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A recent recalibration of the Geneva Copenhagen Survey (GCS) using the infrared flux method finds no discrepancy with solar age stars (Casagrandeetal.2011). but even modern. studies contirm that the age-metallicity relationships (AMRs) of field and solar neighborhood stars are characterized by higher dispersions than expected (Edvardssonetal.1993:Nordstróm2004:Holmbergetal.2007:Casagrande 2011..
A recent recalibration of the Geneva Copenhagen Survey (GCS) using the infrared flux method finds no discrepancy with solar age stars \citep{Casagrande11}, but even modern studies confirm that the age-metallicity relationships (AMRs) of field and solar neighborhood stars are characterized by higher dispersions than expected \citep{Edvardsson93,Nordstrom04, Holmberg07, Casagrande11}.
. In. addition. simple chemical evolution models that divide the Galaxy into independently evolving concentric annuli predict many more low metallicity G-dwarfs in our region of the disk compared to those observed. a discrepancy known as “the local G-dwarf problem” (vandenBergh1962:Schmidt1963).
In addition, simple chemical evolution models that divide the Galaxy into independently evolving concentric annuli predict many more low metallicity G-dwarfs in our region of the disk compared to those observed, a discrepancy known as “the local G-dwarf problem” \citep[][]{vandenbergh62, Schmidt63}.
Evidence seemingly in contradiction to standard galaxy chemical evolution theory is not limited to our own Galaxy.
Evidence seemingly in contradiction to standard galaxy chemical evolution theory is not limited to our own Galaxy.
Metallicity gradients in disk galaxies are shallower than predicted by classical models (e.g..Magrinietal.
Metallicity gradients in disk galaxies are shallower than predicted by classical models \citep[\eg,][]{Magrini07}.
2007)... Ferguson&Johnson(2001). and Fergusonetal.(2007) find unexpectedly old stellar populations on nearly circular orbits in the outskirts of M31 and M33. respectively.
\citet{Ferguson01} and \citet{Ferguson07} find unexpectedly old stellar populations on nearly circular orbits in the outskirts of M31 and M33, respectively.
The outermost regions of NGC300 and NGC7739 show flattened or positive abundance gradients with radius (Vlajiéetal.2009.2011).
The outermost regions of NGC300 and NGC7739 show flattened or positive abundance gradients with radius \citep{Vlajic09,Vlajic11}.
. These perplexing observations cannot be readily explained within the contines of classic galaxy formation models.
These perplexing observations cannot be readily explained within the confines of classic galaxy formation models.
A natural explanation for the observational challenges above arises if the present day radii of many stars could be signiticantly different from their birth radii.
A natural explanation for the observational challenges above arises if the present day radii of many stars could be significantly different from their birth radii.
One difficulty in establishing radial migration as a common phenomenon. from a dynamical standpoint.
One difficulty in establishing radial migration as a common phenomenon, from a dynamical standpoint,
that incorporation of type-I migration without auy conditions causes severe inconsistency with observed data of extrasolar planets aud our Solar System.
that incorporation of type-I migration without any conditions causes severe inconsistency with observed data of extrasolar planets and our Solar System.
If we rely ou type-I migration moclel. we ueed to clarify the condition for the occurrence of type-[ inigration at the same time.
If we rely on type-I migration model, we need to clarify the condition for the occurrence of type-I migration at the same time.
The high deusity of the close-in ext‘asolar elaut planet HD119026b recently discovered by Cialleuges theories of plaiet. formation.
The high density of the close-in extrasolar giant planet HD149026b recently discovered by \citet{Sato05} challenges theories of planet formation.
Iu this paper. we have attempted to «lerive robust constraiits on the planet's compositio1 auc iufer possible routes to explain its formation.
In this paper, we have attempted to derive robust constraints on the planet's composition and infer possible routes to explain its formation.
We have fi‘st sinlated the evolution of HD119026) more exteusively than previous w«kers (Satoetal.2005xFortiey2006) aid corfirmed that the planet coutaius a sustantla alnot ol heavy eeluerts.
We have first simulated the evolution of HD149026b more extensively than previous workers \citep{Sato05,Fortney06} and confirmed that the planet contains a substantial amount of heavy elements.
Preerred values of t1 elota| mass of heavy elements are 50-80 ÀLl (secTOl 2.1 which is €OUSISent wihi the previous calculaious.
Preferred values of the total mass of heavy elements are 50–80 $\mearth$ (section \ref{sec:standard_model}) ), which is consistent with the previous calculations.
We showed that the results a| unciauged or heavy eleijens locatec Lin the central core. or distributed iuside the euvelope. provided they remain deeper thiui the exterial radiative zone.
We showed that the results are unchanged for heavy elements located in the central core, or distributed inside the envelope, provided they remain deeper than the external radiative zone.
In the event of a significant enrichment of tlie outer layers. slightly hieher values of heavy elements coutent are possible (section 2.3)).
In the event of a significant enrichment of the outer layers, slightly higher values of heavy elements content are possible (section \ref{sec:core_or_envelope}) ).
Iu. order ο €lerive luΠΕΙ values of the mass of heavy eements. we have explored tle possibility that tle planet Was slOpec ina relatively cold euviroumeu for sotje time before mierating near to tlie pinet.
In order to derive minimum values of the mass of heavy elements, we have explored the possibility that the planet was stored in a relatively cold environment for some time before migrating near to the planet.
This strict miniΕΜΠ is ~35M.. but is regarced as unlikely because it reques a late migrat1 and no yeieating of the planet by tidal circulizajon (section 2.3)).
This strict minimum is $\sim 35 \mearth$, but is regarded as unlikely because it requires a late migration and no reheating of the planet by tidal circulization (section \ref{sec:cold_storage_hypothesis}) ).
We have then investigated the possillity of subcritical core accreion as envisioned. OotU rantS d Neptune to account for the small eivelope lass as well as the arge core mass.
We have then investigated the possibility of subcritical core accretion as envisioned for Uranus and Neptune to account for the small envelope mass as well as the large core mass.
D1 incipe arge Cor eo. 90-80 M Ca1 be formed by sul(critical Core accretion.
In principle a large core of 50–80 $\mearth$ can be formed by subcritical core accretion.
However we |ave found very unlikelv for at least two reasons: (i) A σιberitical core accreion results in a ratio of te ο bass d Ol1ο total mass a)ove ~OQ.T (section 3.1)). whereas ot revolutio1 calculaious slowed ‘th a hiei ratio to be possible in a very limited range of paraneers (see Table l1) (ii) Te )critical forration of a 50-80 M core requires au extremely massive or meal-rich clisk with dust ‘face deisity 30-50 times tlie vates obtained for the mai ninilass solar uebula (section :3.
However we have found it very unlikely for at least two reasons: (i) A subcritical core accretion results in a ratio of the core mass to the total mass above $\sim 0.7$ (section \ref{sec:critical_mass}) ), whereas our evolution calculations showed such a high ratio to be possible in a very limited range of parameters (see Table \ref{tab:constraints}) ); (ii) The subcritical formation of a 50–80 $\mearth$ core requires an extremely massive or metal-rich disk with dust surface density 30-50 times the values obtained for the minimum mass solar nebula (section \ref{sec:subcritical}) ).
2 'easonably jassive and/or metaI-yich disk eau fori cores of at most ~30K [ar [rom the parent a5
A reasonably massive and/or metal-rich disk can form cores of at most $\sim 30 \mearth$ far from the parent star.
Those facts require us to coSidler (i) the migration ol the planet. (ii) the supy of 1eavy ements to tle planet during or alte "the gas accretion pliase. aud (11) a limited suppvy of cis&e Bas0 oss of the envelope gas to acco1 for the properties o. HD119026b.
Those facts require us to consider (i) the migration of the planet, (ii) the supply of heavy elements to the planet during or after the gas accretion phase, and (iii) a limited supply of disk gas or loss of the envelope gas to account for the properties of HD149026b.
h seclior L.1.. we have discusse how the heavy elements cau be delivered to the planet dWing Or alter the gas accretion pliase according to hese scenarios.
In section \ref{sec:Z_supply}, we have discussed how the heavy elements can be delivered to the planet during or after the gas accretion phase according to these scenarios.
An efficieit delivery. durius tle gas accretlon piase ueeds to be re-invesigated in much more details. becaise the shepherdiug eucs to prevent he planet from accreting plauetesinals (section [.1.1)).
An efficient delivery during the gas accretion phase needs to be re-investigated in much more details, because the shepherding tends to prevent the planet from accreting planetesimals (section \ref{sec:concurrent_Z_supply}) ).
Ou he other hand. scatteriug oL plajetes]tals/planets by one or several outer giant. planets was shown to lead to au elicient accrelion by a close-in giaut. planet. and is a p'omisiug explanation for tje formation of meta|-rich
On the other hand, scattering of planetesimals/planets by one or several outer giant planets was shown to lead to an efficient accretion by a close-in giant planet, and is a promising explanation for the formation of metal-rich
Mi.=Vie/G r=GAL,⋉/2cz"EE,La Myr.)=1. (Vs(rs)>1). (Maier)1). CAEGn)uw1) ParkerUs laa-b: see1909)). Parker ‘acceler (Lamers&Wintersetal.2000) n(f-:|: ACIE=0) Lee) 7. f. Parker {1 p,Cr). η).ei(r) Vy pw. àV dp ®,i M,1 r,1 Ἐν.
$\mach_w=V_w/\cs$ $r_c=GM_\ast/2\cs^2$ $\mach_w(r_c)=1$ $\mach_w(r_c)>1$ $\mach_w(r_c)=1$ $\mach_w(r_c)<1$ \citeauthor{par58}' \ref{fig:init}a \citealp[see also Figure~3.1 of][]{lam99}) \citeauthor{par58}' \citep{lam99,win00} $f=1$ $M_\ast^{\rm eff}=0$ \ref{fig:init}c $r_\ast$ $f$ \citeauthor{par58}' $f=1$ $\rho_w(r)$ $V_w(r)$$\cs(r)$ $\vct{V_w}$ $\rho_w$ $\delta\vct{V}$ $\delta\rho$ $\Phi_p$ $M_p$ $r_p$ $V_p$
tidal torques and infall episodes.
tidal torques and infall episodes.
These questions can likely be addressed: using the Millennium Simulation which has already been used to explore properties of BC'Gs (seec.g.?) Recent surveys such as the SDSS Stripe S2 and future surveys such as LSS (?.LSSTScience.Book)— and PanSTARRS (7?) with deeper photometry will allow us to probe clusters at higher redshifts to study the evolution of alignment over a larger range of cosmic time.
These questions can likely be addressed using the Millennium Simulation which has already been used to explore properties of BCGs \citep[see e.g.][]{2009ApJ...696.1094R} Recent surveys such as the SDSS Stripe 82 and future surveys such as LSST \citep[][LSST Science Book]{2009arXiv0912.0201L} and PanSTARRS \citep{2002SPIE.4836..154K} with deeper photometry will allow us to probe clusters at higher redshifts to study the evolution of alignment over a larger range of cosmic time.
Deeper surveys will also increase the number of galaxies in low-redshift clusters. reducing Poisson noise in our samples.
Deeper surveys will also increase the number of galaxies in low-redshift clusters, reducing Poisson noise in our samples.
This will allow us to study cluster shapes and. alignments for galaxies of dillerent luminosities. a diagnostic of dillerent accretion and dynamical histories.
This will allow us to study cluster shapes and alignments for galaxies of different luminosities, a diagnostic of different accretion and dynamical histories.
Simulations addressing the connection between BCC dominance and alignment will be important in discovering which mechanisms cause the distinct. physical. properties of the BCGs and their relation to their parent cluster.
Simulations addressing the connection between BCG dominance and alignment will be important in discovering which mechanisms cause the distinct physical properties of the BCGs and their relation to their parent cluster.
AMINO is funded. by the Gates Cambridge Trust. the Isaac Newton Studentship fund and the Science and ‘Technology Facilities Council (SPEC).
MNO is funded by the Gates Cambridge Trust, the Isaac Newton Studentship fund and the Science and Technology Facilities Council (STFC).
ALAS was supported in part by NSE grant .AST-0707266.
MAS was supported in part by NSF grant AST-0707266.
. We thank James LE. Gunn and Michael D. Claceders for valuable discussions aud eedback and the anonymous referee for help in clarifying he text.
We thank James E. Gunn and Michael D. Gladders for valuable discussions and feedback and the anonymous referee for help in clarifying the text.
Funding for the SDSS and SDSS-LE has been provided w the Alfred. D. Sloan Foundation. the Participating Institutions. the National Science. Foundation. the U.S. Department of Energy. the National Aeronautics and Space Administration. the Japanese Monbukagakusho. the Max λαο Society. and the Higher Education Funding Council or England.
Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England.
The SDSS Web Site is http:wav.sclss.0re
The SDSS Web Site is http://www.sdss.org/.
is (he consequence of an increase in (he radial component of the pressure gradients of the eas,
is the consequence of an increase in the radial component of the pressure gradients of the gas.
Figure 8 shows (is concept in more detail.
Figure 8 shows this concept in more detail.
In (his figure. we have plotted the density of the gas as a [unction of r on a plane parallel to the midplane at 2=0.2 AU.
In this figure, we have plotted the density of the gas as a function of $r$ on a plane parallel to the midplane at $z=0.2$ AU.
As a result of increasing the temperature of (he gas from 50 Ix to 1000 Ix. the maximum value of the gas density on this plane is increased by (three orders of magnitude.
As a result of increasing the temperature of the gas from 50 K to 1000 K, the maximum value of the gas density on this plane is increased by three orders of magnitude.
Such an increase implies an increase in the radial component of the pressure gradients of the eas which in turn results in more rapid radial migration.
Such an increase implies an increase in the radial component of the pressure gradients of the gas which in turn results in more rapid radial migration.
From equation (5). the radial position of the maxinuun of the gas density al a certain j1ejeht z approaches smaller values bv increasing the gas temperature.
From equation (5), the radial position of the maximum of the gas density at a certain height $z$ approaches smaller values by increasing the gas temperature.
Figure ὃ shows (his or z —0.2 AU.
Figure 8 shows this for $z=$ 0.2 AU.
For an object at this height and at an initial radial position of r=2 AU. such a decrease in the radius of the local density enhancement causes the distance of the initial outward migration of the object to become smaller (Figure 7).
For an object at this height and at an initial radial position of $r=2$ AU, such a decrease in the radius of the local density enhancement causes the distance of the initial outward migration of the object to become smaller (Figure 7).
For the temperature of 1000 Ix. the radial location of the maximum gas density αἱ 2=0.2 AU becomes smaller than 2 AU and as a result. an object at (2.0.2) AU does not undergo an initial outward migration.
For the temperature of 1000 K, the radial location of the maximum gas density at $z=0.2$ AU becomes smaller than 2 AU and as a result, an object at (2,0.2) AU does not undergo an initial outward migration.
While the combined effect of pressure gradients and (he drag force of the gas causes solids to racially migrate toward the location of local densitv enhancement. the vertical component of the gravitational force of the central star attracts solids toward the midplane of the nebula.
While the combined effect of pressure gradients and the drag force of the gas causes solids to radially migrate toward the location of local density enhancement, the vertical component of the gravitational force of the central star attracts solids toward the midplane of the nebula.
Figure 9 shows the vertical motion of solids with different sizes.
Figure 9 shows the vertical motion of solids with different sizes.
The dependence ol the rate of vertical migration on the size of an object can clearly be seen [rom this figure.
The dependence of the rate of vertical migration on the size of an object can clearly be seen from this figure.
As shown here. except for when the objects are in the vicinitv of the midplane. the rates of (heir vertical migrations increase by increasing (their radii.
As shown here, except for when the objects are in the vicinity of the midplane, the rates of their vertical migrations increase by increasing their radii.
Ili (he vicinity of the midplane. however. while I0 centimeter-sized and smaller objects continue their smooth descent. the 1 meter-sized object undergoes an overshoot.
In the vicinity of the midplane, however, while 10 centimeter-sized and smaller objects continue their smooth descent, the 1 meter-sized object undergoes an overshoot.
Such an overshoot aud its corresponding damped oscillatory motion are more pronounced when (he density of (he object is increased.
Such an overshoot and its corresponding damped oscillatory motion are more pronounced when the density of the object is increased.
Figure 10 shows the vertical migrations of a 1 centimeter-sized and a 1 meter-sized object for different values of their densities.
Figure 10 shows the vertical migrations of a 1 centimeter-sized and a 1 meter-sized object for different values of their densities.
As shown here. the rates of vertical descent of both objects increase bv increasing their densities.
As shown here, the rates of vertical descent of both objects increase by increasing their densities.
Once in (he vicinity of the midplane. the meter-sized object undergoes a damped oscillatory motion whose amplitude increases by increasing (he solid's density.
Once in the vicinity of the midplane, the meter-sized object undergoes a damped oscillatory motion whose amplitude increases by increasing the solid's density.
Figure 11 shows the thiree-dimensional paths of these (wo objects.
Figure 11 shows the three-dimensional paths of these two objects.
The clvnamical behavior ofan object along the z-axis and in the vicinity of the midplane can be explainecl bv studying equation (13) in more detail.
The dynamical behavior of an object along the $z$ -axis and in the vicinity of the midplane can be explained by studying equation (18) in more detail.
In this equation. the crag coellicient C5 is Che factor that determines the funetional form of D.
In this equation, the drag coefficient $C_{\rm D}$ is the factor that determines the functional form of ${{\dot {\hat P}}_z}$.
substituting for the drag force of the eas lrom equation (8) and lor the z-component of the relative velocity of
Substituting for the drag force of the gas from equation (8) and for the $z$ -component of the relative velocity of
absorption).
absorption).
The coluun deusitv Nyy is between 1.0 and 1.1107? ?.
The column density $N_{\rm H}$ is between 1.0 and $\times10^{22}$ $^{-2}$.
This range is consistent with the value resulting from the interstellar reddening to Terzan 6: for Epvy=2.21 (Barbuy et al.
This range is consistent with the value resulting from the interstellar reddening to Terzan 6: for $E_{\rm B-V}=2.24$ (Barbuy et al.
1997) with an estimated error of 0.1. ely=6.9L40.31 and Ny=(7940.1)«10214 3 (according to the conversion of Av to Nyy by Predehl Scluuitt 1995).
1997) with an estimated error of 0.1, $A_{\rm V}=6.94\pm0.31$ and $N_{\rm H}=(1.79\pm0.1)\times10^{21}A_{\rm V}=(1.2\pm0.1)\times10^{22}$ $^{-2}$ (according to the conversion of $A_{\rm V}$ to $N_{\rm H}$ by Predehl Schmitt 1995).
The xoad component at 1.6-2.0 keV. which we moceled by black body radiation. affects the determination of Nyy somewha but we estimate that this is limited to 0.2«1072 cu?.
The broad component at 1.6-2.0 keV, which we modeled by black body radiation, affects the determination of $N_{\rm H}$ somewhat but we estimate that this is limited to $0.2\times10^{22}$ $^{-2}$.
Ouly the LECS aud MECS provide data curing the eclipse that are of sufficient quality to allow a meaninetul analvsis.
Only the LECS and MECS provide data during the eclipse that are of sufficient quality to allow a meaningful analysis.
Formally. the spectrunn is consistent iu shape with that outside the eclipse (A2=1.18 for 28 dof).
Formally, the spectrum is consistent in shape with that outside the eclipse $\chi^2_{\rm r}=1.18$ for 28 dof).
Nevertheless.the spectiiun has the appearance of bene somewhat softer.
Nevertheless,the spectrum has the appearance of being somewhat softer.
If. iu the Comptonized model. ki. is allowed to vary during the fit. it converges to a value of 1.54£0.9 keV (42=0.90 for 27 dof).
If, in the Comptonized model, $T_{\rm e}$ is allowed to vary during the fit, it converges to a value of $1.5\pm0.9$ keV $\chi^2_{\rm r}=0.90$ for 27 dof).
The F-est predicts a probability of less than 0.02 for a chance occurrence of the fuprovement in 47.
The F-test predicts a probability of less than 0.02 for a chance occurrence of the improvement in $\chi^2$.
The same kind of improvement can be obtained when leaving free the optical depth instead of the plasma temperature. so we conclude that the nature of the softening is unclear.
The same kind of improvement can be obtained when leaving free the optical depth instead of the plasma temperature, so we conclude that the nature of the softening is unclear.
The 2-10 keV. fiux i 2.34101 Cres ley 2 (2-10 keV) or of that outside the eclipse.
The 2-10 keV flux is $2.3\times10^{-11}$ erg $^{-1}$ $^{-2}$ (2-10 keV) or of that outside the eclipse.
Figs.
Figs.
Loo aud 5. show that there appear to be two different stages of ceress: the fast 35 s rie and a shoulder of a few luudred seconds.
\ref{fignfilc} and \ref{figlczoom} show that there appear to be two different stages of egress: the fast 35 s rise and a shoulder of a few hundred seconds.
We generated separate spectra for these two time intervals aud fitted them with the Comptouization model while keeping all parameters values fixed to those found for the out-of-eclipse spectra (Table 1)) except Ap.
We generated separate spectra for these two time intervals and fitted them with the Comptonization model while keeping all parameters values fixed to those found for the out-of-eclipse spectrum (Table \ref{tabnfifit}) ) except $N_{\rm H}$.
We used only LECS and MECS data because we are primarily interested i whether Ny changes during ceress.
We used only LECS and MECS data because we are primarily interested in whether $N_{\rm H}$ changes during egress.
Frou this we fiud that during the quick rie [Ng=(10+3)<1072 7 and during the shoulder (1.27x0.07)νE1077OD P 7.
From this we find that during the quick rise $N_{\rm H}=(10\pm3)\times10^{22}$ $^{-2}$ and during the shoulder $(1.27\pm0.07)\times10^{22}$ $^{-2}$ .
However. iftq we leave free in addition the nornalizations of the Comptonized spectrum. the scusitivity to measuring Vy
However, if we leave free in addition the normalizations of the Comptonized spectrum, the sensitivity to measuring $N_{\rm H}$
Estimations show that VLP is possibly resulted from the standing slow sausage modes coupling and resonating with the underlying photospheric 5-min p-mode oscillations.
Estimations show that VLP is possibly resulted from the standing slow sausage modes coupling and resonating with the underlying photospheric 5-min p-mode oscillations.
It may be associated with evolutive behaviors of the solar internal structures.
It may be associated with evolutive behaviors of the solar internal structures.
As VLPs have the largest magnitude of emission fluxes, we suggest that the modulations are amplified and form the main framework of the whole flare/CME eruptive processes.
As VLPs have the largest magnitude of emission fluxes, we suggest that the modulations are amplified and form the main framework of the whole flare/CME eruptive processes.
From Equ.(2) and (3), we can estimate the radius of the emission source region: Here, the period (P), radio central frequency (/) and radio emission intensity (F) can be obtained from radio observations, the loop length (L) can be estimated from the optical or other imaging observations approximately.
From Equ.(2) and (3), we can estimate the radius of the emission source region: Here, the period $P$ ), radio central frequency $f$ ) and radio emission intensity $F$ ) can be obtained from radio observations, the loop length $L$ ) can be estimated from the optical or other imaging observations approximately.
Then we may obtain the radius of the radio emission source region (r) by adopting Equ.(7) even if we have no radio imaging observations.
Then we may obtain the radius of the radio emission source region $r$ ) by adopting Equ.(7) even if we have no radio imaging observations.
By substituting the parameters obtained in the above sections, we may get the radius of the emission source region from VLP paragraph A, B, C to D is from 2.04x10° km, 2.06x10° km, 1.62x10° km to 0.92x10? km, ie., the source region is undergoing an evolutive process of expanding at first, and then shrinking.
By substituting the parameters obtained in the above sections, we may get the radius of the emission source region from VLP paragraph A, B, C to D is from $2.04\times10^{5}$ km, $2.06\times10^{5}$ km, $1.62\times10^{5}$ km to $0.92\times10^{5}$ km, i.e., the source region is undergoing an evolutive process of expanding at first, and then shrinking.
However, as we have no radio imaging observations at the corresponding frequencies, we do not know the exact cites of the emission source region, the only thing we can do is to adopt the averaged loop length in our above estimations, which may have much uncertainties.
However, as we have no radio imaging observations at the corresponding frequencies, we do not know the exact cites of the emission source region, the only thing we can do is to adopt the averaged loop length in our above estimations, which may have much uncertainties.
Similar to VLP, LPP and SPP (may include part of slow-VSP) are also caused by MHD oscillations.
Similar to VLP, LPP and SPP (may include part of slow-VSP) are also caused by MHD oscillations.
However, their MHD modes may have some differences.
However, their MHD modes may have some differences.
They may be related with the standing fast sausage or kink modes.
They may be related with the standing fast sausage or kink modes.
The propagating MHD modes and the LRC-circuit resonation of current-carrying plasma loops are also the possible candidates of the generating mechanism.
The propagating MHD modes and the LRC-circuit resonation of current-carrying plasma loops are also the possible candidates of the generating mechanism.
Fast-VSP and most part of slow-VSP are generated by a completely different mechanism: the modulation of the resistive tearing-mode oscillations in electric current-carrying flare loops.
Fast-VSP and most part of slow-VSP are generated by a completely different mechanism: the modulation of the resistive tearing-mode oscillations in electric current-carrying flare loops.
In this mechanism, both, the period and duration of QPP are coupled with the magnetic field, plasma density, electric current, and the loop parameters.
In this mechanism, both, the period and duration of QPP are coupled with the magnetic field, plasma density, electric current, and the loop parameters.
By using their relation, we may deduce the physical conditions of the emission source region.
By using their relation, we may deduce the physical conditions of the emission source region.
The timescale of periods of QPPs implies a limit on the pulsating emission source size.
The timescale of periods of QPPs implies a limit on the pulsating emission source size.
Regardless of generating mechanism the pulsating source must be smaller than that given by the product of speed of light and period (P).
Regardless of generating mechanism the pulsating source must be smaller than that given by the product of speed of light and period $P$ ).
If not, the pulsating structure would be smeared out (Elgardy, 1986).
If not, the pulsating structure would be smeared out $\phi$ y, 1986).
So, it is reasonable to suppose that the short periodic QPP may come from a smaller source region.
So, it is reasonable to suppose that the short periodic QPP may come from a smaller source region.
The broad hierarchy of timescales of QPPs occurred in a flare event may imply that there is a multi-scale hierarchy of sizes of the magnetic configurations in the flaring region, and timescales of the dynamic processes.
The broad hierarchy of timescales of QPPs occurred in a flare event may imply that there is a multi-scale hierarchy of sizes of the magnetic configurations in the flaring region, and timescales of the dynamic processes.
The frequency drift rate and the bandwidth of the pulsating emission are dominated by the emission mechanism which is always related to the magnetic field strength, plasma density, and possibly to the plasma temperature.
The frequency drift rate and the bandwidth of the pulsating emission are dominated by the emission mechanism which is always related to the magnetic field strength, plasma density, and possibly to the plasma temperature.
It may be reasonable to suppose that the frequency drift features of QPPs implies the motion of the pulsating source regions, and the bandwidth of the pulsating emission are related to the dimensional size of the pulsating source regions.
It may be reasonable to suppose that the frequency drift features of QPPs implies the motion of the pulsating source regions, and the bandwidth of the pulsating emission are related to the dimensional size of the pulsating source regions.