{
    "count": 15,
    "next": null,
    "previous": null,
    "results": [
        {
            "id": 3829,
            "url": "https://svs.gsfc.nasa.gov/3829/",
            "result_type": "Visualization",
            "release_date": "2011-05-10T00:00:00-04:00",
            "title": "Aquarius studies Ocean and Wind Flows",
            "description": "Aquarius is a focused satellite mission to measure global Sea Surface Salinity. During its nominal three-year mission, Aquarius will map the salinity at the ocean surface to improve our understanding of Earth's water cycle and ocean circulation. Aquarius will help scientists see how freshwater moves between the ocean and the atmosphere. It will monitor changes in the water cycle due to rainfall, evaporation, ice melting, and river runoff. Aquarius will also demonstrate a measurement capability that can be applied to future operational missions. Ocean circulation is driven in large part by changes in water density, which is determined by temperature and salinity. Cold, high-salinity water masses sink and trigger the ocean's \"themalhaline circulation\" - the surface and deep currents that distribute solar energy to regulate Earth's climate. By measuring salinity, Aquarius will provide new insight into this global process. Aquarius' measurements of ocean salinity will provide a new perspective on the ocean and its links to climate, greatly expanding upon limited past measurements. Aquarius salinity data - combined with data from other sensors that measure sea level, ocean color, temperature, winds and rainfall will give us a much clearer picture of how the ocean works, how it is linked to climate, and how it may respond to climate change.Aquarius will provide information that will help improve predictions of future climate trends and short-term climate events such as El Niño and La Niña. Precise salinity measurements from Aquarius will reveal changes in patterns of global precipitation and evaporation and show how these changes may affect ocean circulation. || ",
            "hits": 93
        },
        {
            "id": 10740,
            "url": "https://svs.gsfc.nasa.gov/10740/",
            "result_type": "Produced Video",
            "release_date": "2011-04-07T09:00:00-04:00",
            "title": "When Neutron Stars Collide",
            "description": "Armed with state-of-the-art supercomputer models, scientists have shown that colliding neutron stars can produce the energetic jet required for a gamma-ray burst. Earlier simulations demonstrated that mergers could make black holes. Others had shown that the high-speed particle jets needed to make a gamma-ray burst would continue if placed in the swirling wreckage of a recent merger. Now, the simulations reveal the middle step of the process—how the merging stars' magnetic field organizes itself into outwardly directed components capable of forming a jet. The Damiana supercomputer at Germany's Max Planck Institute for Gravitational Physics needed six weeks to reveal the details of a process that unfolds in just 35 thousandths of a second—less than the blink of an eye.For the researchers' website, with more video and stills of their simulations, go here. || ",
            "hits": 472
        },
        {
            "id": 3816,
            "url": "https://svs.gsfc.nasa.gov/3816/",
            "result_type": "Visualization",
            "release_date": "2011-01-21T00:00:00-05:00",
            "title": "The Thermohaline Circulation - The Great Ocean Conveyor Belt - Stereoscopic Version",
            "description": "The oceans are mostly composed of warm salty water near the surface over cold, less salty water in the ocean depths. These two regions don't mix except in certain special areas. The ocean currents, the movement of the ocean in the surface layer, are driven primarily by the wind. In certain areas near the polar oceans, the colder surface water also gets saltier due to evaporation or sea ice formation. In these regions, the surface water becomes dense enough to sink to the ocean depths. This pumping of surface water into the deep ocean forces the deep water to move horizontally until it can find an area on the world where it can rise back to the surface and close the current loop. This usually occurs in the equatorial ocean, mostly in the Pacific and Indian Oceans. This very large, slow current is called the thermohaline circulation because it is caused by temperature and salinity (haline) variations.This animation shows one of the major regions where this pumping occurs, the North Atlantic Ocean around Greenland, Iceland, and the North Sea. The surface ocean current brings new water to this region from the South Atlantic via the Gulf Stream and the water returns to the South Atlantic via the North Atlantic Deep Water current. The continual influx of warm water into the North Atlantic polar ocean keeps the regions around Iceland and southern Greenland generally free of sea ice year round.The animation also shows another feature of the global ocean circulation: the Antarctic Circumpolar Current. The region around latitude 60 south is the the only part of the Earth where the ocean can flow all the way around the world with no obstruction by land. As a result, both the surface and deep waters flow from west to east around Antarctica. This circumpolar motion links the world's oceans and allows the deep water circulation from the Atlantic to rise in the Indian and Pacific Oceans, thereby closing the surface circulation with the northward flow in the Atlantic.The color on the world's ocean's at the beginning of this animation represents surface water density, with dark regions being most dense and light regions being least dense (see the animation Sea Surface Temperature, Salinity and Density). The depths of the oceans are highly exaggerated to better illustrate the differences between the surface flows and deep water flows. The actual flows in this model are based on current theories of the thermohaline circulation rather than actual data. The thermohaline circulation is a very slow moving current that can be difficult to distinguish from general ocean circulation. Therefore, it is difficult to measure or simulate.This is a stereoscopic version of the original visualziation. || ",
            "hits": 293
        },
        {
            "id": 3652,
            "url": "https://svs.gsfc.nasa.gov/3652/",
            "result_type": "Visualization",
            "release_date": "2009-10-09T13:24:00-04:00",
            "title": "Sea Surface Temperature, Salinity and Density",
            "description": "Sea Surface TemperatureThe oceans of the world are heated at the surface by the sun, and this heating is uneven for many reasons. The Earth's axial rotation, revolution about the sun, and tilt all play a role, as do the wind-driven ocean surface currents. The first animation in this group shows the long-term average sea surface temperature, with red and yellow depicting warmer waters and blue depicting colder waters. The most obvious feature of this temperature map is the variation of the temperature by latitude, from the warm region along the equator to the cold regions near the poles. Another visible feature is the cooler regions just off the western coasts of North America, South America, and Africa. On these coasts, winds blow from land to ocean and push the warm water away from the coast, allowing cooler water to rise up from deeper in the ocean. || ",
            "hits": 640
        },
        {
            "id": 3658,
            "url": "https://svs.gsfc.nasa.gov/3658/",
            "result_type": "Visualization",
            "release_date": "2009-10-08T00:00:00-04:00",
            "title": "The Thermohaline Circulation - The Great Ocean Conveyor Belt",
            "description": "The oceans are mostly composed of warm salty water near the surface over cold, less salty water in the ocean depths. These two regions don't mix except in certain special areas. The ocean currents, the movement of the ocean in the surface layer, are driven mostly by the wind. In certain areas near the polar oceans, the colder surface water also gets saltier due to evaporation or sea ice formation. In these regions, the surface water becomes dense enough to sink to the ocean depths. This pumping of surface water into the deep ocean forces the deep water to move horizontally until it can find an area on the world where it can rise back to the surface and close the current loop. This usually occurs in the equatorial ocean, mostly in the Pacific and Indian Oceans. This very large, slow current is called the thermohaline circulation because it is caused by temperature and salinity (haline) variations.This animation shows one of the major regions where this pumping occurs, the North Atlantic Ocean around Greenland, Iceland, and the North Sea. The surface ocean current brings new water to this region from the South Atlantic via the Gulf Stream and the water returns to the South Atlantic via the North Atlantic Deep Water current. The continual influx of warm water into the North Atlantic polar ocean keeps the regions around Iceland and southern Greenland mostly free of sea ice year round.The animation also shows another feature of the global ocean circulation: the Antarctic Circumpolar Current. The region around latitude 60 south is the the only part of the Earth where the ocean can flow all the way around the world with no land in the way. As a result, both the surface and deep waters flow from west to east around Antarctica. This circumpolar motion links the world's oceans and allows the deep water circulation from the Atlantic to rise in the Indian and Pacific Oceans and the surface circulation to close with the northward flow in the Atlantic.The color on the world's ocean's at the beginning of this animation represents surface water density, with dark regions being most dense and light regions being least dense (see the animation Sea Surface Temperature, Salinity and Density). The depths of the oceans are highly exaggerated to better illustrate the differences between the surface flows and deep water flows. The actual flows in this model are based on current theories of the thermohaline circulation rather than actual data. The thermohaline circulation is a very slow moving current that can be difficult to distinguish from general ocean circulation. Therefore, it is difficult to measure or simulate. || ",
            "hits": 198
        },
        {
            "id": 39,
            "url": "https://svs.gsfc.nasa.gov/39/",
            "result_type": "Visualization",
            "release_date": "1994-02-12T12:00:00-05:00",
            "title": "Rayleigh-Taylor Instabilities in Supernovae Explosions: Density",
            "description": "The following calculation shows the development and evolution of Rayleigh-Taylor instabilities which develop behind the supernova blast wave on a time scale of a few hours. The initial model was chosen to provide a good representation for the progenitor star for Supernova 1987A. The calculation was performed using the Piecewise-Parabolic Method for hydrodynamics on a two-dimensional spherical grid with rotational symmetry about the vertical axis and equatorial symmetry about the horizontal axis.The grid contained 800 zones in the radial direction and 400 zones in the angular diraction and was allowed to expand homologously with the explosion to maintain as high a resolution as possible in the unstable layer during the evolution. The following sequences show the evolution of the density distribution as well as the distribution of hydrogen, helium, and oxygen within the ejecta to illustrate the amount of mixing caused by the instability. Each sequence shows the evolution in two reference frames.In the first frame, the size of the plot expands with time as the grid expands. For the second reference frame, the size of the plot is kept fixed with the time so that more detail can be seen in the unstable layer. || ",
            "hits": 75
        },
        {
            "id": 41,
            "url": "https://svs.gsfc.nasa.gov/41/",
            "result_type": "Visualization",
            "release_date": "1994-02-12T12:00:00-05:00",
            "title": "Rayleigh-Taylor Instabilities in Supernovae Explosions: Partial Density of Hydrogen",
            "description": "The following calculation shows the development and evolution of Rayleigh-Taylor instabilities which develop behind the supernova blast wave on a time scale of a few hours. The initial model was chosen to provide a good representation for the progenitor star for Supernova 1987A. The calculation was performed using the Piecewise-Parabolic Method for hydrodynamics on a two-dimensional spherical grid with rotational symmetry about the vertical axis and equatorial symmetry about the horizontal axis.The grid contained 800 zones in the radial direction and 400 zones in the angular diraction and was allowed to expand homologously with the explosion to maintain as high a resolution as possible in the unstable layer during the evolution. The following sequences show the evolution of the density distribution as well as the distribution of hydrogen, helium, and oxygen within the ejecta to illustrate the amount of mixing caused by the instability. Each sequence shows the evolution in two reference frames.In the first frame, the size of the plot expands with time as the grid expands. For the second reference frame, the size of the plot is kept fixed with the time so that more detail can be seen in the unstable layer. || ",
            "hits": 38
        },
        {
            "id": 43,
            "url": "https://svs.gsfc.nasa.gov/43/",
            "result_type": "Visualization",
            "release_date": "1994-02-12T12:00:00-05:00",
            "title": "Rayleigh-Taylor Instabilities in Supernovae Explosions: Partial Density of Helium",
            "description": "The following calculation shows the development and evolution of Rayleigh-Taylor instabilities which develop behind the supernova blast wave on a time scale of a few hours. The initial model was chosen to provide a good representation for the progenitor star for Supernova 1987A. The calculation was performed using the Piecewise-Parabolic Method for hydrodynamics on a two-dimensional spherical grid with rotational symmetry about the vertical axis and equatorial symmetry about the horizontal axis.The grid contained 800 zones in the radial direction and 400 zones in the angular diraction and was allowed to expand homologously with the explosionto maintain as high a resolution as possible in the unstable layer during the evolution. The following sequences show the evolution of the density distribution as well as the distribution of hydrogen, helium, and oxygen within the ejecta to illustrate the amount of mixing caused by the instability. Each sequence shows the evolution in two reference frames.In the first frame, the size of the plot expands with time as the grid expands. For the second reference frame, the size of the plot is kept fixed with the time so that more detail can be seen in the unstable layer. || ",
            "hits": 36
        },
        {
            "id": 45,
            "url": "https://svs.gsfc.nasa.gov/45/",
            "result_type": "Visualization",
            "release_date": "1994-02-12T12:00:00-05:00",
            "title": "Rayleigh-Taylor Instabilities in Supernovae Explosions: Partial Density of Oxygen",
            "description": "The following calculation shows the development and evolution of Rayleigh-Taylor instabilities which develop behind the supernova blast wave on a time scale of a few hours. The initial model was chosen to provide a good representation for the progenitor star for Supernova 1987A. The calculation was performed using the Piecewise-Parabolic Method for hydrodynamics on a two-dimensional spherical grid with rotational symmetry about the vertical axis and equatorial symmetry about the horizontal axis.The grid contained 800 zones in the radial direction and 400 zones in the angular diraction and was allowed to expand homologously with the explosion to maintain as high a resolution as possible in the unstable layer during the evolution. The following sequences show the evolution of the density distribution as well as the distribution of hydrogen, helium, and oxygen within the ejecta to illustrate the amount of mixing caused by the instability. Each sequence shows the evolution in two reference frames.In the first frame, the size of the plot expands with time as the grid expands. For the second reference frame, the size of the plot is kept fixed with the time so that more detail can be seen in the unstable layer. || ",
            "hits": 35
        },
        {
            "id": 46,
            "url": "https://svs.gsfc.nasa.gov/46/",
            "result_type": "Visualization",
            "release_date": "1994-02-12T12:00:00-05:00",
            "title": "Instabilities in Very Young Neutron Stars: Density",
            "description": "This simulation shows the first 20 milliseconds in the life of a neutron star which is formed in a Type II supernova. After an initial collapse phase, the neutron star becomes unstable to convection. The resulting convective motions destroy the spherical symmetry of the star and rapidly mix the inner regions. In addition, the neutrino flux from the neutron star will be non-spherical and will be significantly enhanced by the convective motions. This may have major implications for the Type II supernova mechanism. The calculation was performed using the Piecewise-Parabolic Method for hydrodynamics. The computational grid contained 300 zones in radius and 200 zones in angle. The inner 200 zones in radius were uniformly spaced, ranging from the inner boundary at 25 km to 175 km. The outer 100 zones were non-uniformly spaced and stretched to 2000 km. Only the inner 200 zones are plotted. The inner boundary was treated as a hard sphere. At the outer boundary, zero gradients for all the variables were assumed. Periodic boundary conditions were used along the sides of the grid. The following sequence shows the density evolution for 20 milliseconds after the shock stalls. The density is plotted on a log scale. Values range from 10^9 gm/cm^3 at the outer boundary to 1.4 x 10^12 gm/cm^3 at the inner boundary. || ",
            "hits": 54
        },
        {
            "id": 47,
            "url": "https://svs.gsfc.nasa.gov/47/",
            "result_type": "Visualization",
            "release_date": "1994-02-12T12:00:00-05:00",
            "title": "Instabilities in Very Young Neutron Stars: Temperature",
            "description": "This simulation shows the first 20 milliseconds in the life of a neutron star which is formed in a Type II supernova. After an initial collapse phase, the neutron star becomes unstable to convection. The resulting convective motions destroy the spherical symmetry of the star and rapidly mix the inner regions. In addition, the neutrino flux from the neutron star will be non-spherical and will be significantly enhanced by the convective motions. This may have major implications for the Type II supernova mechanism. The calculation was performed using the Piecewise-Parabolic Method for hydrodynamics. The computational grid contained 300 zones in radius and 200 zones in angle. The inner 200 zones in radius were uniformly spaced, ranging from the inner boundary at 25 km to 175 km. The outer 100 zones were non-uniformly spaced and stretched to 2000 km. Only the inner 200 zones are plotted. The inner boundary was treated as a hard sphere. At the outer boundary, zero gradients for all the variables were assumed. Periodic boundary conditions were used along the sides of the grid. The following sequence shows the temperature structure for 20 milliseconds after the shock stalls. The minimum temperature is approximately 1.35 MeV. The maximum temperature varies from 6 MeV at the beginning of the calculation to 10 MeV at the later times. || ",
            "hits": 50
        },
        {
            "id": 48,
            "url": "https://svs.gsfc.nasa.gov/48/",
            "result_type": "Visualization",
            "release_date": "1994-02-12T12:00:00-05:00",
            "title": "Instabilities in Very Young Neutron Stars: Electron Fraction",
            "description": "This simulation shows the first 20 milliseconds in the life of a neutron star which is formed in a Type II supernova. After an initial collapse phase, the neutron star becomes unstable to convection. The resulting convective motions destroy the spherical symmetry of the star and rapidly mix the inner regions. In addition, the neutrino flux from the neutron star will be non-spherical and will be significantly enhanced by the convective motions. This may have major implications for the Type II supernova mechanism. The calculation was performed using the Piecewise-Parabolic Method for hydrodynamics. The computational grid contained 300 zones in radius and 200 zones in angle. The inner 200 zones in radius were uniformly spaced, ranging from the inner boundary at 25 km to 175 km. The outer 100 zones were non-uniformly spaced and stretched to 2000 km. Only the inner 200 zones are plotted. The inner boundary was treated as a hard sphere. At the outer boundary, zero gradients for all the variables were assumed. Periodic boundary conditions were used along the sides of the grid. The following sequence shows the mixing of composition which results from the convective motions. The variable plotted is the electron fraction Ye, which ranges from 0.2 to 0.5. || ",
            "hits": 55
        },
        {
            "id": 1381,
            "url": "https://svs.gsfc.nasa.gov/1381/",
            "result_type": "Visualization",
            "release_date": "1994-02-12T12:00:00-05:00",
            "title": "Instabilities in Very Young Neutron Stars: Density",
            "description": "This simulation shows the first 20 milliseconds in the life of a neutron star which is formed in a Type II supernova. After an initial collapse phase, the neutron star becomes unstable to convection. The resulting convective motions destroy the spherical symmetry of the star and rapidly mix the inner regions. In addition, the neutrino flux from the neutron star will be non-spherical and will be significantly enhanced by the convective motions. This may have major implications for the Type II supernova mechanism. The calculation was performed using the Piecewise-Parabolic Method for hydrodynamics. The computational grid contained 300 zones in radius and 200 zones in angle. The inner 200 zones in radius were uniformly spaced, ranging from the inner boundary at 25 km to 175 km. The outer 100 zones were non-uniformly spaced and stretched to 2000 km. Only the inner 200 zones are plotted. The inner boundary was treated as a hard sphere. At the outer boundary, zero gradients for all the variables were assumed. Periodic boundary conditions were used along the sides of the grid. The following sequence shows the density evolution for 20 milliseconds after the shock stalls. The density is plotted on a log scale. Values range from 10^9 gm/cm^3 at the outer boundary to 1.4 x 10^12 gm/cm^3 at the inner boundary. || ",
            "hits": 33
        },
        {
            "id": 1382,
            "url": "https://svs.gsfc.nasa.gov/1382/",
            "result_type": "Visualization",
            "release_date": "1994-02-12T12:00:00-05:00",
            "title": "Instabilities in Very Young Neutron Stars: Temperature",
            "description": "This simulation shows the first 20 milliseconds in the life of a neutron star which is formed in a Type II supernova. After an initial collapse phase, the neutron star becomes unstable to convection. The resulting convective motions destroy the spherical symmetry of the star and rapidly mix the inner regions. In addition, the neutrino flux from the neutron star will be non-spherical and will be significantly enhanced by the convective motions. This may have major implications for the Type II supernova mechanism. The calculation was performed using the Piecewise-Parabolic Method for hydrodynamics. The computational grid contained 300 zones in radius and 200 zones in angle. The inner 200 zones in radius were uniformly spaced, ranging from the inner boundary at 25 km to 175 km. The outer 100 zones were non-uniformly spaced and stretched to 2000 km. Only the inner 200 zones are plotted. The inner boundary was treated as a hard sphere. At the outer boundary, zero gradients for all the variables were assumed. Periodic boundary conditions were used along the sides of the grid. The following sequence shows the temperature structure for 20 milliseconds after the shock stalls. The minimum temperature is approximately 1.35 MeV. The maximum temperature varies from 6 MeV at the beginning of the calculation to 10 MeV at the later times. || ",
            "hits": 30
        },
        {
            "id": 1383,
            "url": "https://svs.gsfc.nasa.gov/1383/",
            "result_type": "Visualization",
            "release_date": "1994-02-12T12:00:00-05:00",
            "title": "Instabilities in Very Young Neutron Stars: Electron Fraction",
            "description": "This simulation shows the first 20 milliseconds in the life of a neutron star which is formed in a Type II supernova. After an initial collapse phase, the neutron star becomes unstable to convection. The resulting convective motions destroy the spherical symmetry of the star and rapidly mix the inner regions. In addition, the neutrino flux from the neutron star will be non-spherical and will be significantly enhanced by the convective motions. This may have major implications for the Type II supernova mechanism. The calculation was performed using the Piecewise-Parabolic Method for hydrodynamics. The computational grid contained 300 zones in radius and 200 zones in angle. The inner 200 zones in radius were uniformly spaced, ranging from the inner boundary at 25 km to 175 km. The outer 100 zones were non-uniformly spaced and stretched to 2000 km. Only the inner 200 zones are plotted. The inner boundary was treated as a hard sphere. At the outer boundary, zero gradients for all the variables were assumed. Periodic boundary conditions were used along the sides of the grid. The following sequence shows the mixing of composition which results from the convective motions. The variable plotted is the electron fraction Ye, which ranges from 0.2 to 0.5. || ",
            "hits": 24
        }
    ]
}