The long-term temperature evolution of neutron stars undergoing episodic accretion outbursts
Context:Transiently accreting neutron stars in low-mass X-ray binaries undergo episodes of accretion, alternated with periods of quiescence. During an accretion outburst, the neutron star heats up due to exothermic accretion-induced processes taking place in the crust. Besides the long-known deep crustal heating of nuclear origin, a likely non-nuclear source of heat, dubbed ‘shallow heating’, is often present at lower densities. Most of this accretion-induced heat generated in the crust slowly diffuses into the stellar core on a thermal diffusion time scale of years. Photon emission from the surface and neutrino emission from the core cool the star. Over many accretion outburst cycles, an equilibrium state is reached when the core temperature is high enough that the heating and cooling processes are in balance with each other.
Aims:Here we present a parameter study to investigate how stellar characteristics and outburst properties affect the long-term temperature evolution of a transiently accreting neutron star. For the first time the effects of the crustal properties are considered, in particular that of shallow heating.
Methods:We track the thermal evolution of a neutron star undergoing accretion outbursts over a period of yr, using our code NSCool. The outburst sequence is based on the regular outbursts observed from the neutron star X-ray transient Aql X-1. We calculate for each model the time scale over which equilibrium is reached and present these along with the temperature and luminosity parameters of the equilibrium state.
Results: We perform several simulations with scaled outburst accretion rates, to vary the amount of heating over the outburst cycles. The results of these models show that the equilibrium core temperature follows a logarithmic decay function with the equilibrium time scale. Secondly, we find that shallow heating significant contributes to the equilibrium state. Increasing its strength raises the equilibrium core temperature, while the depth of the shallow heat has a smaller effect on the long-term equilibrium state. We find that if deep crustal heating is replaced by shallow heating alone, the core would still heat up, reaching only a 2% lower equilibrium core temperature than if the star was only heated by deep crustal heating. Deep crustal heating may therefore not be vital to the heating of the core. Additionally, shallow heating can increase the quiescence luminosity to values higher than previously expected. The thermal conductivity in the envelope and crust are unable to significantly alter the long-term internal temperature evolution of the neutron star (variation in equilibrium core temperature between the models ). Moreover, a layer with low thermal conductivity at the bottom of the crust, such as a pasta layer, can not provide a heat barrier that prevents heat from flowing into the core. Stellar compactness and nucleon pairing in the core change the specific heat and the total neutrino emission rate from the core as function of temperature, with the consequences for the properties of the equilibrium state depending on the exact details of the assumed pairing models. The presence of direct Urca emission leads to the lowest equilibrium core temperature and the shortest equilibrium time scale.
A neutron star low-mass X-ray binary (LMXB) consists of a neutron star that accretes matter from a low-mass companion (typically with a mass ) via Roche-lobe overflow. While some systems are accreting persistently, most show accretion outbursts alternated with quiescent periods. The latter systems are referred to as transients and the cycles of accretion are often explained to be caused by instabilities in the accretion disc (see, e.g., the extensive review by Lasota 2001). The accretion outbursts can last from weeks to decades, and during the peak of the outbursts the neutron star can accrete matter at a rate (although for some the peak accretion rate is significantly lower, see the discussion in Wijnands et al. 2006). In quiescence, no or very little accretion takes place.
When the neutron star is accreting, a layer of material grows on the surface of the star. This freshly accreted material (composed of hydrogen and helium) is fused into heavier elements via stable and unstable thermonuclear processes in the envelope of the neutron star (see, e.g., the reviews by Bildsten 1998; Strohmayer & Bildsten 2006; Galloway & Keek 2017). As these processes occur at shallow depth (), most of the energy that is generated is rapidly radiated away from the surface and does not heat the core of the star (Brown 2000). Due to the continued buildup of material on the surface, the ashes of the thermonuclear burning are pushed to higher densities which causes reactions in deeper layers of the neutron star crust. This induces a series of deep crustal heating reactions (electron captures, neutron emission, and pycnonuclear reactions) which release energy deep in the crust (Sato 1979; Haensel & Zdunik 1990, 2008; Gupta et al. 2007; Fantina et al. 2018; Lau et al. 2018). Additionally, another heat source is inferred to be present from observational data in the outer crust (i.e., at densities lower than the neutron drip, ) of accreting neutron stars, referred to as the shallow heat source. A physical explanation for this shallow heat source has not been found (see Deibel et al. 2015; Waterhouse et al. 2016, for discussions of the different proposed hypotheses).
The accretion-induced heating processes cause the crust of the neutron star to be heated up during an outburst (Brown et al. 1998; Ushomirsky & Rutledge 2001). Although much of the heat flows towards the core, the crust and core get out of thermal equilibrium, because the specific heat of the core is significantly higher than that of the crust and hence the core takes longer to heat up (Yakovlev & Pethick 2004; Cumming et al. 2017). When the outburst ends, the crust of the neutron star is therefore hotter than the core and the crust will cool down in quiescence in order to restore crust-core thermal equilibrium. Neutron star crust cooling is a phenomenon that has now been observed in 9 systems (although in one system crust cooling needs to be confirmed; for an overview of crust cooling systems see the review by Wijnands et al. 2017, are reference therein). Observational campaigns to follow cooling neutron star transients as soon as they return to quiescence up to decades after the outburst ended are an important tool to gain insight in the physics of neutron stars. The cooling curves that these campaigns provide can be compared with theoretical models to constrain the crustal parameters of a source (e.g. Rutledge et al. 2002; Brown & Cumming 2009; Page & Reddy 2013; Turlione et al. 2015; Merritt et al. 2016; Ootes et al. 2019), and even of the core (Cumming et al. 2017; Brown et al. 2018).
The heat generated in a neutron star crust can be emitted either via photon emission from the surface or via various neutrino emission processes from the interior (see e.g. Page et al. 2006). However, if the cooling rate is unable to keep up with the rate at which heat from the crust flows towards to core, the neutron star will heat up over time. After many outbursts, the temperature of the core is sufficient to support a high enough cooling rate that can balance the heating and hence an equilibrium state is established (Brown et al. 1998; Colpi et al. 2001). Once a steady state has been reached, the core temperature will oscillate around the equilibrium core temperature (albeit only by a very small amount) during one accretion cycle (i.e., outburst and following quiescent period), but the average core temperature stays constant.
The properties of the equilibrium state – i.e., the time needed before equilibrium is reached, the equilibrium core temperature, and the corresponding surface temperature and photon and neutrino luminosities – depend on the outburst accretion rate and recurrence time, the heating and cooling processes at play, and on the characteristics of the neutron star (e.g., stellar compactness, crustal microphysics, nucleon pairing properties). In this research we investigate how these different properties contribute to the equilibrium state. This study complements previous research by Colpi et al. (2001) and Han & Steiner (2017). Whereas the previous studies were mainly focused on the effects of the core properties, we also set out to investigate if the assumed (micro)physical processes and properties of the crust can have a significant effect on the long-term thermal evolution of an accreting neutron star. Additionally, we have incorporated shallow heating in our model, and use an accretion rate that is based on observed light curves, which introduces more variation in accretion rate during and between outbursts.
In a simplified model in which the stellar interior is assumed to be always isothermal at a red-shifted temperature , the long term evolution of a transiently accreting neutron star would determined by the balance between heating and cooling described by
where is the stellar specific heat and , , and are, respectively, the redshifted heating, photon, and neutrino luminosities. The heating luminosity is provided by the amount of deep crustal heating () and shallow heating () that takes place during outburst. The cooling takes place via neutrino emission from the interior (primarily from the core), and photon emission from the surface of the star. For convenience we also define the total sum of the heat sources and sinks as the ‘effective’ heating luminosity . Once equilibrium has been reached, long-term time average values, denoted with , are then simply related by
To calculate the long-term equilibrium state of a transiently accreting neutron star beyond the above isothermal approximation we use our thermal evolution code NSCool (Page 2016). This code calculates the detailed thermal profile as function of time for a spherically symmetric neutron star based on the equations for heat transfer and energy conservation (taking into account general relativistic effects). We compute a basic model in which the star is subjected to short accretion outbursts and follow the thermal evolution of the neutron star towards an equilibrium state over a period of yr. We determine the parameters of this equilibrium state (equilibrium core temperature, luminosity, and the time it takes before equilibrium is reached) and then change the model parameters (describing the outburst accretion rate, recurrence time, stellar compactness, heating processes, crustal microphysics and assumed pairing gap models) one by one, to determine the effect of each of them on the parameters of the equilibrium state.
2.1 Accretion-induced heating
Both the deep crustal heating and the shallow heating are assumed to be directly proportional to the mass accretion rate (). In our models, we use the deep crustal heating processes as calculated by Haensel & Zdunik (2008), assuming an initial composition at the top of the crust (defined in our code as a boundary density ) of Fe, left behind by the thermonuclear burning in the envelope layers (which is defined here as ). A total amount of MeV per accreted nucleon is released by the electron capture, neutron emission, and pycnonuclear processes. Most of this heat () is released in the inner crust (above the neutron drip at ).
Shallow heating is introduced in our models by injecting an amount of heat per accreted nucleon in a region of the outer crust in the density range and is assumed to be uniformly distributed per volume. Because the origin of shallow heating is unknown, we present models in which we vary the strength and depth of the shallow heating. For several crust cooling sources the amount of shallow heating that is required to explain their observed cooling curves has been constrained to be (see for example Degenaar et al. 2011; Merritt et al. 2016; Parikh et al. 2019; Ootes et al. 2019). However, there are also sources that may not require shallow heat (e.g. Degenaar et al. 2015), and at least one that requires a very strong heat source in the outer crust (; Deibel et al. 2015; Parikh et al. 2017). Moreover, while for one source the shallow heating was constrained to be consistent between two consecutive outbursts (Parikh et al. 2019), for two others the amount and depth of this heating mechanism seems to vary between outbursts (Deibel et al. 2015; Parikh et al. 2017; Ootes et al. 2018). In the models presented here, we always assume the shallow heating strength and depth to be constant between outbursts.
To simulate accretion outbursts, we do not assume a constant accretion rate during a fixed outburst period, but instead take into account variations in the outburst history, based on observed light curves (Ootes et al. 2016). In this way we simulate accretion outbursts more realistically, although the effect of accretion rate variability on the long-term temperature evolution of a neutron star is minor (Cumming et al. 2017).111We have compared a model in which a time-dependent accretion rate was used with one in which a constant accretion rate was assumed and found the difference in equilibrium core temperature for the two models to be negligible (¡0.1%). However, the surface temperature in the crust-cooling phase after an outburst is very sensitive to variation in outburst properties. Here we use the accretion history of Aql X-1 observed over the past yr that we calculated for Ootes et al. (2018). We include here a more recent outburst from 2016 as well, extending the number of outbursts of the sequence to 24 (Degenaar et al. 2019). We repeat the sequence of outbursts until we reach a simulation time of yr (except for some models which required a longer simulation time to reach equilibrium). To be able to repeat the sequence of outbursts, we take into account a recurrence time of the last outburst of our sequence of , based on the detection of the start of a new outburst from this source in May 2017 (Vlasyuk & Spiridonova 2017). Aql X-1 is a rather useful source to use the light curve data from, because it shows (contrary to other sources) quite regular outbursts, that have been observed over a relatively long period. One might therefore assume the outburst history has been this regular for a long time. On average, the source shows outbursts with outburst accretion rates of the order that last months, and with recurrence times of almost a year. However, amongst the individual accretion episodes, there is still significant variation in outburst properties (e.g. Güngör et al. 2017; Ootes et al. 2018).
2.2 Heat transport, neutrino emission, and superfluidity
Heat that is generated in the crust during an outburst spreads towards the core of the star and towards the surface. The thermal conductivity inside the crust is set by electrons scattering off lattice impurities, phonons/ions, and other electrons (Yakovlev & Urpin 1980; Gnedin et al. 2001; Shternin & Yakovlev 2006). The first observational crust cooling studies revealed that the crust of an accreting neutron star cooled fast in quiescence, indicating that the thermal conductivity of this region is high (Wijnands et al. 2002, 2004; Cackett et al. 2006; Shternin et al. 2007). This observational result was initially unexpected from a theoretical perspective. To achieve such high thermal conductivity, the crust would need to be structured in a highly ordered lattice, rather than it being amorphous (Horowitz et al. 2009). The innermost region of the crust may host a so-called nuclear pasta region (Ravenhall et al. 1983; Hashimoto et al. 1984). In this region, the density approaches the nuclear density which causes a competition between the attractive nuclear force and the repulsive Coulomb force. As a consequence, matter can take on several pasta-like shapes (see for a recent review Caplan & Horowitz 2017). Such a region is predicted to have a low thermal conductivity, and can hence slow down the heat flow towards the core (e.g. Horowitz et al. 2015). In this study, we consider the effect of the thermal conductivity of the crust on the temperature evolution of the neutron star by changing the impurity factor (which quantifies how impure the lattice structure is) of the crust, and investigate if a low-conductivity pasta region can have a significant effect on the equilibrium core temperature of a transiently accreting neutron star.
The (redshifted) photon luminosity of the neutron star is set by
where is the radius of the neutron star at infinity, the Stefan-Boltzmann constant, and is the redshifted effective temperature. Within NSCool, the internal temperature profile is calculated up to the boundary density . We refer to the layers at as the envelope. This region acts as isolation between the hot crust and the surface. Hence, there is a strong temperature gradient across this thin ( tens of meters) layer. Detailed calculations have been carried out to derive relations to calculate the effective temperature for a particular boundary temperature, (e.g. Gudmundsson et al. 1983). For accreting neutron stars, the envelope strongly varies in composition as function of time during outburst due to thermonuclear burning (the details of which depend on the accretion rate and composition of accreted material) and the accumulation of newly accreted material (e.g. Woosley et al. 2004). As a consequence, the composition of the envelope in quiescence can differ from outburst to outburst (Brown et al. 2002), depending on when the last X-ray burst occurred and the type of the burst. The envelope consists of a layer of accreted material (hydrogen and helium) on top of layers containing heavier elements, which are products of thermonuclear burning. The lighter elements in the envelope have a higher thermal conductivity than heavier elements (e.g., Potekhin et al. 1999). Consequently, for a specific temperature at the bottom of the envelope, the effective temperature will be higher for an envelope that contains a thick light element layer, compared to an envelope containing more heavy elements (Potekhin et al. 1997). We use here the relations as outlined in Ootes et al. (2018), which depend on the light element column depth . The light element layer is assumed to consist of helium and carbon (helium has a stronger effect on the relation than hydrogen). During our simulations we generally keep the envelope composition constant between different outburst, although we use different light element column depths for two simulations to investigate the effect of variation in the thermal conductivity of the envelope on the equilibrium state.
Most of the heat generated in the crust will flow towards the core. The core of a neutron star has a high thermal conductivity (Baiko et al. 2001), such that (contrary to the crust) the core is always practically isothermal. The time scale over which the stellar core heats up depends on the specific heat, which is highly uncertain because it depends strongly on the state of the matter inside the core. If the nucleons are paired into Cooper pairs, forming a superfluid (neutron pairing) or superconductor (proton pairing), the specific heat can be either reduced or enhanced, depending on the temperature (see, e.g., Figure 13 in Page et al. 2004). Several spin-angular momentum channels are possible in which pairing can occur in the core, mostly likely P for neutrons and S for protons, but with large uncertainties about the value of the corresponding, density dependent, critical temperatures . In the temperature evolution models presented here we use several pairing gap models and also investigate the possibility that the nucleons are unpaired. We refer to Page et al. (2014) for an extensive review on superfluidity in neutron stars.
Cooling via neutrino emission can take place throughout the entire neutron star via various neutrino emission processes (see, e.g., the review by Yakovlev et al. 2001). These processes become more efficient at higher temperatures and the strongest neutrino emission processes take place in the neutron star core. The various neutrino processes are generally divided into fast and slow neutrino emission processes. The slow neutrino emission processes include modified Urca (mUrca), Bremsstrahlung and Cooper Pair Breaking and Formation (PBF). Fast neutrino emission in neutron stars can be generated by the direct Urca process (dUrca, Boguta 1981; Lattimer et al. 1991). The emissivity of this process is several orders of magnitude stronger than that of the slow neutrino emission processes, and hence direct Urca results in much stronger cooling (e.g. Page & Applegate 1992), but is expected to occur only in the core of massive neutron stars (see, e.g., Brown et al. 2018 for a recent example). Just as the specific heat, the strength of the neutrino emission in a neutron star depends on the presence of nucleon pairing and can be reduced or enhanced depending on the temperature.
2.3 Calculating the equilibrium state
For our models, we assume the equation of state for the core () from Akmal et al. (1998). This equation of state assumes only ‘regular’ neutron star material, and we do not consider the effects of exotic matter. For the crust, we use the accreted crust model from Haensel & Zdunik (2008). We thus assume that the neutron star has accreted for a period long enough to fully replace its original, catalysed crust (see Wijnands et al. 2013, for a discussion on partly-accreted crusts). Furthermore, we assume an initial (red-shifted) core temperature of at the start of the simulation, and repeat the sequence of accretion outbursts (24 outbursts over a duration of ) until the simulation reaches a time of yr. The heat flow towards the core is very different between outburst and quiescence, because of the very different temperature profiles in the crust in these two phases. Hence, once the star has reached an equilibrium between heating and cooling, the core temperature will still oscillate around the equilibrium temperature over an outburst cycle, even though the variation is very small (see Figure 2 in Colpi et al. 2001). We therefore collect the temperatures and luminosities at the end of each sequence of 24 outbursts, at the last time step before the start of the first outburst of the next sequence. This is (except for the models in which the recurrence time is increased) during the crust-cooling phase, because of the short recurrence time of the outbursts in Aql X-1 (Ootes et al. 2018). Consequently, the surface temperature and photon luminosity are different compared to what these values would be if the star were in crust-core equilibrium, but the core temperature and neutrino luminosity are not affected. We define the time at which equilibrium is reached from this data as the time when the temperature reaches of the maximum temperature.
3 Results and discussion
|( yr)||( K)||(eV)||( erg s)||( erg s)|
|7||13.0||1.026||99||2.78||25.2||No shallow heating (|
|8||16.8||0.972||102||3.12||17.0||No deep crustal heating (|
|13||16.2||1.007||106||4.23||37.9||(minimal cooling paradigm)|
|20||10.1||1.108||107||3.90||44.0||neutron gap model ‘CCDK’|
|21||10.5||1.108||107||3.90||44.1||no neutron pairing|
|22||11.4||1.837||129||8.08||39.4||neutron gap model ‘c’|
|23||37.6||1.732||126||7.47||40.2||no neutron pairing|
|24||7.2||1.109||107||3.90||44.0||proton gap model ‘CCDK’|
|25||10.4||1.109||107||3.90||44.0||no proton pairing|
3.1 The basic model
For our basic model (model 1), we take a neutron star with a mass , which gives for the EOS that we use a stellar radius of 11.5 km. The average outburst accretion rate is , which results in a long-term time-averaged accretion rate of . In addition to the heat released by deep crustal heating during outburst, we set the shallow heating strength to released over a density range . We assume that the crust has a high thermal conductivity and set , throughout the whole crust. For the envelope we take in our basic model . In regards to the pairing of nucleons, we assume the critical temperatures from the neutron superfluid gap calculated by Schwenk et al. (2003). This model has a maximum critical temperature and the gap continues up to the crust-core boundary such that the neutrons can be superfluid in the entire bottom of the crust. For the core pairing, we assume the neutron gap ‘a’ from Page et al. (2004), which is based on the results of Baldo et al. (1998), with , and the proton gap from Amundsen & Østgaard (1985), with . For this neutron star the central density is not high enough for direct Urca to occur, and hence only the neutrino emission processes from the minimal cooling scenario (Page et al. 2004) are allowed.
Fig. 1 shows the results for the temperature evolution (left panel, with the black curve the effective temperature and the red curve the redshifted core temperature) and the luminosities (right panel, with the blue curve the photon luminosity and the orange curve the neutrino luminosity). During one sequence of 24 outbursts, there is significant variation in the temperatures and luminosities. Because we are interested in the equilibrium state that is reached over many cycles of the outburst sequence, we plot in Fig. 1 the values for the temperatures and luminosities once per outburst sequence. We have chosen to plot the parameter values at the end of the quiescence period after the last outburst of our sequence (i.e., after every yr, at the last time step before the onset of the first outburst of the next sequence). However, in Ootes et al. (2018) we found that the recurrence times of the outbursts from Aql X-1 are so short, that – under the assumed heating behaviour and neutron star properties – at the onset of a new outburst some thermal energy remains in the crust of the neutron star, i.e., the temperature profile is not such as what would be the case for crust-core equilibrium. Since the outburst duration, outburst accretion rate, and recurrence time is different for each outburst in the sequence, the temperature profile is different at the onset of each new outburst. Consequently, in particular the effective temperature would be different if it would be ‘measured’ after a different outburst in the sequence. The grey areas in Fig. 1 indicate the spread on the respective curves that is introduced by taking the value after different outbursts. Although the choice of outburst after which the properties defining the equilibrium state are ‘measured’ affects the inferred effective temperature shown in Fig. 1, the core temperature evolution is not affected by this choice. Therefore, from here onwards we only consider the values taken at the end of each last outburst of the outburst sequence.
In Table 1 the properties of the equilibrium state are presented. For the basic model, we find that after yr the heating during the outbursts is in equilibrium with the cooling. The core temperature at equilibrium is . After yr the neutrino luminosity exceeds the photon luminosity in this model. When equilibrium is reached, the neutrino luminosity is a factor higher than the photon luminosity, and it is thus the neutrino luminosity that sets the equilibrium properties.
3.2 Assuming different outburst properties
3.2.1 Recurrence time
First we investigate the effect on the outcome of our simulations when assuming different outburst properties. Aql X-1 has outbursts with relatively short recurrence times of (on average). In model 2 we increase the average recurrence time by a factor 10 to 9.0 yr, while keeping the outburst strength and duration the same compared to model 1. Such a recurrence time in combination with the short outburst duration of Aql X-1, makes the outburst behaviour of model 2 more similar to other neutron star LMXB transients that recur less often than Aql X-1 (Chen et al. 1997). Therefore, this model might be more representative for the average transient neutron star LMXB.
As expected, increasing the recurrence time has a significant effect on the long term temperature evolution. Although the amount of heat that enters the star per outburst cycle remains the same compared to model 1, the source has a longer period over which it can cool, such that the time-averaged heating is lower than for model 1. Hence in its equilibrium state the core temperature is significantly lower () compared to model 1 (see Table 1 and Fig. 6) and the equilibrium photon luminosity is now about 50% of the neutrino one. Furthermore, in model 2 it takes about 4 times longer before the equilibrium state is reached (). From Eq. 2, if heating were balanced predominantly by neutrino losses with and is reduced by a factor ten as we imposed here, then should be reduced by , i.e., from K in model 1 down to K. However, in model 2, contributes significantly and hence a slightly lower is obtained.
One can estimate the thermal equilibration time as (Wijnands et al. 2013)
under the assumption that the specific heat is , with a temperature independent constant, as for a normal degenerate Fermion system. In our case we found that is reduced by a factor from model 1 to model 2 while is suppressed by a factor 10, so we expect that is increased by a factor 4.8 from Eq. 4. This is in good agreement with the numerically found factor .
3.2.2 Average outburst accretion rate
Next, we investigate the effect of assuming different average outburst accretion rates on the thermal equilibrium. In the basic model, the outbursts have an average outburst accretion rate of . In models 3–6 we multiply the accretion rates by a factor , , , and respectively, while keeping the other model properties the same as for model 1. As a result, the time-averaged accretion rates for models 3–6 become respectively , , , and . For model 3, the total time of the simulation was extended to yr, in order to determine the properties of the equilibrium state, because for such a low average accretion rate the core takes longer to heat up.
The results of models 3–6 are shown in comparison with model 1 in Fig. 7 and Table 1. For models 3–6, the results are as expected. Decreasing the average outburst accretion rate decreases the equilibrium core temperature ( for model 3, for model 4, and for model 5) and increases the time in which the equilibrium state is reached ( for model 3, for model 4, and for model 5), and vice versa for model 6 in which the outburst accretion rate is increased by a factor 2 ( and ). In model 3 the equilibrium core temperature is low enough that the photon luminosity is dominant over the neutrino luminosity (see right panel of Fig. 7). In these cases, again, the simple estimates of from Eq. 4 are in good agreement with the numerically found .
In Fig. 2, we plot the relation between the time-averaged accretion rate, equilibrium core temperature, and equilibrium time that we obtained for models 1, 3, 4, 5, and 6. The equilibrium core temperature decreases with time-averaged accretion rate, while the equilibrium time scale increases with time-averaged accretion rate. The equilibrium core temperature as function of the equilibrium time for these models can be fit with a logarithmic decay function.
Model 4 has the same time-averaged accretion rate as model 2 and naturally gives very similar results for the equilibrium values, but with nevertheless interesting differences (see Table 1). Although the time-averaged accretion rate is the same, the models differ in how these are achieved. Model 4 has a short recurrence time, with low outburst accretion rates, and model 2 has long recurrence times with high outburst accretion rates. As a result, during the outburst and the subsequent cooling phase the star in model 2 is hotter and thus energy losses, both and , are larger than in model 4: the effective heating luminosity is hence slightly lower, explaining the slightly lower equilibrium temperature and the slightly longer equilibrium time.
3.3 Changing the heating properties
In our basic model (model 1), heating during the outbursts was assumed to be caused by deep crustal heating (with a total strength of ) and shallow heating (with a strength of , released in the outer crust at a depth between ). In Fig. 8 we present model 7 in which we assume that only deep crustal heating takes place (and no shallow heating), and model 8, in which only shallow heating is activated. The properties of the heating mechanism that remains active in models 7 and 8 have not been changed compared to model 1. The results show that for model 1 the contribution of the deep crustal heating process is most significant to the heating of the neutron star to an equilibrium state, but that shallow heating is significant as well. When shallow heating is deactivated (model 7), the core temperature evolution curve is 7% lower compared to model 1 once equilibrium is reached (, see Table 1). For model 8 (no deep crustal heating), the equilibrium core temperature is 12% lower than for model 1 (). The difference between models 7 and 8 is due to the difference in heating strength assumed for the two heating mechanisms ( for deep crustal heating versus for shallow heating), and due to the difference in depth at which the heat is released (primarily in the inner crust for deep crustal heating and in the outer crust for shallow heating).
Fig. 8 shows that model 7 has a higher core temperature than model 8 throughout the simulation, but that the surface temperatures of these two models are inverted; the surface temperature is higher for model 8 than for model than for model 7. The reason for this trend is that in model 8 (only shallow heating), due to the shallow depth of the heat source, a larger fraction of the heat flows to the surface compared to model 7 (only deep crustal heating). This causes the surface temperature to be higher throughout the crust-cooling phase in quiescence (when the surface temperatures are ‘measured’ from the models).
In models 9 and 10 the effect of varying the depth of heat release is investigated. In these two models both shallow heating and deep crustal heating are assumed to take place. We keep the shallow heating strength equal to that of model 1, but in model 9 we place it at shallower depth (in the density range ), and in model 10 we place it deeper (in the density range ). Fig. 9 shows that the equilibrium core temperature is highest for the model that has the deepest shallow heating (model 10), indicating that a larger fraction of the heat flows towards the core for deeper heating resulting in a slightly larger , although the differences are marginal. The equilibrium core temperature for model 9 (shallower heating) is lower compared to model 1, and for the model that assumes deeper shallow heating (model 10) the core temperature in equilibrium is higher than for the basic model (see Table 1).
The evolution of the effective temperature shown in the left panel of Fig. 9 is compatible with the evolution of , but with a twist. is determined by the temperature in the outer layers of the crust and the values we display in the figure are ‘measured’ at the end of the last cooling phase (which is 236 days after the outburst end) in each sequence of 24 outbursts when the crust as somewhat, but never completely, relaxed from the accretion heating. When compared to model 1, in model 10 a larger part of flowed inward into the core resulting in cooler crust and a lower at the late time when we ‘measure’ it while in model 9 an even larger part of flowed outward resulting in a higher at early time during the post accretion cooling phase (not shown in the figures) but in a cooler crust and lower at late time, when we ‘measure’ it for the figures. As time runs and the star reaches the equilibrium state the difference in between models 1 and 10 becomes insignificant while in model 9 the surface losses are large enough that remains lower. The differences are nevertheless very small and smaller than the statistical error in the observed .
Finally, Fig. 10 shows the effect of varying the shallow heating strength on the equilibrium state of the neutron star. In this figure, model 1 is compared to model 7 in which the shallow heating is deactivated, and model 11 in which the shallow heating strength is increased to . Comparing to models 9 and 10, the effect of variation in shallow heating strength is much more significant than when the depth of the shallow heating was varied. For stronger shallow heating () the equilibrium core temperature is higher () than for model 1. The effect is similar compared to model 6 (see Fig. 7) in which the accretion rate was increased by a factor 2. In both models an effectively similar amount of heat is released in the crust during outbursts (i.e., in total from shallow heating and deep crustal heating for model 11 versus but at a twice higher rate for model 6), but because for model 11 the heating is only increased at shallow depths (while in model 6 more heat is released from deep crustal heating as well), the equilibrium core temperature in model 11 is slightly lower (, see Table 1) compared to model 6.
From these models it becomes clear that shallow heating can play an important role in the long-term heating of transiently accreting neutron stars (assuming that shallow heating is active during all outbursts). The strength of the shallow heating has a stronger effect on the equilibrium core temperature than the depth of the shallow heating. Even when shallow heating is placed just below the envelope, most of the heat released during outbursts will flow towards the core rather than the surface. As shown in Fig. 8, if deep crustal heating is not active, shallow heating heats up the core of the neutron star. The equilibrium core temperature is lower than if only deep crustal heating is assumed to take place and this difference would be smaller () if the same heating strength was assumed for the two heating mechanisms. Moreover, for shallow heating strengths in excess of the heating strength from deep crustal heating, the contribution of shallow heating to the equilibrium state can become dominant. Deep crustal heating is therefore not essential to long-term heating of transiently accreting neutron stars if significant shallow heating takes place.
These results also have consequences for the position of observations from neutron star LMXBs on the plot, where is the observed quiescent luminosity (Yakovlev & Pethick 2004; Heinke et al. 2009, 2010; Wijnands et al. 2013, 2017; Han & Steiner 2017). The heating curves with which the observations are compared to determine which cooling mechanisms might be at play, generally assume that , with for the assumed heating rate. However, if the contribution from shallow heating is taken into account, the heating luminosity, and hence is higher. It should then possible to find sources that have a quiescent luminosity that exceeds the predicted photon cooling-curve.
3.4 Changes in the composition of the neutron star envelope
In models 14 and 15 we vary the envelope composition compared to model 1, in which we assumed . In model 14 we decrease the light element layer to and in model 15 we use a thicker layer of light elements (). For each model we assume that the light element layer consist of a helium layer on top of a carbon layer with .
The results show that these variations do not cause significant changes in the evolution of the core temperature. For the three models shown in Fig. 11 the difference between the models is most prominent in the surface temperature and photon luminosity. This is as expected, since the only difference between the models is how easily the heat that flows towards the surface (while most of the heat flows towards the core) can pass the envelope layer. For the model with the thickest layer of light elements, the thermal conductivity in the envelope is the highest, and therefore heat can leave the star most easily. Consequently, this star has a significantly higher effective temperature for similar crust temperatures. For model 14 and 15, there is only a very small difference in the equilibrium core temperature (i.e., a difference of compared to model 1, see Table 1). This different in equilibrium core temperature is caused by the fact that for the smaller light element layers less heat can flow away from the surface over an outburst cycle, so that relatively more heat flows towards the core. However, this effect is small, since for all envelope compositions most of the accretion induced heat flows towards the core.
3.5 Assuming different thermal conductivity in the crust
The next property that we consider is the thermal conductivity in the crust. By increasing the impurity factor of the crust, impurity scattering is enhanced which decreases the thermal conductivity. In model 1 the impurity factor was set to a low value () throughout the crust. Typical values used for modelling crust cooling are of the order up to (e.g. Brown & Cumming 2009; Merritt et al. 2016; Turlione et al. 2015; Cumming et al. 2017; Parikh et al. 2019), although some sources may have higher impurity (regions) in the crust (Degenaar et al. 2014; Ootes et al. 2019). In models 16 and 17 the impurity is increased to and respectively. This decreases the thermal conductivity in the crust, primarily for , because the effect on the thermal conductivity at lower densities in the crust is largely dominated by electron-ion scattering (depending on the temperature and value of , see left panel of Fig. 3).
|(M)||(km)||(km)||( cm s)||( g cm)|
Fig. 12 shows that the increased impurity (lower thermal conductivity) has a small effect on the thermal evolution of the core. The effect is similar to what was determined for variation in the envelope composition and shallow heating depth. For lower thermal conductivity in the crust (higher ) the core equilibrium temperature is somewhat lower, but the difference in comparison to the basic model is small; for model 16 and for model 17 (see Table 1). Although the thermal conductivity in the crust is decreased for higher , the thermal conductivity in the inner crust is still high compared to the outer crust, and the direction of the heat flow is only slightly altered. Most of the heat generated in the crust is transported to the core in the crust-cooling phase. For the basic model, the internal luminosity at the crust-core boundary is 11 times higher than at the crust-envelope boundary once the long-term equilibrium is reached (at the time in each cycle of outbursts at which we ‘measure’ the thermal properties, which is during the crust-cooling phase). Because the thermal conductivity is mostly decreased in the inner crust by increased , heat flow towards the core is slower compared to the low-impurity crust, and more heat flows towards the surface (as can been seen from the photon luminosity curve in the right panel in Fig. 12). However, even though the luminosity (once the long-term equilibrium is reached) at the crust-envelope boundary is almost twice as high for model 17 compared to model 1, the luminosity at the crust-core interface is still 6 times higher than at than at the crust-envelope boundary. All in all, despite the lower thermal conductivity assumed in models 16 and 17, the variation in the amount of heat that flows towards the core rather than the surface is small (in all cases most of the heat flows towards the core), causing the equilibrium core temperatures of these models to be similar to the model that assumes high thermal conductivity.
The right panel in Fig. 12 shows that at early times ( yrs), the neutrino luminosity is higher for the models with lower thermal conductivity in the crust (models 16 and 17). Due to the lower crustal thermal conductivity for models 16 and 17 the temperature in the crust at the time at which the model quantities are ‘measured’ is higher than the crustal temperature in model 1 at that time (see left panel in Fig. 12). This causes the neutrino emissivity in the crust to be high for models 16 and 17, and, moreover, dominant over the core neutrino processes because the core temperatures at that time are still low. The difference disappears when the core of the neutron star becomes hot enough that the core neutrino emission dominates over that of the crust.
In models 18 and 19 we investigate if a low conductivity pasta layer (with respectively and at ) in the crust is able to affect the heat flow significantly. In these models the conductivity in the outer regions is still assumed to be high (). One might expect that a high-impurity pasta layer can provide a sufficient heat barrier to force less heat to flow to the core. However, Fig. 13 shows that this is not the case. The difference in equilibrium core temperature between models 1, 18, and 19 is negligible. The locally increased impurity does provide a heat barrier, as the thermal conductivity close to the crust-core boundary decreases by a factor of and for models 18 and 19 respectively, compared to model 1 (see right panel of Fig. 3). The heat flow in this specific region is hence slowed down, causing the crust to be hotter compared to model 1 in the crust-cooling phase (as can be observed from the surface temperature in Fig. 13). However, the heat barriers in these two models are not strong enough to prevent heat from flowing into the core. Even when the impurity factor in the pasta layer is increased to , the difference in equilibrium core temperature compared to model 1 is only .
3.6 Variation in stellar compactness
In this subsection we will discuss the effect that the stellar compactness has on the thermal evolution of transiently accreting neutron stars. However, we restrict ourselves here to cases that assume the minimal cooling paradigm (i.e., no direct Urca emission, see Page et al. 2004). Enhanced cooling will be discussed in Section 3.8. In our basic model we assume a neutron star with mass and radius . Such a neutron star has a surface gravity , a crust thickness , a surface redshift , and a central density . In model 12 we assume a lower mass of and in model 13 a heavy neutron star with a mass of is assumed. The stellar properties of these models are summarised in Table 2.
The results presented in Fig. 14 show that for smaller masses (smaller compactness), the core reaches higher equilibrium temperatures and that the equilibrium state is reached faster than for heavier stars ( and yr for model 12, and and yr for model 13, see Table 1). For a higher stellar mass, the neutron star has a higher central density. For the chosen neutron paring gap in our models, this means that critical temperature curve for the pairing gap is cut off in the centre at temperatures K for the star, while the curve continues to lower temperatures for the heavier stars ( K for and K for ). This has consequences for both the specific heat and neutrino luminosity of these stars. On the one hand, the specific heat as function of temperature (over the range of temperatures considered here), is higher for the higher mass stars, because neutron pairing enhances the specific heat for the range (see Page et al. 2004). As a consequence, the amount of heat needed to heat up the core is larger, and therefore the equilibrium state is reached over a longer time scale. On the other hand, the neutrino luminosity is stronger for higher mass stars at the same temperature (as can be determined from comparison of the left and right panels in Fig. 14), because PBF neutrino emission becomes significant (and consequently dominant over all other neutrino emission processes in these models) for core temperatures in the range . This causes an equilibrium to be established at lower core temperatures for the models in which a higher mass is assumed. Finally, the difference in neutrino luminosity as function of temperature for the different models, in combination with the different specific heat as function of temperature is what causes the neutrino luminosity curves in the right panel of Fig. 14 to intersect.
3.7 Assuming different pairing characteristics
In this subsection we focus on pairing characteristics and their effects on the neutrino emission. In model 1 we use for the pairing of dripped neutrons in the crust the ‘SFB’ gap model from Schwenk et al. (2003). In model 20 we use instead the neutron gap model ‘CCDK’ calculated by Chen et al. (1993), which has a maximum critical temperature . Contrary to the gap model used in model 1, the one assumed in model 20 closes inside the crust at a density . For a neutron gap that closes in the crust, a layer of non-paired neutrons may exist at the bottom of the crust as long as the neutron gap starts at a higher density. Such a layer affects the heat flow towards the core (see the discussion by Deibel et al. 2017). In model 20 we force the existence of a layer of non-paired neutrons at the bottom of the crust by cutting off the pairing gap at the crust-core boundary. In model 21 we assume that the crust is completely non-superfluid. The effects of the two models compared to the basic model are very small; the equilibrium core temperature of models 20 and 21 is lower than for model 1. A small difference in core temperature evolution over the period in which the equilibrium state is reached can be observed from Fig. 15. Due to the pairing of neutrons in the inner crust the specific heat is lower in this region for models 1 and 20 than for model 21, such that the thermal diffusivity in this region is higher. The heat flow towards the core is somewhat higher and hence the core can heat up faster. This explains why for the models in which the free neutrons in the crust are paired the equilibrium temperature is reached faster (see Table 1; the equilibrium time scales for models 1 and 20 are respectively 4% and 8% lower compared to model 21).
Next, we consider neutron pairing in the core of neutron stars. In model 1, we assume the gap model ‘a’ from Page et al. (2004), which has a maximum critical temperature K. In model 22 we assume the neutron gap model ‘c’ from Page et al. (2004) instead. This gap model has a higher maximum critical temperature of K, such that the neutrons form stronger pairs for the typical temperatures reached in the core for these models. To consider the other extreme situation, we assume in model 23 that no neutron pairing takes place. In that case the neutrons in the core are unpaired (but the protons can still form pairs), except for the outermost region of the core up to where the assumed neutron gap reaches. The neutron gap closes at a density . For both models 22 and 23 the core reaches an equilibrium core temperature that is much higher than for our basic model ( K for model 22 and K for model 23, see Fig. 16). There are two factors that are responsible for the observed effects between these three models. First of all, there is a difference in specific heat between the three models. In models 1 and 22, the core specific heat is suppressed for K due to neutron pairing. Because the critical temperatures for the pairing gap assumed in model 22 are higher than for model 1, the specific heat suppression is stronger. As a consequence, for model 22 the equilibrium time scale is the shortest, and vice versa model 23 (no pairing, and hence the highest specific heat) has the lowest core heating rate. Secondly, there are large differences in neutrino emission from the PBF process between these models (see Fig. 4). This neutrino emission process is – in the absence of direct Urca emission – the strongest neutrino emission process, and it becomes important when the temperature increases and reaches values close to the critical temperature for nucleon pairing (see, e.g., Fig. 5 in Page et al. 2006, for the control function for this process as function of ). For gap ‘a’ (Page et al. 2006, assumed in model 1), the neutron pairing gap comprises the entire core with critical temperatures in the range K, while for gap ‘c’ (Page et al. 2006, assumed in model 22) the critical temperatures are in the range ( K). This means that in model 1, the PBF process becomes efficient at much lower core temperatures than for model 22, such that equilibrium between heating and cooling is reached at much lower core temperatures (see Fig. 4). For model 23, neutrino emission from the PFB process for neutron is absent, and it is the modfied Urca process that controls the equilibrium core temperature. Because the neutrons in the core are unpaired, the modfied Urca luminosity is higher as function of temperature than for models 1 and and 22. The total neutrino luminosity as function of temperature is higher in model 23 than in model 22, and hence equilibrium is reached at a lower core temperature.
The last pairing gap to consider is the proton gap, which controls the pairing of the small fraction of protons in the core of neutron stars. We compare here again the pairing gap assumed in model 1 with a model in which proton pairing is stronger (model 24) and a model in which protons are assumed to be unpaired (model 25). In model 1 we assumed the proton pairing gap from Amundsen & Østgaard (1985), which has a critical temperature function that reaches from the crust-core boundary up to a density . The function peaks at low density in the core and has a maximum critical temperature K. The gap model assumed in model 24 is from Chen et al. (1993). It comprises the full core and has a maximum critical temperature of K. The results in Fig. 17 show that the assumptions for proton pairing do not affect the equilibrium core temperature. As can be seen from the left panel in Fig. 4, the neutrino emission from the proton PBF process is negligible compared to the neutron PBF emission in model 1 at temperatures K. This depends on the assumed gap model and on the temperature. For models 24 the contribution from proton PBF neutrino emission at K to the total neutrino emissivity is similar compared to model 1 and absent for model 25, such that for both models the total neutrino emission at the temperature at which equilibrium is reached is the same as for model 1 (see right panel in Fig. 17). There is however a difference in the time scale over which equilibrium is reached between models 24, 25 and the basic model due to differences in specific heat. At temperatures K, the specific heat for the two models that assume proton pairing (model 1 and model 24), is reduced because the internal temperature is significantly lower than the critical temperature for the assumed proton gap models. Because the critical temperatures assumed in model 24 are the highest, the specific heat reduction in model 24 is stronger than for model 1. The neutron star in model 24 can therefore heat up most easily as can be observed from Fig. 17. Similarly, because no proton pairing takes place in model 25, this model has the highest specific heat at low temperatures such that the star heats up more slowly and the equilibrium time scale is higher than for models 1 and 24.
3.8 Direct Urca emission
The direct Urca process is the most efficient neutrino emission process that might take place in the cores of heavy neutron stars, under the requirement that the fraction of protons is large enough (; Lattimer et al. 1991). In model 26 we assume a neutron star with a mass , which allows for direct Urca to take place in the innermost part of the core at densities . The direct Urca emission region covers % of the total volume of the star.
In Fig. 18 we compare the results of this model with model 13, in which a neutron star with the same mass is assumed, but where the direct Urca emission process is not taken into account. In model 26 the neutron star reaches a very low equilibrium core temperature of K in yr. Fig. 5 shows the neutrino emission luminosity as function of temperature as used in NSCool for model 26, which shows that direct Urca emission is orders of magnitude larger than any of the other neutrino emission processes. Due to the fast increase in luminosity with temperature, the neutrino emission from direct Urca emission is at low temperatures already strong enough to balance the heating rate from the accretion outbursts. The equilibrium time scale of model 26 is more than 50 times lower than that of model 13. If direct Urca takes place in only a small fraction of the core such as might be the case for MXB 1659-29 (Brown et al. 2018), the neutron star is very likely to be in equilibrium state.
In this parameter study we investigated the long-term temperature evolution of transiently accreting neutron stars. Assuming that such neutron stars have regular accretion outbursts, they will reach an equilibrium state in which the heat generated in the crust by accretion-induced reactions is balanced by neutrino cooling from the core and photon emission from the surface. We compared how the properties of neutron star crusts and cores, as well as their outburst properties affect the equilibrium state : the equilibrium core temperature and the time over which this temperature is reached. We used the observed outburst behaviour of Aql X-1 over the past yr and extrapolated the sequence of 24 outbursts observed in this period to model such outbursts for yr using our thermal evolution code NSCool, assuming that the neutron star heats up from a cold state.
The results are as follows.
As expected, we find that outburst properties, such as the recurrence time and outburst accretion rate strongly affect the equilibrium state. These properties set the time-averaged accretion rate and thus affect the amount of heat that is injected into the neutron star over an outburst cycle. We find that the equilibrium core temperature as function of equilibrium time scale follows a logarithmic decay trend, for decreasing accretion rate.
We find that shallow heating can significantly increase the equilibrium core temperature, independently of the depth of this heat source. Even when the shallow heating is injected just below the envelope on the neutron star, most of the heat flows towards the core rather than the surface. The strength of the shallow heating therefore affects the long term equilibrium state of neutron stars. When shallow heating is the only assumed heating process, the core will reach an equilibrium core temperature similar to the equilibrium core temperature for a model that only takes into account deep crustal heating (the difference is in our models if the same heating strength is assumed for the heat sources). From these models we conclude that it is possible to heat accreting neutron stars only by the shallow heating process. Additionally, stars with strong shallow heating may be observed to have higher quiescent luminosities than the calculated photon cooling curves, because these curves do not take into account the additional contribution of shallow heating to the heating rate of the star.
Under the assumption of minimal cooling and for the assumed pairing gaps in our basic model, increasing the stellar compactness increases the neutrino luminosity. Naturally, the stronger cooling by neutrino emission from the core in these models leads to a lower equilibrium core temperature for heavier (more compact) neutron stars.
The envelope composition has a strong effect on the observed surface temperature neutron stars, but changing the thermal conductivity of this heat blanket (by changing the composition of the envelope) has minor effects on the equilibrium state (differences in equilibrium core temperature are compared to the basic model).
Decreasing the thermal conductivity (by increasing the impurity factor) in the crust delays the heat flow, but even an impurity factor is not strong enough to change the equilibrium state significantly. The difference in equilibrium core temperature is compared to the model in which is assumed. Moreover, a low-conductivity pasta layer (with or ) can not provide a heat barrier that prevents heat from flowing towards the core.
It is well known that the exact details of the assumed pairing gaps for neutron superfluidity and proton superconductivity in the core of neutron stars have strong consequences of the properties of the core. At the pairing of nuclei affects the specific heat and the neutrino emission. However, both quantities are very sensitive to the local , and hence the effects on the equilibrium core temperature and the time scale over which equilibrium is reached depend strongly on the exact details of the assumed models for neutron and proton pairing. The pairing of dripped neutrons in the crusts of neutron stars has a minor effect on the equilibrium state.
The strongest effect on the equilibrium state of neutron stars is provided by the direct Urca neutrino emission process. If the nucleons in the innermost region of the core of massive neutron stars are unpaired, and the proton fraction of the material is large enough this process can take place. Because it has a much higher emissivity at a given temperature than any of the other neutrino emission processes, it can balance the heat that enters the core at much lower temperatures than when via the minimal cooling processes are assumed to take place. Hence a neutron star in which direct Urca emission takes place will be much colder when it reaches the equilibrium state. We have shown here that in the presence of a direct Urca process the time to reach the equilibrium state is much shorter, by one to three orders of magnitude, and can be as small as a few hundred years, confirming the estimates of Wijnands et al. (2013)
Acknowledgements.LSO and RW receive support from a NWO TOP Grant, module 1, awarded to RW. DP is partially supported by the Mexican Conacyt (CB-2014-01, #240512). This work was supported in part by the National Science Foundation under Grant No. PHY-1430152 (JINA Center for the Evolution of the Elements).
- Akmal et al. (1998) Akmal, A., Pandharipande, V. R., & Ravenhall, D. G. 1998, Phys. Rev. C, 58, 1804
- Amundsen & Østgaard (1985) Amundsen, L. & Østgaard, E. 1985, Nuclear Physics A, 437, 487
- Baiko et al. (2001) Baiko, D. A., Haensel, P., & Yakovlev, D. G. 2001, A&A, 374, 151
- Baldo et al. (1998) Baldo, M., Elgarøy, Ø., Engvik, L., Hjorth-Jensen, M., & Schulze, H.-J. 1998, Phys. Rev. C, 58, 1921
- Bildsten (1998) Bildsten, L. 1998, in NATO Advanced Science Institutes (ASI) Series C, Vol. 515, NATO Advanced Science Institutes (ASI) Series C, ed. R. Buccheri, J. van Paradijs, & A. Alpar, 419
- Boguta (1981) Boguta, J. 1981, Physics Letters B, 106, 255
- Brown (2000) Brown, E. F. 2000, ApJ, 531, 988
- Brown et al. (2002) Brown, E. F., Bildsten, L., & Chang, P. 2002, ApJ, 574, 920
- Brown et al. (1998) Brown, E. F., Bildsten, L., & Rutledge, R. E. 1998, ApJ, 504, L95
- Brown & Cumming (2009) Brown, E. F. & Cumming, A. 2009, ApJ, 698, 1020
- Brown et al. (2018) Brown, E. F., Cumming, A., Fattoyev, F. J., et al. 2018, Physical Review Letters, 120, 182701
- Cackett et al. (2006) Cackett, E. M., Wijnands, R., Linares, M., et al. 2006, MNRAS, 372, 479
- Caplan & Horowitz (2017) Caplan, M. E. & Horowitz, C. J. 2017, Reviews of Modern Physics, 89, 041002
- Chen et al. (1993) Chen, J. M. C., Clark, J. W., Davé, R. D., & Khodel, V. V. 1993, Nuclear Physics A, 555, 59
- Chen et al. (1997) Chen, W., Shrader, C. R., & Livio, M. 1997, ApJ, 491, 312
- Colpi et al. (2001) Colpi, M., Geppert, U., Page, D., & Possenti, A. 2001, ApJ, 548, L175
- Cumming et al. (2017) Cumming, A., Brown, E. F., Fattoyev, F. J., et al. 2017, Phys. Rev. C, 95, 025806
- Degenaar et al. (2011) Degenaar, N., Brown, E. F., & Wijnands, R. 2011, MNRAS, 418, L152
- Degenaar et al. (2014) Degenaar, N., Medin, Z., Cumming, A., et al. 2014, ApJ, 791, 47
- Degenaar et al. (2019) Degenaar, N., Ootes, L. S., Page, D., et al. 2019, submitted to MNRAS
- Degenaar et al. (2015) Degenaar, N., Wijnands, R., Bahramian, A., et al. 2015, MNRAS, 451, 2071
- Deibel et al. (2015) Deibel, A., Cumming, A., Brown, E. F., & Page, D. 2015, ApJ, 809, L31
- Deibel et al. (2017) Deibel, A., Cumming, A., Brown, E. F., & Reddy, S. 2017, ApJ, 839, 95
- Fantina et al. (2018) Fantina, A. F., Zdunik, J. L., Chamel, N., et al. 2018, A&A, 620, A105
- Galloway & Keek (2017) Galloway, D. K. & Keek, L. 2017, in Timing Neutron Stars: Pulsations, Oscillations and Explosions, ed. T. Belloni & M. Mendez, Astrophysics and Space Science Library (Springer International Publishing), to be published, arXiv e–print: 1712.06227
- Gnedin et al. (2001) Gnedin, O. Y., Yakovlev, D. G., & Potekhin, A. Y. 2001, MNRAS, 324, 725
- Gudmundsson et al. (1983) Gudmundsson, E. H., Pethick, C. J., & Epstein, R. I. 1983, ApJ, 272, 286
- Güngör et al. (2017) Güngör, C., Ekşi, K. Y., & Göğüş, E. 2017, New A, 56, 1
- Gupta et al. (2007) Gupta, S., Brown, E. F., Schatz, H., Möller, P., & Kratz, K.-L. 2007, ApJ, 662, 1188
- Haensel & Zdunik (1990) Haensel, P. & Zdunik, J. L. 1990, A&A, 227, 431
- Haensel & Zdunik (2008) Haensel, P. & Zdunik, J. L. 2008, A&A, 480, 459
- Han & Steiner (2017) Han, S. & Steiner, A. W. 2017, Phys. Rev. C, 96, 035802
- Hashimoto et al. (1984) Hashimoto, M., Seki, H., & Yamada, M. 1984, Progress of Theoretical Physics, 71, 320
- Heinke et al. (2010) Heinke, C. O., Altamirano, D., Cohn, H. N., et al. 2010, ApJ, 714, 894
- Heinke et al. (2009) Heinke, C. O., Jonker, P. G., Wijnands, R., Deloye, C. J., & Taam, R. E. 2009, ApJ, 691, 1035
- Horowitz et al. (2015) Horowitz, C. J., Berry, D. K., Briggs, C. M., et al. 2015, Physical Review Letters, 114, 031102
- Horowitz et al. (2009) Horowitz, C. J., Caballero, O. L., & Berry, D. K. 2009, Phys. Rev. E, 79, 026103
- Lasota (2001) Lasota, J.-P. 2001, New A Rev., 45, 449
- Lattimer et al. (1991) Lattimer, J. M., Pethick, C. J., Prakash, M., & Haensel, P. 1991, Physical Review Letters, 66, 2701
- Lau et al. (2018) Lau, R., Beard, M., Gupta, S. S., et al. 2018, ApJ, 859, 62
- Merritt et al. (2016) Merritt, R. L., Cackett, E. M., Brown, E. F., et al. 2016, ApJ, 833, 186
- Ootes et al. (2016) Ootes, L. S., Page, D., Wijnands, R., & Degenaar, N. 2016, MNRAS, 461, 4400
- Ootes et al. (2019) Ootes, L. S., Vats, S., Page, D., et al. 2019, MNRAS, 487, 1447
- Ootes et al. (2018) Ootes, L. S., Wijnands, R., Page, D., & Degenaar, N. 2018, MNRAS, 477, 2900
- Page (2016) Page, D. 2016, NSCool: Neutron star cooling code, Astrophysics Source Code Library
- Page & Applegate (1992) Page, D. & Applegate, J. H. 1992, ApJ, 394, L17
- Page et al. (2006) Page, D., Geppert, U., & Weber, F. 2006, Nuclear Physics A, 777, 497
- Page et al. (2004) Page, D., Lattimer, J. M., Prakash, M., & Steiner, A. W. 2004, ApJS, 155, 623
- Page et al. (2014) Page, D., Lattimer, J. M., Prakash, M., & Steiner, A. W. 2014, in Novel Superfluids: Volume 2, ed. K. Bennemann & J. B. Ketterson (198 Madison Avenue, New York, NY 10016, United States of America: Oxford University Press), 505–579, arXiv:1302.6626
- Page & Reddy (2013) Page, D. & Reddy, S. 2013, Physical Review Letters, 111, 241102
- Parikh et al. (2017) Parikh, A. S., Homan, J., Wijnands, R., et al. 2017, ApJ, 851, L28
- Parikh et al. (2019) Parikh, A. S., Wijnands, R., Ootes, L. S., et al. 2019, A&A, 624, A84
- Potekhin et al. (1999) Potekhin, A. Y., Baiko, D. A., Haensel, P., & Yakovlev, D. G. 1999, A&A, 346, 345
- Potekhin et al. (1997) Potekhin, A. Y., Chabrier, G., & Yakovlev, D. G. 1997, A&A, 323, 415
- Ravenhall et al. (1983) Ravenhall, D. G., Pethick, C. J., & Wilson, J. R. 1983, Physical Review Letters, 50, 2066
- Rutledge et al. (2002) Rutledge, R. E., Bildsten, L., Brown, E. F., et al. 2002, ApJ, 580, 413
- Sato (1979) Sato, K. 1979, Progress of Theoretical Physics, 62, 957
- Schwenk et al. (2003) Schwenk, A., Friman, B., & Brown, G. E. 2003, Nuclear Physics A, 713, 191
- Shternin & Yakovlev (2006) Shternin, P. S. & Yakovlev, D. G. 2006, Phys. Rev. D, 74, 043004
- Shternin et al. (2007) Shternin, P. S., Yakovlev, D. G., Haensel, P., & Potekhin, A. Y. 2007, MNRAS, 382, L43
- Strohmayer & Bildsten (2006) Strohmayer, T. & Bildsten, L. 2006, in Compact Stellar X-Ray Sources, ed. W. H. G. Lewin & M. van der Klis (Cambridge University Press), 113 – 156
- Turlione et al. (2015) Turlione, A., Aguilera, D. N., & Pons, J. A. 2015, A&A, 577, A5
- Ushomirsky & Rutledge (2001) Ushomirsky, G. & Rutledge, R. E. 2001, MNRAS, 325, 1157
- Vlasyuk & Spiridonova (2017) Vlasyuk, V. V. & Spiridonova, O. I. 2017, The Astronomer’s Telegram, 10441
- Waterhouse et al. (2016) Waterhouse, A. C., Degenaar, N., Wijnands, R., et al. 2016, MNRAS, 456, 4001
- Wijnands et al. (2013) Wijnands, R., Degenaar, N., & Page, D. 2013, MNRAS, 432, 2366
- Wijnands et al. (2017) Wijnands, R., Degenaar, N., & Page, D. 2017, Journal of Astrophysics and Astronomy, 38, 49
- Wijnands et al. (2002) Wijnands, R., Guainazzi, M., van der Klis, M., & Méndez, M. 2002, ApJ, 573, L45
- Wijnands et al. (2004) Wijnands, R., Homan, J., Miller, J. M., & Lewin, W. H. G. 2004, ApJ, 606, L61
- Wijnands et al. (2006) Wijnands, R., in’t Zand, J. J. M., Rupen, M., et al. 2006, A&A, 449, 1117
- Woosley et al. (2004) Woosley, S. E., Heger, A., Cumming, A., et al. 2004, ApJS, 151, 75
- Yakovlev et al. (2001) Yakovlev, D. G., Kaminker, A. D., Gnedin, O. Y., & Haensel, P. 2001, Phys. Rep, 354, 1
- Yakovlev & Pethick (2004) Yakovlev, D. G. & Pethick, C. J. 2004, ARA&A, 42, 169
- Yakovlev & Urpin (1980) Yakovlev, D. G. & Urpin, V. A. 1980, Sov. Ast., 24, 303