A&A 450, 747758 (2006)
DOI: 10.1051/00046361:20054260
J. L. Zdunik^{1}  M. Bejger^{1}  P. Haensel^{1}  E. Gourgoulhon^{2}
1  N. Copernicus Astronomical
Center, Polish
Academy of Sciences, Bartycka 18, 00716 Warszawa, Poland
2  LUTH, UMR 8102 du CNRS, Observatoire de Paris, 92195 Meudon Cedex, France
Received 27 September 2005 / Accepted 6 January 2006
Abstract
Aims. We analyze potentially observable phenomena during spin evolution of isolated pulsars, such as back bending and corequakes resulting from instabilities, which could result from phase transitions in neutron star cores.
Methods. We study these aspects of spin evolution of isolated compact stars by means of analytical models of equations of state, for both constantpressure phase transitions and the transitions through the mixedphase region. We use highprecision 2D multidomain spectral code LORENE for the calculation of the evolutionary sequences of rotating neutron stars. This allows us to search the parameter space for possible instability regions, and possible changes in the stability character of rotating stars with phase transitions in their cores.
Results. We determine the conditions on the density jump in constantpressure phase transitions which leads to the back bending phenomena or to the existence of the unstable segments in the evolutionary sequences of spinning down isolated normal neutron stars. We formulate the conjectures concerning the existence of two disjoint families of nonrotating and rotating stationary configurations of neutron stars. To clarify the effect of rotation on the stability of neutron star we present the particular case of EOSs leading to marginal instability of static and rotating configurations: marginal instability point in nonrotating configurations continues to exist in all evolutionary spindown tracks. We discuss the fate of rotating stars entering the region of instability calculating the change in radius, energy release, and spinup associated with the corequake in rotating neutron star, triggered by the instability. The energy release is found to be very weakly dependent on the angular momentum of collapsing star.
Key words: dense matter  equation of state  pulsars: general  stars: neutron  stars: rotation
Particularly interesting method of searching for the phase transition in neutron star cores via pulsar timing was proposed by Glendenning et al. (1997). As a pulsar spins down, its central density increases, and for a certain density a new phase of matter can appear. In the case considered by Glendenning et al., the new phase consisted of quark matter. The authors suggested, that the softening of the EOS, induced by the formation of the new dense phase, leads to a temporary spinup era, the phenomenon called backbending. Originally, the name comes from nuclear physics, where the phenomenon of "backbending'' was observed in the systematics of the moment of inertia of excited states of rapidly rotating nuclei, see e.g., Ring & Schuck (1980). The calculations of Glendenning et al. were performed within the slowrotation approximation (Hartle & Thorne 1968; Hartle 1967), supplemented with additional relations resulting from accounting for the rotational stretching and framedragging effects (Weber & Glendenning 1992,1991).
Several other authors (Chubarian et al. 2000; Heiselberg & HjorthHensen 1998) carried out their calculations of the backbending phenomenon using the slow rotation approximation of Hartle. However, as shown by Salgado et al. (1994), the Hartle method, when compared with results of exact 2D numerical codes, breaks down for angular velocity close to the Keplerian one.
First calculations concerning backbending based on exact 2D code were performed by Cheng et al. (2002). These authors used the version of KEH code (Komatsu et al. 1989a,b), improved by Stergioulas & Friedman (1995) (see also references therein). In their work, Cheng et al. focused on the role of the crust for the very existence of the backbending. Indeed, as they show, even a slight change in the physical state of the crust (for example, a change in the crustcore transition pressure) may significantly affect the results. This shows that high precision is mandatory for reliable calculation of the backbending phenomenon.
Another important article containing results based on 2D computations with the Stergioulas & Friedman (1995) code, was published by Spyrou & Stergioulas (2002). They showed that the results obtained by Glendenning et al. (1997) are plagued by large numerical uncertainties. For example, the very same EOS as that used by Glendenning et al. (EOS from Table 9.2 of Glendenning 2000) did not yield the backbending phenomenon at all in Spyrou & Stergioulas (2002)! It became evident that the backbending problem is much subtler than previously considered, and that it requires careful handling as well as highprecision 2D computations. Also, Spyrou & Stergioulas pointed out some errors in previous papers on backbending. For example, the formula for braking index must be corrected by taking into account the rotational flattening of the star (Sect. 6 in Spyrou & Stergioulas 2002).
Most recently, Zdunik et al. (2004) showed that the backbending phenomenon can also occur for the EOSs different from the mixedphase one. They also pointed out importance of the stability with respect to the axisymmetric perturbations. The appearance of hyperons in the dense matter (Balberg & Gal 1997) softens some EOSs so much, that pulsars losing angular momentum actually spin up during a period of time. Paradoxically, during this spinup phase pulsars could lose a significant amount of their angular momentum.
It should be stressed that all previous works (except of Zdunik et al. 2004) considered backbending as a feature of the the dependence of the moment of inertia, I, on the rotation frequency, : I=I(f). As shown recently by Zdunik et al. (2004), this can easily lead to incorrect determination of the stability of rotating stars. Zdunik et al. (2004) pointed out that many cases claimed before to correspond to the backbending, actually cannot be realized in nature because of the instability with respect to the axisymmetric perturbations.
One of the aims of the present work is to determine reliably and precisely the stability regions on the back bending segments of the spin evolution tracks. In order to avoid any precision problem, and to investigate large and possibly complete parameter space, we will work with analytical EOSs of dense matter exhibiting a softening at supranuclear density. Two examples of softening by a phase transition will be considered. We will study EOSs with constantpressure phase transitions, characterized by a density jump obtained using the Maxwell construction. Second, we will use EOSs with phase transition extending over a finite pressure range in which two pure phases coexist forming a mixedphase state. Such EOSs are obtained for the firstorder phase transition between two pure phases by relaxing the condition of local electric charge neutrality and replacing it by less stringent condition of the global neutrality (Glendenning 1992). Mixedphase state can be realized provided the surface tension at the interface between the two pure phases is not too large.
The plan of the article is as follows. In Sect. 2 we present various types of analytical EOSs used in the calculation of the spin evolution tracks. Numerical methods used in exact 2D calculations are briefly presented in Sect. 3. Our numerical results are described in Sect. 4. We first describe the general criteria for the backbending and the stability for spinningdown stars. Then, the results for the EOSs with a mixedphase segment are studied in Sect. 4.1, and those with constant pressure phase transition with density jump are reviewed in Sect. 4.2. In Sect. 5 we describe a link between the existence of unstable segments in the families of static and rotating configurations of neutron stars. Change in neutron star parameters, accompanying transitions between two rotating configurations, triggered by instabilities of isolated rotating neutron stars, are studied in Sect. 6. Modifications in the pulsar timing and pulsar age evaluations, due to phase transitions in spinningdown isolated neutron stars are studied in Sect. 7. A summary of our results and their discussion is presented in final Sect. 8. Formulae referring to analytical models of EOSs with phase transitions are collected in the Appendix.
Figure 1: Two examples of EOSs employed in our calculations: constantpressure phase transition (dashed line, EOS with , the density jump , can be found in Table 2) and transition through a mixedphase state (solid line, thin dotted lines mark n_{1} and n_{2}; EOS MM of Table 1). Top panel: diagrams. Bottom panel: density profiles in stellar core, for a neutron star models.  
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Table 1: Main parameters of the EOSs with mixedphase segment. Below the mixedphase transition point n_{1} a polytropic EOS with is used. Mixed phase extends within , and is described by a polytrope with adiabatic index . Above the density n_{2} we assume pure quark matter with MIT bag model EOS . In all cases the dimensionless polytropic pressure coefficient Kwas equal 0.025 (see Appendix A for details). and denote the maximum allowable baryon and gravitational mass of the nonrotating star. The EOSs are labeled as follows: MSt produces a stable back bending, MUn  an unstable one, and MM produces a marginally stable case (for more details see the text).
Figure 2: Examples of the three EOSs with phase transition through the mixedphase state, considered in the present paper. The parameters of the EOSs are given in Table 1.  
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Figure 3: Total angular momentum versus rotation frequency f ( left panels), and moment of inertia versus f ( right panels), for EOSs MSt and MUn. The stability criterion is easily applied to left panels. It is clear that for the MSt EOS back bending feature is not associated with an instability, with all configurations being stable. On the contrary, the MUn EOS produces back bending with a large segment of unstable configurations. Simultaneously, the I(f) curves for both EOSs are very similar and apparently show very similar backbending shapes.  
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In the present paper we will show examples of the EOS with constantpressure (densityjump) phase transitions and mixedphase transitions for which one of the two situations occurs:
Figure 4: Upper panel: stellar baryon mass vs radius for fixed frequency (dashed lines) and fixed total angular momentum (solid lines), for the model MM of mixedphase EOS. Bottom panel: stellar angular momentum as a function of the rotational frequency for fixed baryon mass (indicated as a label, in solar masses) for the same MM EOS. This EOS corresponds to the marginal case from the point of view of stability  the curves and have flat horizontal regions. The regions of backbending are drawn by thick lines.  
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Figure 5: The same as in Fig. 4 but for the MSt EOS. Phase transition does not result in the stability loss  all configurations are stable.  
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Figure 6: The same as in Fig. 4 but for MUn EOS. Upper panel: Phase transition results in the stability loss, i.e., there exists a region where decreases with increasing at fixed J (marked by dotted lines). Lower panel: plots in the Jf plane, where on the dotted segments J increases as increases (unstable configurations). Both features indicate instability with respect to the axisymmetric perturbations.  
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In Figs. 46 we present the results obtained for three different choices of the EOS's parameters describing mixed phase phase transitions (Table 1).
Model MM has been chosen to be a marginal one  there exists a
region where the curve
(or
)
is
locally horizontal, which means the marginal stability. The
function
for nonrotating star (f=0) has an
inflection point slightly above the configuration where phase
transition to the mixed phase occurs. At a fixed f, the
condition for an inflection point reads,
Model MSt produces a set of stable configurations in the region of phase transition (Fig. 5). Here, is an increasing function of for nonrotating configurations, as well as along all rotating sequences of configurations at fixed J, terminating at the global maximum mass configurations. The backbending phenomenon is in this case limited to the baryon masses larger than that for which the curve (or equivalently ) at fixed f has a flat, horizontal region (strictly speaking, where the Eq. (1) is fulfilled). The numerical values for the MSt model are and f>400 Hz (see for example the bottom panel of Fig. 5, where the curve for has almost vertical part).
For the MUn model there exists a region for which configurations are unstable, i.e. the baryon mass is decreasing function of at fixed J (Fig. 6). In some sense this instability is not very strong  the difference between the maximum (local) and minimum mass is of the order of 0.3%. However this feature (existence of instability region) is characteristic to all rotational frequencies  for all values of angular momentum (J) fixed along the curve, the baryon mass has local maximum and local minimum connected by an unstable sequence of stellar configurations.
For nonrotating configurations, the reaction of the star to a constant pressure (first order) phase transition has been studied in detail in the second half of 1980s (see Zdunik et al. 1987, and references therein). The appearance of a new, dense phase in the center of the star results in the change of the derivatives of the global stellar parameters with respect to (see the formula B6 in the appendix of an article by Zdunik et al. 1987). Two important dimensionless parameters are: fractional density jump and the relativistic parameter . There exists a critical value of , , such that for configurations with an infinitesimally small Bphase core are unstable with respect to collapse into a new configuration with a large core of the dense phase. Putting it differently, a phase transition with destabilizes the star at central pressure at which the phase transition occurs. It should be stressed that while guarantees stability of smallcore configurations, it does not assure the stability of configurations with a finite, or  in an extreme case  a large core. In such a case the instability would result from the softness of the phase somewhat above and not directly from an overcritical . In other words the compressibility of a matter leads to the larger mean density in the core than the value at the phase boundary. The response of the whole star to the appearance of the dense core built of the Bphase of the matter is determined by the mass and radius of this core (strictly speaking, this statement is true for nonrotating configurations, see Zdunik et al. (1987); for rotating ones also rotation rate and resulting oblateness play role). As a result even if the first order phase transition can lead to the unstable configurations for finite size of the core. As a result for the given model of the matter in the phases and there exist the maximum value of density jump for which all configurations below maximum mass are stable. Of course and the difference between and is larger for softer EOS in the phase .
Table 2: Selected sets of EOSs with a constant pressure phase transition. For all cases, and (see the text and Appendix A.1).
Numerical results for a collection of sets of EOSs with constant pressure phase transition are collected in Table 2. The parameters presented in this table correspond to the onset of back bending, i.e., the rotational frequency and baryon mass for which the curve or starts to have a flat region. More precisely, at these values of frequency and mass, an inflection point appears in the curves under consideration. We also included parameters of those EOSs for which all nonrotating stars with are stable. This means that for such EOSs the curve for static configurations increases monotonically up to . The parameter gives then the maximum value of the density jump for a fixed set of other EOS parameters (adiabatic indices and , number density threshold ) for which this property of neutron stars is valid; in other words corresponds to the "marginally stable'' case. Increasing implies increasing softening of the EOS by the phase transition. If , the phase transition leads to the existence of an unstable branch of the nonrotating stellar configurations. This unstable branch separates stable family of neutron stars with Aphase cores from a second family of superdense neutron stars with Bphase cores: these are two distinct neutronstar families. It should be mentioned that this feature (existence of the unstable region) does not depend on rotation  the unstable branches exist also for rotating configurations (strictly speaking for any value of a total angular momentum of the star J there exist a region with ). We have tested this feature (existence of or the lack of unstable regions) for very small departures from marginally stable case ( ). From numerical results it follows that if all rotating configurations are stable (before loosing stability at maximum mass point) and if we have two branches of stable configuration for rotating stars (for any J).
Picking up the onset parameters is visualized in
Fig. 7 where we display
for one of the EOSs from Table 2. The curves are
plotted for three frequencies, with middle one corresponding to the
back bending onset,
.
Last column of Table 2
gives the maximum allowable baryon mass for static configurations,
.
We restrict ourselves to
back bending for the normal (non
supramassive) stars, which appears during the spindown evolution which terminates
eventually by a nonrotating stable configuration. The dependence between the
backbending onset parameters 
,
corresponding baryon mass, and the
intrinsic parameters of the EOS 
the density jump ,
as well as the "departure'' from the
critical configuration (
)
is presented in Fig. 8. The three families of curves visualize
the data from Table 2 (solid lines for
,
dotted for
and dashed for
).
The value
defines the onset of backbending at the limit f=0;
in this case the backbending phenomenon is present for any rotational
frequency.
As it can be seen on the left panel, the onset frequency
depends very weakly on the EOS in the dense core  the main parameter describing the
reaction of the star to the appearance of this phase transition is the density jump.
The right panel presents the same data not normalized with respect to the
maximum density jump
 the results can be very well approximated by
the dependence
.
These two plots can be treated as
a slice through the parameter space to search for regions of the backbending appearance

in the right panel, the backbending is present, for a particular model, above
its curve.
Figure 7: The definition of the frequency of the onset of backbendingphenomenon, , and corresponding mass, . In the presented example Hz and .  
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Figure 8: The onset frequency of the backbending phenomenon as a function of the departure of from the maximum density jump ( , left panel), and density jump as a function of ( right panel) for models presented in Table 2 (solid lines for , dotted for and dashed for ).  
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Rotation can influence stability of a star of a given baryon mass. In particular, stars with cannot exist without rotation and collapse into black hole as soon as their rotation frequency falls below a certain minimum value. Here, however, we will restrict to a different problem of stability, which will concern the normal configurations only.
In what follows we will use the term "stability'' in a restricted sense. Namely, by stability (instability) of an equilibrium configuration we will mean stability (instability) with respect to radial perturbations in the nonrotating case, and with respect to axisymmetric perturbations for rotating configurations.
We studied a very large set of EOSs with phase transitions at
constant pressure, as well as those with transition through a
mixed phase state. We then produced static sequences and normal
rotating sequences for these EOSs. Our calculations were very
precise, because we used analytic forms of the EOSs. The results
for both constant pressure phase transitions, and those proceeding
through mixed phase, turned out to be qualitatively the same. In
all cases, if nonrotating configuration were stable
(monotonically increasing
and
), then for
any value of the total angular momentum J the functions
and
were monotonically
increasing, too. Thus, when all nonrotating configuration with
were stable (with respect to radial
perturbations), all normal rotating configurations were stable too
(with respect to axisymmetric perturbations). On the other hand,
if for nonrotating stars there existed a region with decreasing
and
,
even extremely small one with a
very shallow minimum, then an unstable region persisted within the
rotating configurations, at each value of J. These two cases are
illustrated in the Figs. 5 and 6.
Figure 9: The mass M of the star as a function of central baryon number density for the MM model. Dashed lines  fixed rotational frequency. Solid lines  fixed total angular momentum J. The vertical lines correspond to the densities of transitions from the nuclear to mixed phase and from the mixed phase to the pure denser phase (here  quark matter)  see Table 1.  
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We studied also the case of marginally stable EOS. An inflection point, witnessing marginal stability, present in the or curves for nonrotating stars continued to exist in the or curves for normal rotating stars (Figs. 4, 9).
The analysis of numerical results leads us to an interesting conclusion. Namely, for an EOS with a phase transition (constant pressure one or through mixedphase state), rotation neither stabilizes nor destabilizes normal sequences of stationary configurations based on this EOS. We define a family of configurations as a compact set (in mathematical sense) of configurations. Similarly, an EOS leading to a marginally stable point for nonrotating stars, produces also spindown evolution tracks with a marginally stable point. Our result can be formulated as three conjectures:
Figure 10: Evolution of an isolated pulsar loosing angular momentum, after it reaches the instability region in Jf plane and then collapses. Arrows lead from unstable configuration to a collapsed stable one, with the same baryon mass and angular momentum. Dotted lines  unstable configurations.  
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Figure 11: Changes of stellar parameters of a rotating solitary neutron star, due to a collapse which occurs after a pulsar loosing angularmomentum reaches an unstable configuration.  
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Figure 12: Total gravitational mass of the star as a function of its angular momentum, for fixed baryon number of the star for our EOS model MUn. The central density is increasing along this curve as marked by the arrows. The upper segment (dotted) corresponds to the unstable configurations. Two cusps reflect strict property that the mass and angular momentum have simultaneous extrema along the path with fixed baryon number.  
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As was shown on many occasions in the preceding sections, an isolated neutron star, loosing its angular momentum, moves down along the line of fixed in the J(f) plane, and can at some moment reach the instability point (i) (the minimum of J(f)at fixed in Fig. 10). The subsequent behavior of the star cannot be described by our stationary rigidly rotating model. In real world, the star has to collapse, rearranging the angular momentum distribution in its interior. What we can do, is determining the final stationary state (f), which by construction will be a stable rigidly rotating neutron star. We assume, that the transition conserves the baryon number of the star and is sufficiently rapid so that the angular momentum loss can be neglected. Therefore, the final stable configuration will have the same and J as the unstable initial one, , . The difference between the parameters of these configurations (mass, equatorial radius, moment of inertia) gives us the energy release, and changes in equatorial radius and rotation frequency, due to the collapse implied by the instability.
Examples are shown in Figs. 1012 for our EOS model MUn. As the star becomes more compact, collapse is accompanied by the decrease of the equatorial radius (by a few kilometers) and by a significant spin up. For a given EOS, the changes in radius, energy release, and spinup are function of the angular momentum at the instability point: , , . As we see in Fig. 11, the energy release depends rather weakly on the rotation of the unstable configuration (i.e., on the value of ). It should be mentioned that our MUn model is only an example of the EOS resulting in the instability region within the hydrostatic equilibria. For EOSs with a weaker phase transition this instability region would be narrower and the changes of stellar parameters in the collapse would be smaller. However an approximate constancy of the energy release (i.e., its very weak dependence on ) seems to be a generic property of rotating neutron stars undergoing a collapse due to a first order phase transition.
In order to discuss in more detail the energy release during collapses , we plotted in Fig. 12 the gravitational mass of the star, M, as a function of angular momentum, J, at fixed baryon mass: . Consider an initial configuration C_{1}. As the star looses angular momentum, it moves down along line a, and reaches eventually the cusp (corresponding to the value of ). To continue moving on the dotted segment , the star would have to gain angular momentum and energy!
As we already mentioned, the evolution of the star beyond the instability point cannot be described by our hydrostationary model. The star can only collapse to the final configuration , with the same values of and J, i.e. along vertical arrow in Fig. 12. Then, it evolves down the line b.
We notice a very special role played by the point at which line a and b cross. This is a degeneracy point, which corresponds to two very different configurations of the same , M, and J. However, transitions between these two configuration are prevented by the huge energy barrier.
The existence of the sharp cusps at and C_{2} on the track is a very stringent test of the precision of the numerical code: it means that mass and angular momentum extrema (for fixed baryon mass) are reached exactly at the same point. This property follows from the general relativistic relation (Bardeen 1972). Here, , and is the "injection energy per unit mass''. This relation has to be strictly fulfilled by the stationary configurations. A graph, analogous to Fig. 12, can be plotted in the plane for configurations with fixed angular momentum J. Also in this case, the existence of sharp cusps proves the correctness of the numerical code.
It has been already pointed out by Spyrou & Stergioulas (2002), that the back bending phenomenon, resulting from the growth of a densephase core, can lead to significant difference between the actual pulsar age, , and that inferred from the measurements of the period P and period derivative, , and denoted . The calculation of is based on quite strong assumptions. Firstly, pulsar kinetic energy loss due to radiation is given by the magnetic dipole formula. Secondly, nonrelativistic approximation is used, with pulsar kinetic energy given by , where and pulsar moment of inertia is constant, independent of .
Following Spyrou & Stergioulas (2002), we will use general relativistic
notion of total pulsar energy, Mc^{2}. Then the pulsar energy
balance is
Let us consider the increase of stellar energy due to a spin up
to frequency f at constant .
In general relativity, the
increase is given by
.
In the standard model, we neglect the effect of f on stellar
structure, so that
.
This is a good approximation when rotation is
slow and EOS is smooth (no phase transition).
Equation (2) can be then rewritten as
In the case when angular momentum loss leads to
a phase transition at the stellar
center, the situation is much more complicated, because
of the strong fdependence of the pulsar structure in
the vicinity of the phase transition. This difference
is illustrated in Fig. 13, where we plotted
the quantity
resulting from
our calculations, and compared it with results
given by standard nonrelativistic model with constant I.
Figure 13: The relative massenergy increase due to rotation of the star at fixed baryon mass, , for the EOS with a phase transition (MM model), is shown using solid line. Three solid lines are labeled by the gravitational mass of the nonrotating configuration (in solar masses). Dotted lines correspond to the curves calculated for the standard model, Eq. (3), with .  
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Figure 14: The evolution of the pulsar period when the energy loss is described by the magnetic dipole braking with . Solid curves  results for our model MM with different values of baryon mass. Dotted lines correspond to the standard model, Eq. (3). The unit of time (horizontal axis) is where [cgs] is parameter entering Eq. (2).  
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We limited ourselves to the case of normal rotating configurations, which are connected with the nonrotating ones by the angular momentum loss. We considered two types of instabilities which bound the sets of stationary configurations: mass shedding and instability with respect to the axisymmetric perturbations. The EOSs split into two sets: those producing a single family of stable stationary configurations (static and rotating) of neutron stars and those producing two disjoint families of stable stationary configurations. Conjectures concerning normal configurations based on EOSs with a phase transition has been formulated. If an EOS yields two disjoint families of static configurations containing "twin neutron stars'' of the same baryon mass but different radius, then also stable normal rotating configurations form two disjoint regions in the mass  equatorial radius plane, which contain "twin neutron stars'' of the same baryon mass but of different compactness.
Very often, neutron stars are called a second family of compact stars, the first family being composed of white dwarfs. Therefore, our conjecture means that an EOS which produces a third (disjoint) family of static compact stars, produces also a third (disjoint) family of rotating normal stars.
We have also shown the existence of a very special class of "fine tuned EOSs'' with phase transitions which produce marginally stable stationary configurations of normal neutron stars, which form a boundary separating stable stationary configurations (a line in the massequatorial radius plane).
Conditions on the density jump in constantpressure phase transitions were derived, under which their presence in the EOS produces the back bending phenomenon in the spin down evolution.
The case when a spinning down normal neutron star reaches an unstable configuration was studied in detail. The instability leads to neutron star collapse, associated with an energy release in a "corequake'', decrease of radius, increase of central density, and spin up of the star. We have shown that the energy release associated with such a "corequake'' depends rather weakly on the initial rotation frequency at the instability point. In our examples, energy release was of the order of a few times 10^{50} erg.
In the present paper we put accent on the numerical precision and mathematical strictness. We hope that in this way we prepared ground for further studies of the impact of the phase transitions in dense matter on the structure, evolution, and dynamics of rotating neutron stars. These further studies will be performed using realistic EOSs available in the literature and taking into account important microscopic aspects of the phase transitions. The kinetics of the phase transition coupled with stellar spindown, and the ensuing neutron star corequake are now being studied. These topics will be the subject of our subsequent papers.
Acknowledgements
This work was partially supported by the Polish MNiI grant no. 1P03D00827 and by the PAN/CNRS LEA AstroPF.
Dense matter is strongly
degenerate, so that the T=0 approximation is valid.
First Law of Thermodynamics implies then expression for energy
per baryon
The energy density
is thus given by
Figure A.1: Examples of phase transitions considered in the text; constant pressure phase transition ( left), and the phase transition through the mixedphase state ( right). 
From the continuity of pressure, baryon density, and energy density
at the
transition point, we get