50% of the progenitors of type IIn SNe have experienced at least one pre-explosion eruption. In several cases, the progenitors of SNe could expel at least two shells and/or winds. Thus, it is expected that the interactions between the ejecta and multiple shells/winds would power bumpy LCs showing two or more peaks.
In this paper, we have studied two such bumpy SNe, iPTF15esb and iPTF13dcc, which show the undulation features that clearly deviate from the smooth rising and fading. The LC of iPTF15esb has two peaks and a post-peak plateau, while the LC of iPTF13dcc shows an early-time bump and a late-time rebrightening. All the previous energy-source models cannot account for these exotic features. We suggested that the LC undulations of iPTF15esb could arise from SN ejecta interacting with multiple dense CSM shells, which may be expelled by the eruptions of the progenitors. The interaction model for the LC of iPTF15esb is also favored by the broad Hα emission lines in the late-time spectra, which might be produced by the interaction of SN ejecta with a hydrogen-rich CSM shell located at a large distance from the progenitor star and was ejected by the progenitor star about 16.4 years before explosion.
To solve these problems, we generalize the "single" ejecta-CSM interaction model to the multiple-interaction model involving multiple CSM shells and/or winds. By employing this new model to fit the LC of iPTF15esb, we achieved rather good results and found that the first peak of the LC of iPTF15esb might be powered by the interaction between the SN ejecta and stellar wind, while both the second peak and the plateau might be powered by the two CSM shells at different radii. By fitting the LC, we found that the masses expelled by the progenitor of iPTF15esb are 0.49 and 1.46 M_{⊙} and that the mass-loss rate () of the wind is 0.19–1.9 M_{⊙} yr^{−1}, which is comparable to that of SN 1994W (, Chugai et al. 2004), SN 1995G (, Chugai & Danziger 2003), and iPTF13z (, Nyholm et al. 2017).
Furthermore, we also fitted the observed double-peaked LC of iPTF13dcc using this model and got rather satisfactory results. In this fit, the LC of iPTF13dcc was powered by the interactions between the ejecta and two CSM shells whose masses are 18.3 and 7.1 M_{⊙}, respectively. This positive result suggests that this model is also promising to account for the LCs of several double-peak SLSNe/SNe (e.g., Nicholl et al. 2016; Roy et al. 2016; Vreeswijk et al. 2017).
It is necessary to discuss the origins of the shells expelled by the progenitors of iPTF15esb and iPTF13dcc. Several models have been proposed to explain the violent pre-supernova eruptions. Woosley et al. (2007) and Woosley (2017) suggested that a very massive progenitor may undergo several episodes of pulsational pair instabilities and eject several massive shells before the SN explosion. The second possible origin is related to the binary interaction in which the shells ejected by the progenitor are supposedly formed by the large mass ejections from the progenitor interacting with its companion star (Podsiadlowski et al. 1992). In these successive collisions, the masses of the CSM shells decrease from the outermost shell (the first eruption) to the innermost shells (the final eruption), but their density increases (since the densities of the interiors of the progenitors are larger than that of the exterior). This property is consistent with the variation trend of our inferred densities of the ejected shells and the winds of the progenitors of iPTF15esb and iPTF13dcc.
The mass loss of the progenitor of an SN is an important process of stellar evolution. However, our understanding of the stellar mass loss mechanism remains incomplete. Further investigations for SLSNe/SNe like iPTF15esb and iPTF13dcc should shed light on the nature of the mass-loss histories of their progenitors.
We thank the anonymous referee for constructive suggestions that have allowed us to improve our manuscript significantly. We also thank Hai Yu, Weikang Zheng, Bing Zhang, Can-Min Deng, and Xue-Feng Wu for helpful discussions. This work was supported by the National Basic Research Program ("973" Program) of China (grant No. 2014CB845800), the National Key Research and Development Program of China (grant No. 2017YFA0402600), and the National Natural Science Foundation of China (grant No. 11573014). L.J.W. was also supported by the National Program on Key Research and Development Project of China (grant no. 2016YFA0400801). L.D.L. and S.Q.W. are supported by the China Scholarship Program to conduct research at UNLV and UCB, respectively.
Footnotes
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- 7
, where ρ(r) is the average density at a radius r, and v_{w} ≈ 100–1000 km s^{−1} is the terminal velocity of stellar wind of a hydrogen-poor star (He star; Smith 2014).